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9804 | astro-ph9804298_arXiv.txt | We have produced radio maps, using the ATCA, of the central regions of six southern Seyfert 2 galaxies (NGC 1365, 4945, 6221, 6810, 7582, and Circinus) with circumnuclear star formation, to estimate the relative contribution of star formation activity compared to activity from the active galactic nucleus (AGN). The radio morphologies range from extended diffuse structures to compact nuclear emission, with no evidence, even in the relatively compact sources, for synchrotron self--absorption. In each case the radio to far--infrared (FIR) ratio has a value consistent with star formation, and in all but one case the radio to [FeII] ratio is also consistent with star formation. We derive supernova rates and conclude that, despite the presence of a Seyfert nucleus in these galaxies, the radio, FIR, and [FeII] line emission are dominated by processes associated with the circumnuclear star formation (i.e. supernova remnants and H~II regions) rather than with the AGN. | Circumnuclear star formation is common in Seyfert galaxies, but the relationship between the Seyfert nucleus and the surrounding star formation is not well understood (see Filippenko 1992), and both evolutionary and causal relationships have been suggested. For example, a nuclear starburst may develop into a massive star cluster or black hole (Norman \& Scoville 1988), or conversely the energy output from an active galactic nucleus (AGN) may trigger circumnuclear star formation (Sanders \& Bania 1976). More recently, studies have indicated that star formation is occurring in and around the torus (e.g. Cid Fernandes \& Terlevich 1992; Davies \etal 1997) that is thought to dictate the type of Seyfert nucleus we observe. Dust is a common feature of the circumnuclear environment of active galaxies. At optical wavelengths, it can obscure our view of the nucleus and hide any evidence of an AGN. However, radio observations are not only unaffected by dust but also have the advantage of high spatial resolution, and can be important in understanding the role of the various processes in active galaxies. For example, starburst galaxies generally have diffuse radio morphologies dominated by synchrotron emission from cosmic rays accelerated by supernovae, while Seyfert galaxies sometimes exhibit well--collimated radio jets and a compact nuclear radio source (see review by Condon 1992). High-resolution radio synthesis images, with sub-arcsec resolution, have found compact sources at the nucleus, and extended emission consisting of radio jets, ouflows, diffuse emission, and discrete sources such as HII regions and SNRs (e.g. Ulvestad \& Wilson 1984; Kronberg \etal 1985; Antonucci \& Ulvestad 1988; Carral, Turner \& Ho 1990; Condon \etal 1991; Forbes \etal 1994; Collison \etal 1994; Sandqvist \etal 1995). High spatial resolution, multi--frequency radio observations can also be used to test the idea of advection-dominated accretion flows around black holes. For such flows the radio emission depends strongly on the mass of the black hole and is characterised by an inverted radio spectrum (Yi \& Boughn 1997). Here we present 3 and 6cm radio continuum images from the Australia Telescope Compact Array (ATCA) of 6 such galaxies which show evidence for narrow high excitation optical forbidden lines, and are classified as Seyfert 2s. In the case of NGC 1365, broad lines have also been detected from the nuclear region. All six galaxies show evidence for circumnuclear star formation, and several are barred. We discuss the radio morphologies of these galaxies and possible emission mechanisms, and we compare the radio data with that from other wavelengths to assess the relative contributions to the radio flux from star formation and the AGN. | In Figures 1 to 12 we show the 3 and 6cm radio images for the six galaxies in our sample. These galaxies do not represent a complete sample in any sense, but rather were chosen as relatively well--known southern Seyfert 2 galaxies that generally lacked high resolution radio maps. Our radio data for the Circinus galaxy have been presented elsewhere along with near--infrared line imaging (Davies \etal 1997), although we include it here for comparison purposes with the other galaxies. The images show a variety of radio morphologies which range from Circinus, with a strong, compact nucleus, to NGC 1365, with a extended region of discrete sources or hot-spots. The beam shape is shown in the lower left of each image. Care is needed when interpreting these images, as (a) our observations are optimised for studying the nuclear region, and so our images will not correctly represent the extended diffuse emission from the disk of the galaxy, and (b) most of the galaxies have high inclinations so that confusion effects may play some role in the observed radio morphology. Flux measurements in 2 arcsec and 6 arcsec diameter apertures for each galaxy are given in Table 2. We also list the 6cm to 3cm spectral index after smoothing the 3cm image to match the 6cm resolution. The spectral indices (F $\propto \nu ^{\alpha}$) range from flat ($\alpha$ $\sim$ 0) to steep ($\alpha$ $\sim$ --1). The total 6cm flux in Table 2 is given both for our images and for the single--dish observations by Wright \etal (1994, 1996). The difference between these indicates the amount of diffuse emission which is missing from our images. We also calculated the maximum brightness temperatures for each image (both 3 and 6 cm) and the maximum value is given for each source. In Table 3 we list various derived quantities for the sample including Hubble type, distance, inclination, 6cm and [FeII] line luminosities and the SN rate. The [FeII] measurements are from Moorwood \& Oliva (1988) in a 6 arcsec diameter aperture, except for the Circinus galaxy in which we use a nuclear [FeII] flux from Davies \etal (1997). None of the [FeII] fluxes have been corrected for extinction. The SN rate is calculated separately from both the 6cm and [FeII] line flux in the 6 arcsec aperture (corresponding to $\sim$ 500 pc at a typical distance of 20 Mpc). The 6cm SN rate is calculated assuming that all of the 6cm flux is non--thermal emission from cosmic rays accelerated by SNRs (e.g. Condon \& Yin 1990). This may give an overestimate of the rate because it ignores any nuclear flux (which may be significant) and the contribution from thermal emission (which is unlikely to be significant). For the [FeII] SN rate we simply assume an average luminosity of 2 $\times$ 10$^{36}$ erg s$^{-1}$ over an adiabatic lifetime of 2 $\times$ 10$^{4}$ yrs (e.g. Norris \& Forbes 1995). For most galaxies the two methods give rates within a factor of two, the notable exception being NGC 4945 (which is discussed further below). \subsection{Radio Spectral Indices} The radio spectral index of Seyfert and starburst galaxies is determined by four mechanisms. 1. Cosmic rays within the galaxy are generated and re--accelerated by supernovae and supernova remnants, and then interact with the interstellar magnetic field to emit synchrotron emission. This synchrotron emission, with a typical spectral index of $\alpha$ $\sim$ --0.7, is expected to dominate the radio power of starburst galaxies, and should appear as a diffuse component in radio images of these sources. 2. Relativistic particles ejected from the massive black hole at the nucleus of a galaxy may generate intense synchrotron emission, similar to that seen in radio--loud galaxies and quasars. The cores in these radio-loud objects typically have a flat--spectrum core, indicating synchrotron self--absorption, and steep--spectrum extended radio--lobes, suggesting cooling of high--energy electrons. However, synchrotron self--absorption is important only for brightness temperatures greater than 10$^{10}$ K (Condon 1992). Most Seyfert galaxies, on the other hand, are observed to have much lower brightness temperatures than this in the core, so that synchrotron self-absorption is not significant in these sources. This is confirmed by the observed core spectral indices, which are frequently in the region of --0.7. 3. When the radiative efficiency in the accretion disk is low, an advection-dominated accretion flow (ADAF) may operate. The radio emission in these sources is dominated by synchrotron emission from a hot plasma, and the emission from such flows is predicted to give rise to inverted spectra with typical indices of +0.4 (Yi \& Boughn 1997). The ADAF radio emission mechanism has only recently been proposed and should be regarded as untested at this stage. The data here are unlikely to provide a definitive test because of insufficient resolution, and all radio spectral indices measured here are negative. We will therefore not consider this mechanism further, except to note that this mechanism would be indicated by inverted-spectrum emission from a low-brightness-temperature core. No source here falls into this category. 4. H II regions in our galaxy generate free--free emission from hot electrons. Most are optically thin, giving a flat spectrum, although some compact H II regions become optically thick at centimetre wavelengths, giving a spectral index $\sim$+2. However, the integrated flux of such regions is generally insignificant compared to the synchrotron emission of the host galaxy. 5. The radio emission from ultra--luminous infrared galaxies is optically thick to free--free absorption, so that the typical synchrotron spectrum of these galaxies is flattened at low frequencies (Condon et al. 1991). The combined result of these effects in Seyfert and starburst galaxies is to produce a typical radio spectral index of --0.7 (from the extended synchrotron emission) with a flattening at low frequencies in some starburst sources because of free--free absorption. Table 2 shows that the nuclear spectral indices of three of the galaxies (NGC 1365, NGC 6221, NGC 7582) is --0.5 or steeper on both the 2 arcsec and the 6 arcsec scale, showing evidence for neither free--free nor synchrotron absorption. In the other three galaxies, the cores have flatter spectra, but the brightness temperatures ($\le$ 8300 K) are too low for synchrotron self--absorption, indicating that free--free absorption is responsible for the flattening. Of course, we cannot rule out the presence of a weak synchrotron self--absorbed core in the nucleus of any of these galaxies. However, comparison of the radio fluxes in a 2--arcsec aperture with the flux in a 6--arcsec aperture in Table 2 shows that the luminosity of any such core is small compared to the surrounding emission. Therefore any such core does not contribute significantly to the overall energy budget of the nuclear region of the galaxy, and is not responsible for the overall flat spectrum of the nuclear region.. This degree of free-free absorption flattening indicates either a high star formation rate (Condon \etal 1991) or that we are viewing the AGN through optically--thick narrow-line-region clouds (Roy \etal 1994). \subsection{The Radio - [FeII] Correlation} Forbes \& Ward (1993) discovered that the 6cm radio emission in the central regions of active galaxies is strongly correlated with the near--infrared [FeII] 1.64$\mu$m line emission. This relation exists over several orders of magnitude. With a larger sample, Simpson \etal (1996) were able to show that Seyfert and starburst galaxies follow slightly different radio--[FeII] relations. For starburst galaxies the relation, with slope $\sim$ 1, can be reasonably explained by SNRs which are responsible for both the non--thermal radio emission and the fast shocks that provide the [FeII] excitation. However, the situation for Seyfert galaxies (which reveal a correlation slope of $\sim$ 0.7) is less clear. Simpson \etal argued that photo-ionisation from the Seyfert nucleus can cause this relationship, with a contribution from radio--jet induced shocks in some cases. In Fig. 13 we show the [FeII]/6cm ratio for our sample galaxies, compared with the Seyfert and starburst relations of Simpson \etal (1996). The 1$\sigma$ dispersion of galaxies about the relations is $\sim$ $10^{0.5}$. For the Seyferts studied here, we find a large degree of star--formation activity compared to photo--ionisation from an AGN, and so we might expect them to lie closer to the starburst relation than the Seyfert one. This indeed appears to be the case for four galaxies, although one (NGC 7582) is closer to the Seyfert relation and NGC 4945 falls well away from either relation. The [FeII]/6cm ratio of NGC 4945 is about a factor of 100 lower than typical active galaxies, and we discuss this further in Section 3.4 below. We note however that given the dispersion in the relations, and the low luminosities ($\le$ 10$^{40}$ erg s$^{-1}$) of the galaxies studied here, this is not a sensitive test of the excitation mechanism. \subsection {The Radio -- FIR Correlation} Normal spiral and starburst galaxies show a tight correlation between their radio and FIR luminosity (e.g. Wunderlich et al. 1987). This correlation, which extends over five orders of magnitude, is true for both flux density and luminosity, and cannot be accounted for by selection effects, or by a simple ``richness effect''. While a detailed mechanism to explain this correlation has yet to be established, it is almost certainly the result of star formation, which generates both the synchrotron radio emission and the thermal FIR emission. This is supported by the fact that all objects that are dominated by star formation (HII galaxies, normal spirals, starburst galaxies) do follow the correlation. On the other hand, Sopp \& Alexander (1991) showed that radio--loud quasars and radio galaxies clearly do not follow the radio--FIR correlation. Thus whether or not a galaxy follows this correlation may be used as an indicator of the dominant radio luminosity source of the galaxy. Norris \etal (1988) and Roy \etal (1997) showed that Seyfert galaxies, unlike radio--loud quasars, do roughly follow the radio--FIR correlation, but with a looser fit than normal spirals and starbursts. This suggests that the bolometric luminosity of Seyfert galaxies may be dominated by star formation. This is supported by off--nuclear optical and infrared observations of Seyferts, which show the same line ratios and luminosities as starburst galaxies (Bransford et al. 1997). Thus, although the nuclear optical spectra of Seyfert galaxies are clearly dominated by an AGN, in many cases the integrated radio emission and the FIR emission are dominated not by the AGN but by circumnuclear star formation. The degree to which an individual galaxy follows this correlation is most conveniently expressed by the parameter q -- the logarithm of the FIR to radio ratio. The conventional definition of q follows that of Helou \etal (1985), who define q in terms of the 1.49 GHz radio flux. For our purposes, we adapt Helou's definition to our observing frequency of 4.8 GHz by assuming a spectral index of --0.7, and therefore define it as \hspace{.5in}q $\equiv$ log[(FIR/3.75 x 10$^{12}$ Hz)/(2.26 x S$_{\rm 4.8 GHz}$]\hspace{1.5in}(1) \hspace{.5in}where FIR $\equiv$ 1.26 x 10$^{-14}$(2.58S$_{\rm 60\mu}$+S$_{\rm 100\mu}$)\hspace{2in}(2) Typical values of q from the IRAS Bright Galaxy Sample are 2.34 for normal spirals, 2.21 for starburst galaxies, and less than 2 for radio--loud AGNs (Condon \etal 1991). All the galaxies studied here except NGC 4945 have q in the range 2.2 to 2.3, which places them firmly in the middle of the radio--FIR correlation, and suggests that most of their radio and FIR luminosity is produced by star formation. We discuss the case of NGC 4945 (q = 1.88) below. \subsection{Individual Galaxies} Here we discuss each galaxy in turn, starting with an extended discussion of NGC 4945. To avoid repetition, we note that in every case other than NGC 4945, the spectral index, radio--FIR ratio, and [FeII]--radio ratio are all consistent with star formation, rather than an AGN, being the dominant source of radio emission. \noindent {\bf NGC 4945} This infrared luminous galaxy is nearly edge--on and is located in a nearby loose group. Although we list it as a barred galaxy in Table 3, there is a continuing debate about the reality of the bar (e.g. Harnett \etal 1989). Koornneef (1993) described NGC 4945 as a post--starburst galaxy with no evidence for an AGN. However Moorwood \& Oliva (1994) have argued that the central regions do show signs of ongoing young star formation. Evidence for a heavily obscured AGN now come from the variable hard X--rays (Iwasawa \etal 1993), and the presence of a compact radio core in VLBI observations (Sadler \etal 1995). The galaxy contains a thick torus or ring with a radius of $\sim$ 150 pc (Koornneef 1993; Moorwood \etal 1996). Harnett \etal (1989) found that the radio emission has a strong central contribution with emission extended 10 arcmin perpendicular to the major axis. Multi--frequency observations have been carried out by Elmouttie \etal (1997). They focused on the large scale structure using a beam size of $\sim$ 25$^{''}$, and found that the spectral index steepens from the central region to the main disk of the galaxy. Furthermore, NGC 4945 is notable for the fact that it is one of the few galaxies (along with Circinus) known to contain water megamasers. Such megamasers have been cited in NGC 4258 (Miyoshi \etal 1995) as the strongest evidence known for a massive black hole in an AGN. Preliminary VLBI imaging (Greenhill \etal 1997) of the megamasers supports the model that they are in a Keplerian disk surrounding the black hole. Our radio image is dominated by strong nuclear emission, and emission extended along the disk of the galaxy. However there is also evidence of some filamentary structure perpendicular to the major axis. Such extended emission may be associated with the outflowing superwind in this galaxy (Nakai \etal 1989; Lipari, Tsvetanov \& Macchetto 1997). The extended emission has a steeper spectral index ($\alpha \sim -0.8$) than the radio nucleus, however we note that our data are less sensitive to extended emission (particularly at 3cm) which makes the spectral index somewhat less certain. The nucleus has a relatively flat spectral index of $\alpha$ = --0.3. The brightness temperature in the central 2 arcsec of this source (i.e. 7000 K) is still far too low for synchrotron self--absorption, indicating that the star formation activity is particularly intense, to provide the necessary free--free absorption. The obscuration inferred from the X--ray data suggest that the extinction towards the AGN could be as high as A$_V$ $\sim$ 2500. We noted above (in Section 3.3) that the radio--FIR ratio for this galaxy is unusually low (i.e. q = 1.88), which at first sight appears to suggest that an AGN is responsible for much of the radio emission. However, this galaxy is so near that not all the FIR flux was in the single IRAS aperture, and so the IRAS flux listed in the IRAS Point Source Catalog may be an underestimate. Rice \etal (1988) have estimated the total FIR flux by co--adding IRAS images and obtain a higher value for the FIR fluxes, which raises the value of q to 2.1, suggesting that the radio emission in this galaxy is again dominated by star formation rather than by an AGN. An interesting property of NGC 4945 is that the ratio of the [FeII] line luminosity to 6cm radio emission is only 0.63, which is almost a factor of 100 less than is typical for active galaxies (see Fig. 13). We now consider a number of possible reasons for this. \begin{itemize} \item The reduced ratio could be due to extinction (by dust) of the [FeII]. However, Moorwood \& Oliva (1994) estimate that the extinction in the [FeII] line emitting zone is 1.8 mag or a factor of five, which is insufficient to produce the observed effect. \item It could be because of nuclear radio emission from an AGN which is not accompanied by [FeII] line emission. We have shown above that the radio/FIR ratio for the galaxy as a whole is consistent with star formation activity. However, the central 6 arcsec (over which we measure the [FeII]/6cm ratio) contributes only 9\% of the total radio flux, and we have no information on the radio/FIR ratio in the nucleus, so the radio flux from the AGN could be abnormally large. In this case, we would expect the [FeII]/6cm ratio to approach the usual value as we increase the area over which we integrate the flux (for both 6 cm and [FeII]). However, Moorwood \& Oliva (1994) quote a total [FeII] flux over an emitting region of 18 $\times$ 21 arcsec to be 12 $\times$ 10$^{-14}$ erg s$^{-1}$ cm$^{-2}$, or an observed log luminosity of 38.81 erg s$^{-1}$. The 6cm radio luminosity over a similar area is 40.21 erg s$^{-1}$, giving a [FeII]/6cm ratio of 0.03, which is even lower than the value in the nucleus, indicating that the [FeII]/6cm ratio falls off with distance from the galaxy centre, and that the low value is not a consequence of nuclear radio emission. \end {itemize} Thus the abnormally low [FeII]/6cm ratio in NGC 4945 of 0.63 is produced in the region surrounding the nucleus, where the radio emission (with a spectral index of $\sim$ --0.7) is due to SNRs in the galaxy disk, the outflowing starburst superwind discussed above, or perhaps a radio jet. Pure SNRs produce ratios of about 500 i.e. well in excess of typical galaxy values, so this would tend to give an enhanced ratio. In a $6^{''} \times 6^{''}$ aperture, the superwind galaxies M82 and NGC 253 have [FeII]/6cm ratios of 75 and 52 respectively, although the superwind itself in NGC 253 does not seem to produce significant [FeII] line emission (Forbes \etal 1993). Again such ratios are significantly higher than seen in NGC 4945. The data for Seyfert galaxies with clear radio jets are limited. For NGC 4151 and NGC 1068 the measured ratios are 28 and 9. This is closer to the NGC 4945 value but still a factor of at least 10 too high. We conclude that the abnormal [FeII]/6cm ratio in NGC 4945 is due to either (a) a starburst superwind, which produces substantial radio emission but very little [FeII] line flux (due perhaps to an unknown excitation effect or low density in the wind), or (b) a radio jet which dominates the extended radio emission on the few--arcsec scale but which does not produce significant [FeII] emission. \noindent {\bf NGC 1365} This is a well--studied barred galaxy. The central region reveals broad and narrow emission lines (Veron \etal 1980) surrounded by a circumnuclear ring of star formation (Edmunds \& Pagel 1982; Saikia \etal 1994). The star formation, combined with the obscuring effects of dust, give the appearance of a hot-spot nucleus (Sersic \& Pastoriza 1965). A high excitation outflow from the nucleus has been seen (e.g. Hjelm \& Lindblad 1996). High resolution radio continuum observations have been reported by several workers (e.g. Sandqvist, Jorsater \& Lindblad 1982, 1995). In particular, Sandqvist, Jorsater \& Lindblad (1995) observed it with the VLA at 20, 6, and 2 cm. Their radio images revealed a weak nucleus surrounded by a elongated $\sim$ 8 $\times$ 20 arcsec (a/b = 0.4) ring of hot-spots, or components. They labelled a number of components A--H, of which B is blended with A and C is blended with D at $\ge$ 1 arcsec resolutions. Our radio image, shown in Fig. 1, is consistent with theirs, except that we identify one additional component to the SW, which we call `J'. The hot-spots generally have steep spectra with 6cm luminosities of $\sim$ 10$^{36}$ erg s$^{-1}$ which suggests that the radio emission from each component is made up of several SNRs. A combined radio and X--ray study of the nucleus and surrounding regions has been carried out by Stevens, Forbes \& Norris (1998). The radio nucleus does not appear to have an X--ray counterpart. Furthermore the X--ray emission is consistent with star formation processes. Stevens \etal conclude that if NGC 1365 harbours a black hole it is largely inactive. \noindent {\bf NGC 6221} Located in a small group, NGC 6221, is a barred galaxy with a weak Seyfert nucleus. The galaxy may be interacting with NGC 6215 and has a nuclear bar (Koribalski 1996). The radio emission from NGC 6221 is extended in an symmetric bar--like structure. The spectral index of the nucleus and bar are non--thermal with $\alpha$ $\sim$ --0.6, indicative of SNRs. The radio morphology and other properties are all consistent with star formation being the dominant source of radio emission. \noindent {\bf NGC 6810} This early type spiral is the most distant in our sample and has not been well studied to date. It may contain a bar and ring structure (Buta 1995), and does not appear to have been imaged before at radio wavelengths. Our radio image reveals a dominant nucleus surrounded by diffuse extended radio emission. Both the nucleus and surrounding region have flat spectral indices, but the brightness temperature is too low for this to be attributable to synchrotron self--absorption, which suggests that the radio spectrum is flattened by free--free absorption from young star formation. The radio morphology and other radio properties are all consistent with star formation being the dominant source of radio emission. Interestingly, recent high resolution optical spectra do not confirm the status of NGC 6810 as a Seyfert galaxy (Heisler 1998), thus it appears to have been misclassified. \noindent {\bf NGC 7582} This narrow line X--ray galaxy is located in the Grus loose group along with NGC 7590 (a Seyfert 2; Ward \etal 1980), NGC 7552 (a starburst; Forbes \etal 1994) and NGC 7599. Several HI bridges connect group members (Koribalski 1996). Morris \etal (1985) provide evidence for both a rapidly--rotating $\sim$ 1 kpc ring of circumnuclear star formation and high excitation gas moving outwards from the nucleus. Ulvestad \& Wilson (1984) imaged the galaxy at 6 cm using the VLA with a beam size of $\sim$ 1.5 arcsec. They measured a total 6cm flux of 69 mJy. We find a linear, double--peaked morphology to the radio emission. The southern peak appears to lie at the centre of the outer radio isophotes and is presumably the true nucleus. The nucleus has a steep spectral index ($\alpha$ = --0.7) indicating non--thermal emission. To the NW by $\sim$ 3 arcsec lies a second peak, which could be a second nucleus. However, it lies roughly along the bar/major axis position angle (P.A. $\sim$ 150$^{\circ}$). It has a 2 arcsec diameter 6cm flux of 10 mJy and a spectral index of --0.9. This second peak could therefore be a radio jet or simply a discrete star formation region occurring along the galaxy bar. The inferred SN rate in the central 6 arcsec (870 pc) is the highest in our sample (except possibly for NGC 4945) at about 1 SN every 8 years. Although the radio morphology suggests a linear Seyfert jet, the spectral index, radio--FIR ratio, and [FeII]--radio ratio are all consistent with star formation being the dominant source of radio emission. \noindent {\bf Circinus} The Circinus galaxy is perhaps the closest Seyfert galaxy known but is difficult to observe due to its proximity to the Galactic plane and large internal obscuration. Confirmation of an AGN comes from the the presence of high excitation coronal lines (Oliva \etal 1994), X--ray emission (Matt \etal 1996), and a compact radio core (Heisler \etal 1998). Like NGC 4945, Circinus is one of the few water megamaser galaxies. The megamasers in Circinus are stronger but less extreme than those in NGC 4258, and have the curious property of fluctuating on a time scale of minutes (Greenhill \etal 1997), indicating a particularly compact source. Preliminary VLBI imaging of the megamasers (Ellingsen \etal 1998) indicates that the maser region is extended with a velocity gradient aligned with that of the parent galaxy, and perpendicular to the jet. We regard this as strong evidence for a massive black hole in this galaxy. Marconi \etal (1994) found both a circumnuclear starburst ring and an ionisation cone. The [OIII] ionisation cone is asymmetric extending only to the NW, with some high excitation lines also seen in the cone region. They estimated the extinction to the nucleus to be A$_V$ $\sim$ 20. The HI gas distribution shows spiral arms, a bar and a central `hole' (Koribalski 1996, Jones \etal 1998). High resolution observations of the central region indicate a rapidly--rotating gas ring with a diameter of $\sim$ 400 pc (Koribalski 1996). Observations with the ATCA have been carried out at 13 and 20 cm by Elmouttie \etal (1995). They found extended radio lobes perpendicular to the galaxy major axis (position angle = 30$^{\circ}$) with a spectral index of $\alpha \sim$ --0.7. Our 3 and 6cm radio images, observed as part of this project, have been published, along with near--infrared line images, by Davies \etal (1997). The radio data indicate that the nucleus is marginally resolved with a flat spectral index. The low brightness temperature indicates that this is due to free--free absorption rather than synchrotron self--absorption in a compact AGN source. There are also faint hints of extended emission which may be associated with outflowing material. Despite the clear indication of a compact AGN in the radio images, the other radio indicators are consistent with the more extended radio emission being dominated by star formation activity.\\ | 98 | 4 | astro-ph9804298_arXiv.txt |
9804 | astro-ph9804251_arXiv.txt | The innermost regions of quasars can be resolved by a gravitational-lens {\lq}telescope{\rq} on scales down to a few AU. For the purpose, X-ray observations are most preferable, because X-rays originating from the innermost regions, can be selectively amplified by microlensing due to the so-called `caustic crossing'. If detected, X-ray variations will constrain the size of the X-ray emitting region down to a few AU. The maximum attainable resolution depends mainly on the monitoring intervals of lens events, which should be much shorter than the crossing time. On the basis of this idea, we performe numerical simulations of microlensing of an optically-thick, standard-type disk as well as an optically-thin, advection-dominated accretion flow (ADAF). Calculated spectral variations and light curves show distinct behaviors, depending on the photon energy. X-ray radiation which is produced in optically thin region, exhibits intensity variation over a few tens of days. In contrast, optical-UV fluxes, which are likely to come from optically thick region, exhibit more gradual light changes, which is consistent with the microlensing events so far observed in Q2237+0305. Currently, Q2237+0305 is being monitored in the optical range at Apache Point Observatory. Simultaneous multi-wavelength observations by X-ray sattelites (e.g., ASCA, AXAF, XMM) as well as HST at the moment of a microlens event enable us to reveal an AU scale structure of the central accretion disk around black hole. | The high power output from quasars is usually attributed to the combination of a supermassive black hole with a surrounding accretion disk. This belief is supported by a number of observations that indicate the presence of a deep gravitational potential well or a hot gas disk at the center of quasars or other active galactic nuclei; e.g., measurements of stellar velocity dispersion clearly showed a peculiar increase toward the center (Young et al. 1978; Sargent et al. 1978; see also Ford et al. 1994; Harms et al. 1994). Malkan (1983) found that the optical to UV spectra are well fitted by the standard-type accretion disk model (Shakura \& Sunyaev 1973). Recently, by far the best evidence of a supermassive black hole has been found by radio observations of nuclear H$_2$O maser sources in NGC4258 (Miyoshi et al. 1995). Alternatively, we can infer the presence of a relativistic object from the asymmetric Fe line profile (Tanaka et al. 1995). These observational results are all attractive, but still the real vicinity of a putative black hole has not been resolved. Q2237+0305 (e.g., Huchra et al. 1985) is the first object, in which the quasar microlensing events were detected (Corrigan et al. 1987; Houde \& Racine 1994; see also Ostensen et al. 1996). These observations suggest that microlensing events take place roughly once per year. This rather high frequency is consistent with the microlens optical depth of $\tau \sim 0.8$ obtained by the realistic simulation of the lensing galaxy (i.e., Wambsganss \& Paczy\'nski 1994). We consider, here, specifically the microlensing events of this source caused by the so-called `caustic crossings' (see Yonehara et al. 1997 for single-lens calculations). Several authors have already analyzed this `caustic' case based on a simple model for quasar accretion disk (e.g., Wambsganss \& Paczy\'nski 1991; Jaroszy\'nski, Wambsganss \& Paczy\'nski 1992). So far, however, only the standard-type disk, which is too cool to emit X-rays, has been considered, and thus the property of an X-ray microlensing of quasar, e.g., Q2237+0305, has not been predicted. We stress here the significance of X-ray observations to elucidate the physics of the innermost parts of the disk, since X-rays specifically originate from a deep potential well. The observations allow us to assess the extension of hot regions on several AU scales and resultantly to deduce the mass of a central massive black hole. In this $Letter$, we propose to investigate quasar central structure by using X-ray microlensing of Q2237+0305. In section 2, we describe the method for resolving X-ray emission properties of the inner disk structure on a scale down to a few AU. In section 3, we calculate the disk spectra and light curves during microlensing. We here use realistic disk models: the optically-thick, standard disk (Shakura \& Sunyaev 1973) and the optically-thin, advection-dominated accretion flow (ADAF, Abramowicz et al. 1995; Narayan \& Yi 1995; see also Ichimaru 1977). | 98 | 4 | astro-ph9804251_arXiv.txt |
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9804 | cond-mat9804137_arXiv.txt | We study with Monte Carlo methods an ensemble of $c=-5$ gravity graphs, generated by coupling a conformal field theory with central charge $c=-5$ to two-dimensional quantum gravity. We measure the fractal properties of the ensemble, such as the string susceptibility exponent $\gamma_s$ and the intrinsic fractal dimensions $d_H$. We find $\gamma_s = -1.5(1)$ and $d_H = 3.36(4)$, in reasonable agreement with theoretical predictions. In addition, we study the critical behavior of an Ising model on a {\it quenched} ensemble of the \mbox{$c=-5$} graphs and show that it agrees, within numerical accuracy, with theoretical predictions for the critical behavior of an Ising model coupled {\it dynamically} to two-dimensional quantum gravity, provided the total central charge of the matter sector is $c=-5$. From this we conjecture that the critical behavior of the Ising model is determined solely by the average fractal properties of the graphs, the coupling to the geometry not playing an important role. | Randomness in statistical systems arises in a variety of situations and is a very rich and complex subject. Quenched randomness is frequently used in studying the role of impurities and inhomogeneities in real physical systems where the characteristic time-scale of the disorder is much longer than other dynamics of the system. Annealed randomness, on the other hand, arises naturally in studies of fluctuating geometries, such as two-dimensional quantum gravity or fluid membranes, where the disorder is dynamically modified by interaction between the geometry and matter fields living on the surfaces. For a statistical system coupled to random disorder, either in a quenched or annealed approach, the main question is to assess the effect randomness has on the critical behavior of the pure system. One prediction in this direction is the Harris conjecture \cite{harris} which states that randomness changes the values of critical exponents only if the specific heat exponent $\alpha$ of the pure system is positive. This conjecture has been studied in many models with quenched disorder, such as the $2d$ Ising model \cite{dots} (where the Harris criterion is ambiguous as $\alpha = 0$) and the Potts model \cite{bondpott}. For both models a change in the critical behavior is observed. All the above mentioned studies deal with weak disorder. More recently the critical behavior of systems on lattices with fractal structure very different from a flat surface has been investigated. Such systems arise naturally when matter, in the form of conformal field theories, is coupled to two-dimensional quantum gravity. These models can be studied either in a continuum formulation, by Liouville field theory, or using discretized approaches like, for example, models of dynamical triangulations, formulated either as matrix models or studied with numerical simulations. For these systems the disorder is, however, different from the one discussed above in that it is annealed, i.e.\ the models couple dynamically to fluctuations in the geometry. A remarkable degree of universality does emerge for models coupled to two-dimensional quantum gravity. Namely, the change in the critical behavior of the systems, and their back-reaction on the geometry, only depends on the total central charge of the matter sector. This manifests itself in the so-called KPZ scaling relation which describe how the scaling dimensions of conformal operators are changed by the interaction with gravity \cite{kpz}. Moreover, this universality also extents to the fractal structure of the surfaces, from which we derive the string susceptibility exponent $\gamma_s$ and the fractal dimension $d_H$. In view of this universality it is tempting to conjecture that the critical behavior of a particular system, when coupled to a fluctuating geometry, only depends on the (average) fractal structure of the surface. Details of the interaction between the system and the geometry, or the geometrical fluctuations, are not important as such --- they only serve the purpose of defining the average fractal geometry. If this conjecture is true it implies that how the average over disorder is performed, i.e.\ that the disorder is annealed, is not essential. In particular, predictions of the KPZ scaling relation for the change in the critical behavior should just as well apply to models with quenched disorder, {\it provided the quenched average is taken over the same ensemble of disorder as is generated in the annealed approach}. There are some recent simulations that have addressed the question of the critical behavior of spin models on a quenched ensemble of graphs generated by two-dimensional quantum gravity. Both the Ising model \cite{bhj} and the 10-state Potts model \cite{bjj} have been studied on an ensemble of pure gravity graphs ($c=0$). For the Ising model a critical behavior compatible with an Ising model coupled dynamically to gravity was found, although the accuracy of the results is not sufficient to rule out the conjecture discussed above. The goal of this paper is two-fold. First, we want to investigate the fractal geometry of two-dimensional quantum gravity coupled to a conformal field theory with central charge $c=-5$. More precisely, we want to determine the fractal dimension of the corresponding surfaces, using recently developed finite-size scaling methods \cite{hausd,janhaus} and to compare it to the (contradictory) theoretical predictions that exist \cite{anhaus1,anhaus2}. Second, we want to investigate the critical behavior of an Ising model on a quenched ensemble of $c=-5$ graphs and to compare it with predictions from Liouville theory, for the critical behavior of an Ising model coupled dynamically to two-dimensional quantum gravity, and to verify, or disprove, our conjecture about the effect of the disorder. Our motivation for choosing $c=-5$ is that both its predicted fractal structure and the critical behavior of the Ising model is substantially different from both a flat space and for a pure two-dimensional quantum gravity. This makes these different critical behavior easier to distinguish in numerical simulations. The paper is organized as follows: In Section~2 we study the fractal properties of a $c=-5$ conformal field theory coupled to two-dimensional gravity. We define the model in Section~2.1 and discuss the details of the simulations in Section~2.2. In Sections~2.3 and 2.4 we present our measurements of the string susceptibility exponent $\gamma_s$ and of the fractal dimension $d_H$. And in Section~2.5 we comment on how this particular ensemble of graphs differs from other types of graphs frequently used in studying disordered system. The second part of the paper deals with an Ising model on the $c=-5$ graphs in a quenched approach. In Section~3.1 we discuss the prediction from Liouville theory for the critical behavior of an Ising model coupled dynamically to two-dimensional gravity. In Section~3.2 we discuss details of the simulations and the observables we use to probe the critical behavior. In Sections~3.3 and 3.4 we determine the critical temperature of the Ising model and the corresponding critical exponents. Finally, in Section~4 we summarize and discuss our results. | The main results of the work presented in this paper can be summarized as follows: \begin{itemize} \item[({\it a})] The fractal dimension of surfaces, defined by a conformal field theory with central charge $c=-5$ coupled to two-dimensional quantum gravity, is $d_H = 3.36(4)$. This is in reasonable agreement with, and supports, the theoretical prediction Eq.~(\ref{andH2}), whereas it definitely rules out Eq.~(\ref{andH1}). \item[({\it b})] The critical behavior of an Ising model on a {\it quenched} ensemble of $c=-5$ graphs agrees well with the predictions, from the KPZ scaling relation, for an Ising model on an {\it annealed} ensemble of graphs with {\it identical} fractal properties. \end{itemize} The first result, especially combined with the recent simulations of $2d$--gravity for $c=-2$ \cite{cm2}, lends a strong support to Eq.~(\ref{andH2}) as a correct description of the fractal structure of two-dimensional quantum gravity for $c \leq 0$. This makes, however, its disagreement with numerical simulations in the region $0 < c \leq 1$ all the more surprising. What is it in derivation of Eq.~(\ref{andH2}) that breaks down for $c>0$? Or are the simulations dominated by finite-size errors and simulations of larger systems will eventually agree with Eq.~(\ref{andH2})? The result for the Ising model is even more interesting. As the theoretical predictions are obtained for an Ising model coupled dynamically to the disorder, this supports the conjecture put forward in the Introduction about the equivalence between annealed and quenched averages over disorder. That is, the only thing relevant for the critical behavior of the Ising model are the average fractal properties of graphs the spins ``see''. How the statistical average over graphs is performed, quenched or annealed, is not relevant. It is also worth noting that we can continuously change the average fractal properties of the graphs by changing the embedding dimension $D$ in Eq.~(\ref{eq213}). This allows a continuous interpolation between a flat surface and surfaces corresponding to pure gravity. If the prediction of Liouville theory, the KPZ formula, holds for all those models, this implies that the critical behavior of the Ising model should change continuously in the process. In the language of the renormalization group this implies a continuous line of fixed points, rather than isolated points. There are well known examples of this; the low-temperature phase of the two-dimensional $XY$--model or the critical line of the Ashkin-Teller model. But is this statement also true for very weak disorder? If we change the fractal dimension infinitesimally, from 2 to $2+\epsilon$, is that enough to change the critical behavior of the Ising model? Or, alternatively, does there exist some central charge $c^{\prime} < -5$ were the geometrical disorder is not strong enough and we always get the Onsager exponents? This point deserves further study. One could also look at the examples of weak disorder that have been studied recently, for example the site or bond-diluted Ising model, and ask if that kind of disorder can also be classified according to some average fractal properties of the lattices. And, moreover, if one could observe some kind of universality in the critical behavior, depending on the fractal structure, akin to what we have presented in this paper. In view of how dramatically the critical behavior of the Ising model changes on surfaces with such strong disorder, one might ask if such change could be observed in real physical systems. Possible candidates for such systems could be, for example, electrons trapped on the interfaces between two liquids, or on the surface of some porous material, were the surfaces had some well defined non-trivial fractal structure. As our results indicate, it is only the average geometry of the surfaces that is important for the Ising model, not its fluid nature or curvature fluctuations. Thus the relative time-scale between the interactions of the particles and the change in the geometry should be irrelevant. \vspace{20pt} \noindent {\bf Acknowledgments:} The work of G.T.\ was supported by the Humboldt Foundation. The work of P.B.\ was partially supported by KBN grants 2P03 B19609 and 2P03 B04412. | 98 | 4 | cond-mat9804137_arXiv.txt |
9804 | astro-ph9804100_arXiv.txt | I describe a general framework that could allow to understand the broad band spectra of blazars and lead to a unified picture of the emission from relativistic jets, in BL Lac objects as well as in flat spectrum, radio loud Quasars. The scheme serves as a useful basis to introduce and discuss some of the most interesting results so far obtained on Blazars with {\it Beppo}SAX. | The "blazar phenomenon" is due to the presence of relativistic flows (jets) emanating from the nuclei of active galaxies which are radio loud. The power to energize the radio lobes is transported in the jets. Blazars are the subset of radio loud AGN for which the relativistic jet happens to point at small angles to the line of sight. Since the radiation emitted by the jet is relativistically beamed into a narrow cone along the direction of motion, the aligned observer will receive a strongly enhanced flux. For a bulk Lorentz factor $\Gamma\simeq 10$ at an angle $\theta\simeq 1/\Gamma$ the flux enhancement factor is $10^3-10^4$. The evidence in favor of this picture has been accumulating and is now solid (e.g. \cite{pu} and refs therein) although the origin of the jets is poorly understood and their physical parameters are highly uncertain. The relativistically amplified, non thermal emission from high energy particles in the jet can account for the extreme properties of blazars concerning variability, polarization and energy distribution of the continuum, which extends from the radio to the gamma-ray band. Traditionally BL Lac objects, where no prominent emission lines are observed (with an upper limit of 5 \AA), were thought to represent a separate class perhaps more extreme than Quasar-like blazars. The latter include optically violently variable and highly polarized quasars (OVV, HPQ) or more generally Quasars with flat radio spectrum (FSRQ) indicating strong emission from the self absorbed core. It has become clear however that BL Lacs have on average lower luminosity than quasar like blazars (\cite{pado}) and that the distribution of emission line equivalent widths is continuous (\cite{scafal}). We will therefore in the following consider blazars as a single class of objects, implicitly assuming that, irrespective of the emission line properties which derive from the surrounding gas, the same physical mechanisms operate in relativistic jets over a wide range of luminosities. By studying the blazar continuum we expect to learn about the radiation mechanisms in the jets, about the processes of particle acceleration and energy transport along the jets and ultimately about their origin and evolution. \begin{figure*}[bt] \vspace{9pt} \psfig{file=sed_medie.ps,width=15.0truecm,height=11.5truecm,rheight=8.7truecm} \caption{\small\sf Average SEDs for the ``total blazar sample'' binned according to radio luminosity irrespective of the original classification. The overlayed dashed curves are analytic approximations obtained assuming that the ratio of the peak frequencies is constant and that the luminosity at the second peak is proportional to the radio luminosity (from \protect\cite{foss98} } \label{fig:sed_medie} \end{figure*} \section {The broad band spectra of blazars} The discovery by EGRET on board CGRO of copious $\gamma$-ray emission from blazars caused a "Renaissance" in this field. In fact it had been noted early on that the high density of relativistic electrons necessary to produce the observed compact synchrotron emission would lead to strong, even catastrophic inverse Compton radiation (\cite{hbs}). X-ray measurements were used to constrain the amount of inverse Compton emission allowed and to derive minimum values for the necessary beaming factors (e.g. \cite{gp93}). At present the $\gamma$-ray observations allow to measure the intensity and spectral shape of a component which contains a substantial fraction and in some cases the bulk of the emitted power leading to strong constraints on the physical parameters of the emitting region. It is clearly important, besides observing single objects, to try to derive general properties of the continuum and understand whether and how they differ for instance in BL Lac objects and FSRQ. It is especially interesting to discuss whether the gamma-ray emission is a general property of the whole class. We have recently addressed this problem (\cite{foss98}) collecting multifrequency data for three complete samples of blazars: the 2 Jy sample of FSRQs, the 1Jy sample of BL Lac objects and the sample of BL Lacs selected in the X-ray band from the Einstein Slew Survey. Systematic differences in the shape of the continuum in specific spectral bands among different subclasses of blazars were noted early on (e.g. \cite{gg86,imp,ww,smu,umu}. In particular we note that the percentage of objects detected with EGRET (100 MeV - 10 GeV) is significantly larger for the sample of FSRQ than for the two BL Lac samples (40 \% vs. 26\% and 17\% respectively. Nevertheless plotting the average SEDs as shown in Fig.~\ref{fig:sed_medie}, we can see that the shapes are "globally similar". In Fig.~\ref{fig:sed_medie} all blazars in the three complete samples were merged and grouped in luminosity classes irrespective of their original classification and the dashed lines drawn for comparison derive from an analytic parametric representation (\cite{foss98}). The main results of this work are the following: \begin{itemize} \item two peaks are present in all the SEDs \item the first peak occurs at lower frequencies for the highest luminosity objects \item the frequency at which the second peak occurs correlates with that of the first one. The dashed curves correspond to a constant ratio between the two peak frequencies. \end{itemize} For the most luminous objects the first peak is at frequencies lower than the optical band while for the least luminous ones the reverse is true. Thus highly luminous objects have a "red" (steep) IR to UV continuum while objects of lower luminosity have a bluer IR to UV continuum. For this reason and to recall intuitively the location of the peaks on the frequency axis we will briefly call "red" blazars the objects in the three highest luminosity classes and "blue" blazars those in the two lower luminosity classes. The present data suggest a continuous spectral sequence and no absolute separation between red and blue blazars. Considering the continuum of different objects in a fixed spectral range its shape changes systematically with luminosity along the sequence, as the peak frequency approaches and moves across the chosen frequency interval. In particular the X-ray spectrum becomes steeper and the gamma-ray spectrum (in the EGRET range) becomes flatter from "red" to "blue" blazars, as the two peaks march to higher frequencies. The different location of the gamma-ray peak can account for the different detection rates of BL Lacs and FSRQs by EGRET. Objects whose $\gamma$-ray emission peaks in the EGRET range are more easily detected. Recently ground based observations in the TeV range performed with Cherenkov telescopes have detected two of the X-ray brightest "blue" blazars (refs). We expect that with the progress in sensitivity many more will be detected giving access to the study of the highest energies from ground. A final comment concerns variability. It is interesting to note that the largest variability is usually observed close to or above each of the two peaks and is usually in the sense of a hardening of the spectrum at higher intensity. These statements are based mostly on observations at high energies (X-rays and gamma-rays) and concern a limited number of sources (e.g. \cite{umu}) therefore they should be considered as tentative suggestions rather than established facts. \section {Interpretation} It is generally thought that the first spectral component peaking at far infrared up to X-ray frequencies is due to synchrotron radiation. The spectra from the radio to the submm range most likely involve contributions from different regions of the jet with different self absorption turnovers. However, from infrared frequencies upwards the synchrotron emission should be thin and could be produced essentially in a single homogeneous region. Inverse Compton scattering of soft photons by the high energy electrons emitting the thin synchrotron radiation could be responsible for the second ( high frequency) component of the SED, peaking in the gamma-ray band. The soft photons could be the synchrotron photons themselves (SSC) or photons outside the jet (EC), possibly produced by an accretion disk or torus and scattered or reprocessed by the surrounding gas (e.g \cite{sbr}, \cite{umu} and refs therein). If the same region is responsible for the two spectral components then, irrespective of the nature of the seed photons, {\it the two peaks must derive from the same high energy electrons}. Therefore a change in the density and/or spectrum of those electrons is expected to cause correlated variability at frequencies close to the two peaks. In the SSC model the inverse Compton intensity is expected to vary more than the synchrotron one, approximately as the square of it in the simplest case while in the EC model one expects a linear relation (\cite{miami}). Measuring the two peaks simultaneously is thus the best means to determine the physical parameters of the emission region and studying the variability of the spectra around the peaks yields unique insight into the mechanisms of particle acceleration and energy loss in the jet. The variability correlation should enable to disentangle the contribution of different sources of seed photons (SSC vs. EC). The "spectral sequence" discussed above could be attributed to a systematic dependence of the critical electron energy (the break energy) and/or of the magnetic field on luminosity. Assuming that the beaming factors are not significantly different along the sequence, the trend in apparent luminosity is also a trend in intrinsic luminosity. In the SSC model the break energy of the electrons is univocally determined by the ratio of the frequencies of the two peaks and should therefore be approximately constant. "Red" blazars should then have lower magnetic field than "blue" blazars. Systematic model fitting of all the $\gamma$-ray detected blazars with sufficient multifrequency data suggest that as the magnetic energy density decreases the external photon energy density becomes important so that a smooth transition between the SSC and EC scenario takes place (\cite{gg98}). \section {SAX observations} The X-ray band is crucial for a discussion of the above problems in that the synchrotron and inverse Compton components which have different spectral shapes may both be relevant. Simultaneous observations over a broad energy range are required to disentangle the two mechanisms. The unique characteristics of the {\it Beppo}SAX instrumentation appear therefore ideal for blazar studies. Observations of bright blazars detected in $\gamma$--rays were proposed with the main aims of: \begin{itemize} \item determining the spectral shape up to the 100 keV range, thus exploring the connection between X-rays and gamma-rays \item studying the variability in relation with other wavebands, especially $\gamma$--rays. \end{itemize} \begin{figure*}[bt] \vspace{9pt} \psfig{file=sed_3c279.ps,width=15.0truecm,height=13truecm,rheight=11.2truecm} \caption{\small\sf Spectral energy distribution during the 1997 campaign, compared with previous ones. The 1997 data are from: X--rays = {\it Beppo}SAX data; $\gamma$--rays = Hartman, private comm.; R-band = Raiteri \& Villata, private comm.} \label{fig:sed_3c279} \end{figure*} In the following I will briefly mention and comment some of the most interesting results obtained so far. Besides 3C 273 (\cite{htv}) which is probably intermediate between a blazar and a "normal" quasar, two flat spectrum radio quasars, 3C 279 and PKS 0528+134, were observed (before June 1997) and found in a low intensity state. In the scheme presented above these are "red" blazars. I will discuss here the first source (also \cite{mtv}), while for the second one I refer to \cite{geltv}. Finally I will consider results on "blue" blazars ( Mrk 421, 1ES 2344+514 and Mrk 501 (see the contributions in this volume \cite{ftv}, \cite{gitv} and \cite{gtv} respectively). It is important to remember that in "red" blazars the X-ray emission represents the lower energy end of the inverse Compton emission, while for "blue" blazars it is the high energy end of the synchrotron emission. \subsection {3C 279} The X-ray spectrum measured with {\it Beppo}SAX in January 1997, the simultaneously measured optical flux and the quasi simultaneous gamma-ray flux (Hartman, private communication) are shown in Fig.~\ref{fig:sed_3c279} together with other simultaneously measured SEDs obtained at other epochs: the high state observed in June 1991 (\cite{hart}), the low state observed in January 1993 (\cite{m94}) and the preflare and flare states observed in January - February 1996 (\cite{wehrle98}). At the epoch of the {\it Beppo}SAX observations the $\gamma$--ray flux was a factor 6 and 20 weaker respectively than measured in June 1991 and early February 1996. In X-rays the amplitude is smaller but there is a good correlation between the X-ray and gamma-ray fluxes especially in the 2-10 keV band (see Fig. 3 in Maraschi et al. this volume). Note that the fluxes at 1 keV measured by {\it ROSAT} and {\it Beppo}SAX in the 1993 and 1997 low states are similar However the spectrum measured by {\it Beppo}SAX is significantly flatter providing a good connection with the higher $\gamma$--ray flux in 1997. The simultaneity (within one day) of the X-ray (XTE) and $\gamma$--ray peaks during the 1996 flare suggests that the X-ray to gamma-ray emission originates in a single region and that the spectrum hardens with increasing intensity. It is possible that the $\gamma$-ray peak in the SED moves to higher energies at the flare peak. The situation is much more complex at lower energies. Although there is still a general correlation of the IR-optical-UV fluxes with the gamma-ray intensity on long timescales, the flux variation at optical wavelengths corresponding to the rapid 1996 flare is quite small. Note also that for 3C 279 the (presumed) peak of the synchrotron component falls in an unexplored region of the spectrum, between $10^{12}$ and $10^{14}$ Hz. In the SSC model one expects that the inverse Compton emission varies with the square of the amplitude of the synchrotron emission due to the same electrons (e.g. at the two peaks of the SED). This is compatible with the long term variations but not with the strong rapid flare observed in 1996, where the amplitude in gamma-rays was larger than the square of the optical one. On the other hand if the seed photons for the inverse Compton process are external to the jet, they should not be rapidly variable and the inverse Compton emission is expected to vary linearly with the synchrotron one. Thus neither of the two "simple" models can adequately account for the multifrequency variability behaviour. A possible way out is that the seed photons derive from backscattering and/or reprocessing of radiation produced in the jet by gas clouds closely approaching the jet itself (\cite{ggmadau}). This model is attractive and needs to be studied in more detail. Another possible way out is that the region emitting the synchrotron radiation is inhomogeneous so that the observed variability is diluted by a more stationary component. \subsection {Blue blazars} The X-ray emission from these objects has been observed to vary dramatically on short timescales at least in the brightest prototypes, PKS 2155-304 and Mrk 421. This can be understood recalling that the X-ray emission represents the high energy end of the synchrotron component: it is therefore due to radiation from the highest energy electrons which have the shortest lifetimes and can vary very rapidly. A continuous acceleration mechanism must be responsible for maintaining in the source particles which have lifetimes of the order of hours. However the injection mechanism may be continuous only in an average sense or at a low intensity level and episodes of increased injection rate may occur causing variability. The study of flares and of spectral variability associated with them gives direct information on the spectra of the freshly injected/accelerated particles and their subsequent decay to a state of quasi-equilibrium. The photons upscattered through the inverse Compton process by the highest energy electrons reach TeV energies so that the "bluest" and brightest sources can be detected from ground based Cherenkov telescopes. This is the case up to now for three objects: Mrk 421, 1ES 2344+514 and Mrk 501. All of them were observed with SAX in the first half of AO1. The simultaneous observation of a source in X-rays and at TeV energies should allow to determine unambiguously the energy of the radiating electrons, the beaming factor, the magnetic field and the energy density of the seed photons. In some cases it may be necessary to take into account that scattering will occur in the Klein Nishina regime. \begin{figure*}[bt] \vspace{9pt} \psfig{file=sed_mkn421.ps,width=15.0truecm,angle=270,height=13truecm,rheight=11.2truecm} \caption{\small\sf The May 97 data from {\it Beppo}SAX are compared with the energy distributions measured in 1994 (\protect\cite{macomb}} \label{fig:sed_mkn421} \end{figure*} \subsection {Mrk 421} This source has been repeatedly observed with ASCA and at other relevant wavelengths including TeV observations. (\cite{macomb}) The {\it Beppo}SAX observations of Mrk 421 (Fossati et al. this volume) show a decay between two intensity states closely similar to those previously observed with ASCA. The higher state shows a flatter spectrum indicating that the emission peaks at higher energies. In Fig.~\ref{fig:sed_mkn421} the {\it Beppo}SAX data are compared with the ASCA and TeV data presented by \cite{macomb}. Mastichiadis and Kirk \cite{mk}have shown that this behaviour could be the result of an injection mechanism in which the maximum energy of the injected particles increases. The two "states" would therefore represent equilibrium states for injection spectra extending to different maximum energies. The fact that the spectral variability observed in 1994 is closely reproduced in the {\it Beppo}SAX observations indicates that the involved region is the same (same physical parameters). In addition SAX detected the source at higher energy with the PDS. The preliminary analysis yields a flat spectrum at high energies suggesting that the inverse Compton component becomes dominant in this band. The variation in the PDS is much smaller than in the MECS which is consistent with attributing the hard emission to IC from electrons of much lower energies. An extensive campaign for simultaneous ASCA and TeV observations will take place in the spring of 1998, to which {\it Beppo}SAX observations could add significantly. \subsection {1ES 2344+514, Mrk 501} This source is not as well studied as Mrk 421 but shows a similar, more extreme behaviour (\cite{gitv}). The emission peak was in the medium X-ray range in December 96 and shifted to 20-50 keV during the flare of a factor 2 observed between December 3-7 96. The most extraordinary spectral variation was found in Mrk 501 (\cite{pian98}, \cite{gtv}). The source was in a state of strong activity at TeV energies although the 1 keV flux had not increased dramatically. However compared to previous observations {\it the X-ray spectrum had changed dramatically} being flatter than or close to 1 up to 100 keV. In correspondence to a large TeV flare , the 1 to 100 keV spectrum hardened still indicating that the emission peak was at or above 100 keV, while past multifrequency measurements showed it to be below 1 keV. Thus in Mrk 501 the shift of the peak frequency was more than a factor 100. Some theoretical implications are discussed by Ghisellini (\cite{gtv}). This unprecedented behaviour may be less uncommon than judged at first sight. In fact in the medium X-ray band the intensity behaviour was not exceedingly dramatic and good spectral capabilities up to the hard X-ray band, necessary to reveal the phenomenon, have only recently become available with {\it Beppo}SAX. It is interesting to mention that a spectral survey of BL Lacs selected in relatively soft X-rays yields evidence of some hard X-ray spectra (\cite{wolter}). Sources which have more or less permanent emission peaks in the hard X-ray range may also exist and have gone undetected so far. | The results from the first part of the {\it Beppo}SAX AO1 Core Program on Blazars are extremely exciting. Bright $\gamma$-ray blazars can be usually detected up to 50 - 100 keV allowing detailed study of the synchrotron and inverse Compton emission and of the correlations between their temporal and spectral variability. The observations presented above suggest and support: \begin{itemize} \item the correlation between the X--ray (1 keV) and $\gamma$--ray (0.1-10 GeV) fluxes, which now holds over a period of 6 years and about a factor 30 in flux change for $\gamma$-rays (3C 279 -- PKS 0528 + 154) \item the smaller amplitude of variability of the synchrotron component compared to the high energy one. \item the synchrotron emission peak seems to shift systematically to higher energies during flares (Mrk 421, 1ES 2344+541, Mrk 501) \end{itemize} These results will undoubtedly stimulate new observations and new theoretical approaches. In particular genuinely time dependent models for the acceleration of particles and the spectral evolution of the emitted radiation are needed. | 98 | 4 | astro-ph9804100_arXiv.txt |
9804 | astro-ph9804336_arXiv.txt | In this paper we present a detailed study of the radio galaxy J1324$-$3138, located at a projected distance of 2$^{\prime}$ from the centre of the Abell cluster of galaxies A3556, belonging to the core of the Shapley Concentration, at an average redshift z=0.05. We have observed J1324$-$3138 over a wide range of frequencies: at 327 MHz (VLA), at 843 MHz (MOST), and at 1376 MHz, 2382 MHz, 4790 MHz and 8640 MHz (ATCA). Our analysis suggests that J1324$-$3138 is a remnant of a tailed radio galaxy, in which the nuclear engine has switched off and the radio source is now at a late stage of its evolution, confined by the intracluster gas. The radio galaxy is not in pressure equilibrium with the external medium, as it is often found for extended radio sources in clusters of galaxies. We favour the hypothesis that the lack of observed polarised radio emission in the source is due to Faraday rotation by a foreground screen, i.e. the source is seen through a dense cluster gas, characterised by a random magnetic field. An implication of the head-tail nature of the source is that J1324$-$3138 is moving away from the core of A3556 and that possibly a major merging event between the core of A3556 and the subgroup hosting J1324$-$3138 has already taken place. | Extended radio emission associated with galaxies in clusters is often characterised by morphologies which reflect the interaction between the radio emitting plasma and the local environment in the cluster. Head-tail sources are usually associated with non-dominant cluster galaxies moving at a considerable speed within the cluster. Their morphologies are then explained as the result of ram pressure exerted by the intergalactic medium on the double sided radio emission (see for example O'Dea \& Owen 1985a and 1985b, and Owen 1996, for a recent review). Wide-angle tail radio galaxies, on the other hand, are more difficult to account for with the above mentioned model, since they are usually associated with massive and dominant cluster galaxies, with much lower peculiar velocities with respect to the cluster mean. Beyond ram pressure, it is now accepted that large flows of hot gas could provide a wind within clusters of galaxies able to bend straight jets into wide-angle tail morphologies (Owen 1996 and references therein). \noindent The study of extended galaxies in clusters is important for a variety of reasons. The morphology and the direction of the extension may give important information on the dynamics of the galaxy, such as, for example, the direction of the motion projected on the plane of the sky. Furthermore, the non-thermal pressure in the tails of radio emission can be compared to the thermal pressure exerted by the intracluster gas, in those cases where X-ray data are available to provide estimates of the temperature and pressure. This is crucial for studying the interaction between the radio emission and the external gas, and for deriving information on the evolution of radio sources in clusters of galaxies, as well as the influence of the cluster dynamics (such as, for example, merging processes) on the radio properties of the cluster. Last but not least, the observed polarisation properties of the radio emission may give information on the intracluster magnetic field and its structure. \medskip In this paper we present a detailed study of the extended radio galaxy J1324$-$3138 (RA$_{J2000} = 13^h24^m01^s$, DEC$_{J2000} = -31^{\circ}38^{\prime}$), located in the central region of the Abell cluster A3556. It was first observed in a radio survey of the clusters of galaxies in the Shapley Concentration core carried out at 843 MHz with the Molonglo Observatory Synthesis Telescope (MOST) and at 1376 MHz with the Australia Telescope Compact Array (ATCA) (Venturi et al. 1997, hereinafter Paper I). This work is part of a larger project whose aim is to study the radio/optical properties of the clusters in the core of the Shapley Concentration, in particular the chain formed by A3556-A3558-A3562 (Venturi et al. 1998), both from a statistical point of view and through a detailed analysis of the physical properties of the extended radio galaxies in these clusters. In Figure 1 the superposition of the radio isophotes on the Digitised Sky Survey shows that the radio component located at the south west extremity of the extended radio emission is coincident with the nucleus of the 15.6 magnitude galaxy \#5975 in the COSMOS catalogue (RA$_{J2000}$ = $13^h23^m57.5^s$, DEC$_{J2000}$ = $-31^{\circ}38^{\prime}45^{\prime\prime}$). Its radial velocity velocity v = 15054 km s$^{-1}$ (Stein 1996) establishes that it belongs to A3556 ($<v> = 14357$ km s$^{-1}$, Bardelli et al. 1998). This, coupled with the fact that only a few very faint optical objects fall within the envelope of the radio emission, led us to the conclusion that J1324$-$3138 is an extended, possibly head-tail, radio galaxy located in the vicinity of the cluster centre (see Paper I). In Section 2 we present the observational data. In Section 3 the morphology of the source is described and analysed, and in Section 4 a detailed study of the synchrotron spectrum is carried out. The nature of the source, its relation to the cluster of galaxies A3556 and its implications for cluster merging and formation is discussed in Section 5. Throughout the paper we use a Hubble constant of H$_0$ = 100 km s$^{-1}$Mpc$^{-1}$. At the redshift of the cluster this implies that 1$^{\prime\prime}$ = 0.67 kpc. \begin{figure} \epsfysize=8.5cm \epsfbox{FIG1.PS} \caption[]{ 4790 MHZ radio isophotes of the extended radio galaxy J1324$-$3138 and of the nearby radio galaxy J1324$-$3140, associated with the dominant cD galaxy in the centre of A3556 (see Paper I), superimposed on the DSS optical image. The resolution of the image is $20^{\prime\prime} \times 10^{\prime\prime}$, p.a. $0^{\circ}$.} \end{figure} \vskip 1.0truecm \noindent | We have presented observations of the radio galaxy J1324$-$3138, located in the central region of the Abell cluster A3556, over a wide range of frequencies and resolutions. We can briefly summarise our results as follows: \noindent {\it (a)} J1324$-$3138 is an example of a {\it remnant} of a radio galaxy, i.e. a source in which the engine of the radio emission has switched off. The evolution of the radio emission is presently dominated by synchrotron losses. \noindent {\it (b)} The radio source is not in pressure equilibrium with the intracluster gas. In particular it is underpressured. \noindent {\it (c)} We suggest that the lack of polarisation in the source is due to the presence of an intervening Faraday screen, i.e. a cluster scale magnetised medium, as it is now often observed in clusters of galaxies, which depolarises the radio emission. \noindent {\it (d)} Under the hypotesis that cluster mergers influence the radio emission of a galaxy, the properties of J1324$-$3138, coupled with the peculiarities of A3556 at radio and optical wavelengths (Paper I and Bardelli et al. 1998), suggest that the core of A3556 and the subgroup hosting J1324$-$3138 have already undergone a major merging event. \vskip 1.0truecm \noindent {\bf Acknowledgments} We wish to thank D. Dallacasa for his suggestions and discussion while this work was carried out, and R. Fanti for careful reading of the manuscript. We are grateful to S. Ettori for providing unpublished results, and to P. Stein for giving us the spectrum of J1324$-$3138. T.V. acknowledges the receipt of two grants from CNR/CSIRO (Prot. n. 119816 and Prot. n. 088864). The MOST is operated by the University of Sydney, with support from the Australian Research Council. The Australia Telescope Compact Array is operated by the CSIRO Australia Telescope National Facility. The National Radio Astronomy Observatory (NRAO) is operated by Associated Universities, Inc., under contract with the National Science Foundation. This research has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, Caltech, under contract with the National Aeronautics and Space Administration. | 98 | 4 | astro-ph9804336_arXiv.txt |
9804 | astro-ph9804046_arXiv.txt | The CAT (Cherenkov Array at Th\'emis) imaging telescope, equipped with a very-high-definition camera (546 fast phototubes with $0.12^{\circ}$ spacing surrounded by 54 larger tubes in two guard rings) started operation in Autumn 1996 on the site of the former solar plant Th\'emis (France). Using the atmospheric Cherenkov technique, it detects and identifies very high energy $\gamma$-rays in the range $250\:{\mathrm GeV}$ to a few tens of TeV. The instrument, which has detected three sources (Crab nebula, Markarian 421 and Markarian 501), is described in detail. | 98 | 4 | astro-ph9804046_arXiv.txt |
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9804 | astro-ph9804270_arXiv.txt | We report results of a \rosat\ High-Resolution Imager (HRI) observation of the X-ray error box given by the \sax\ Wide Field Camera for the gamma-ray burst that occurred on 1997 February 28. The observation started 10 days after the burst and ended three days later, with a total exposure of 34.3~ks. An X-ray source was detected within the 3$'$ WFC error box and its position determined with a 10$''$ radius accuracy. The source position is in the \sax\ Narrow Field Instrument source error box and is coincident (to within 2$''$) with the optical transient associated with GRB970228. This is the most precise position obtained for an X-ray afterglow and confirms that the X-ray and optical afterglows have the same origin. We present the 0.1--2.4~keV combined HRI and \sax\ Low-Energy Concentrator Spectrometer decay light curve which can be well fit with a power-law. The decay is consistent with that measured at higher energies (2--10~keV) with the \sax\ Medium-Energy Concentrator Spectrometer. | Observations of celestial Gamma-Ray Bursts (GRB) performed over the last 25 years had not, until recently, succeeded in finding counterparts in other wavelength bands. The ability of the \sax\ satellite to provide arc minute precision positions (Piro et al. 1998) and to observe these positions within hours of the GRB changed this situation in 1997 when the X-ray afterglow of GRB970228 was measured (Costa et al 1997a). The burst was detected (Costa et al. 1997a) with the Gamma-Ray Burst Monitor (GRBM) (40--70~keV, Frontera et al. 1997a) on 1997 February 28.123620 UT and also detected in the 1.5--26 keV energy range by one of the two Wide Field Cameras (WFC No. 1) aboard the same satellite (Jager et al. 1997). Its position was determined with an error circle of 3~arcmin (3$\sigma$) radius, centered on $\alpha_{2000}\,=\,05^h01^m57^s$, $\delta_{2000}\,=\,11^\circ46'24''$. Eight hours after the GRB trigger, from February 28.4681 to February 28.8330 UTC, the Narrow Field Instruments (NFI) on board \sax\ (Boella et al. 1997a) were pointed to the WFC error box. An X-ray source, SAX J0501.7+1146, was detected (Costa et al 1997b) in the field of view of both the Low Energy (0.1--10~keV) and Medium Energy (2--10~keV) Concentrators Spectrometers (LECS and MECS) (Parmar et al. 1997; Boella et al. 1997b). The source position ($\alpha_{2000}\,=\,05^h01^m44^s$, $\delta_{2000}\,=\,11^\circ46'42''$) is consistent with the GRB error circle. The source was again observed about three days later, from March 3.7345 to March 4.1174. During this observation, the 2--10~keV source flux had decreased by about a factor 20, while in the 0.1--2~keV energy range the source was not detected. Following the discovery of the GRB, searches for radio and optical counterparts to GRB970228 were conducted with most of the ground based telescopes in the northern hemisphere. Groot et al. (1997) reported the discovery of an optical transient at a position ($\alpha_{2000}\,=\,05^h01^m46.70^s$, $\delta_{2000}\,=\,11^\circ46'53.0''$), consistent with both the \sax\ WFC and NFI error boxes and with the long baseline timing \ulysses/\sax\ and \ulysses/\wind\ error annuli, of 31$''$ and 30$''$ half-width, respectively (Hurley et al. 1997; Cline et al. 1997). While the association of the transient X-ray source with the afterglow of GRB970228 was compelling on the basis of the properties of its decay curve when extrapolated backwards to the burst time (Costa et al. 1997c), the association of the optical transient with the burst afterglow was less strong. In spite of the positional consistency and temporal behaviour of the optical transient, it was not possible to exclude the possibility that the optical transient was unrelated to the GRB (see discussion by van Paradijs et al. 1997), like in the case of the radio source discovered in the earliest error box of GRB970111, which showed a time behaviour consistent with that expected from radio afterglows of GRBs, but later resulted to be unrelated tho the burst (Frail et al. 1997). The \rosat\ satellite, thanks to its HRI focal plane instrument, offered the possibility of imaging the X-ray afterglow at 10$''$ angular resolution (David et al. 1997). A Target of Opportunity observation was thus requested and obtained. Here we report on results of that observation and its consequences. Preliminary results were already previously reported (Frontera et al. 1997b). | The \rosat\ HRI observation of the \sax\ WFC error circle of GRB970228 clearly shows the presence of a new X-ray source. Its position within the error box of the \sax\ source, the low probability of a chance coincidence ($\sim 1 \times 10^{-3}$) and the better imaging capabilities of the \rosat\ HRI compared to the \sax\ NFI, indicate that the \rosat\ source and the \sax\ source are the same object. The source position derived from the \rosat\ observation is the most precise position of a GRB X-ray afterglow obtained thus far. Its position is also coincident with the optical transient associated to GRB970228 within 2$''$. This result confirms that X-ray source and the optical transient are the same object. The X-ray source hows a 0.1--2.4 keV decline according to a power law decline with index $\alpha \,=\, 1.50^{+0.23}_{-0.35}$ for at least 13 days. This slope is fully consistent with that estimated in the 2--10 keV energy band (Costa et al. 1997c) and is marginally consistent with that reported by Fruchter et al. (1997, 1998) in the optical band. Thus it appears that from the X-ray to the optical band the GRB afterglow has the same decline law. \begin{figure} \epsfig{figure=grb282_l_tot.ps,height=8.5cm,width=8.5cm,angle=0} \caption{As in Fig. 2, but extrapolated to the first second from the burst onset. The two arrows on the left delimit the time interval, that corresponds to the GRB last three pulses, when the X-ray afterglow is expected to start (see text). } \label{figure:decayb} \end{figure} | 98 | 4 | astro-ph9804270_arXiv.txt |
9804 | astro-ph9804264_arXiv.txt | Including nucleon--nucleon correlations due to both Fermi statistics and nuclear forces, we have developed a general formalism for calculating the charged--current neutrino--nucleon absorption rates in nuclear matter. We find that at one half nuclear density many--body effects alone suppress the rates by a factor of two and that the suppression factors increase to $\sim$5 at $4\times10^{14}$ g cm$^{-3}$. The associated increase in the neutrino--matter mean--free--paths parallels that found for neutral--current interactions and opens up interesting possibilities in the context of the delayed supernova mechanism and protoneutron star cooling. | The neutrino absorption and scattering opacities in the post--shock core of a supernova, in which nuclei are largely disintegrated into nucleons, determine the duration, spectrum, and flavor distribution of the emerging neutrino pulse. It has been known for some time that the interactions among nucleons in the denser regions can change these opacities significantly, but to date there has been no comprehensive treatment given in the literature and present calculations of the complete supernova process do not include the effects of interactions on the opacities. The neutrino--matter interaction rates can be related to the space-- and time--dependent correlations among the set of density operators for the separate nuclear constituents (to find the Gamow--Teller parts we must consider separate spin--up and spin--down densities). In the case of neutral--current interactions \cite{BS}, there is an instructive limit, which also provides an estimate of the effects, in which the combined limits of large nucleon mass and small neutrino energy allow the use of long--wavelength limits of equal--time correlation functions, in turn expressible in terms of the second derivatives of an energy density functional with respect to various densities. This approach is the direct multichannel generalization of the familiar results for light scattering from the thermal density fluctuations in a fluid, where it is the compressibility that determines the long--wavelength opacity, and it was used in references \cite{s2Ray} and \cite{iwa} to find significant reductions of neutral--current opacity in certain regions. In Burrows \& Sawyer \cite{BS}, an approach based on ring graphs was used to encompass these results and to extend them to domains in which the equal--time and long--wavelength limits are not clearly applicable. The use of the equal--time and long--wavelength limits to express correlation functions in terms of static susceptibilities cannot be extended to the charged--current interactions when there is a large chemical potential difference between protons and neutrons. Furthermore, there do not exist in the present literature systematic estimates of the effects of interactions on the charged--current opacities for electron neutrinos. In the present work, we give a theoretical framework for addressing these opacities, based on summing ring graphs, together with the results of calculations with input parameters taken from the current phenomenology of nuclear matter. | We have developed a new formalism for incorporating the effects of many--body correlations on the charged--current rates of neutrino--matter interactions. This formalism reveals that these rates are considerably suppressed in the densest regions of protoneutron stars and supernova cores. Assuming that the nucleons are non--relativistic, our formalism incorporates the full kinematics of the interaction, Pauli blocking by final--state nucleons (protons), and correlation due to nucleon--nucleon interactions. We have employed the ring approximation (RPA) and assumed the near--validity of Fermi Liquid Theory. It would desirable to include ladder diagrams and to perform the calculations in the context of a better numerical method for solution of the nuclear equation of state (EOS), since the solution of the EOS is intimately related to the derivation of the scattering/absorption rates. However, those who perform detailed nuclear EOS calculations and address many--body correlations in nuclear matter do not as yet provide the requisite spin and density structure functions, even for the static case. These results for charged currents, when combined with the results from Burrows \& Sawyer \cite{BS} for neutral currents, strongly suggest that energy and lepton number will leak from supernova cores at a rate that is higher than heretofore estimated. This implies that the neutrino luminosities during the epoch after bounce for which the inner core is the major energy source ($> 0.5 - 1.5$s) will be enhanced, perhaps by as much as 50\% \cite{BS}. The consequences of this increased transparency for the neutrino--driven supernova explosion mechanism \cite{bhf} may be interesting, but have yet to be clarified. | 98 | 4 | astro-ph9804264_arXiv.txt |
9804 | astro-ph9804052_arXiv.txt | ACO 3627 is a rich, nearby cluster of galaxies at the core of the Great Attractor. At the low galactic latitude of $b = -7.2\deg$ the galactic extinction is significant. Nevertheless, its proximity makes it a prime target for studies of environmental effects on its cluster members. Here, we report on a multi-wavelength study of a Seyfert 1 galaxy at 30 arcmin from the centre of ACO 3627. Its Seyfert nature was discovered spectroscopically and confirmed in X-rays. We have obtained B$_{\rm J}$ and R$_{\rm C}$ CCD photometry as well as J, H, K and L aperture photometry at the SAAO, low and high resolution spectroscopy (ESO and SAAO), 21 cm line observations (Parkes Observatory) and X-ray ROSAT PSPC data. The Seyfert 1 galaxy is of morphology SBa(r). It has a nearby companion (dS0) but shows no signs of interaction. A consistent value for the galactic extinction of A$_{\rm B}$ = 1.6 mag could be determined. The nucleus of the Seyfert is very blue with a strong (B$_{\rm J}$ -- R$_{\rm C}$) colour gradient in the inner 2.5 arcsec. The extinction-corrected near-infrared colours of WKK 6092 are typical of a Seyfert 1 and the X-ray spectrum conforms to the expectation of a Seyfert as well. The galaxy has a very low \HI\ flux. This could be explained by its morphology, but also -- due to its very central position within the rich Norma cluster -- to ram pressure stripping. | Dust and stars in the Milky Way obscure a large fraction of the extragalactic sky, creating a ``Zone of Avoidance'' (ZOA) in the distribution of galaxies. In an effort to reduce the size of the ZOA and thus coming closer towards an all-sky distribution of galaxies, we have embarked on a deep optical galaxy search behind the southern Milky Way (Kraan-Korteweg \& Woudt 1994). This has led to the recognition that ACO 3627 (Abell \etal\ 1989), also called the Norma cluster after the constellation it is located in, is a massive, nearby cluster of galaxies at the core of the Great Attractor (GA) $(\ell,b,v) = (325\deg, -7\deg, 4882$ \kms) (Kraan-Korteweg \etal\ 1996). The Norma cluster appears to be the central, dominant component of a ``great wall''-like structure and would be the most prominent overdensity of galaxies in the southern sky, were it not obscured by the Milky Way (Woudt \etal\ 1997). Recent observations of the Norma cluster with the ROSAT PSPC have confirmed the massive nature of ACO 3627; it is the 6$^{th}$ brightest cluster in the ROSAT sky (B\"ohringer \etal\ 1996). The X-ray contours furthermore suggest the existence of a subcluster. The merging scenario is independently supported by the radio continuum emission of the central cD galaxy PKS1610-608. The emission from this wide-angle-tail (WAT) radio source (Jones \& McAdam 1992) seems to encircle the X-ray subcluster (\cf\ Fig.~3 of Kraan-Korteweg \etal\ 1997) and is indicative of a strong motion of the cluster gas due to the ongoing merging process (Jones \& McAdam 1996, Burns \etal\ 1994). Roughly 30{\arcmin} from the centre of this cluster -- taken as the central cD galaxy PKS1610-608 -- we have identified a Seyfert 1 galaxy. It is a member of ACO 3627. In the following sections we describe the various observations of this galaxy: the discovery of the galaxy in section 2, the multicolour photometry obtained at the South African Astronomical Observatory (SAAO) in section 3, the spectroscopy obtained at the European Southern Observatories (ESO) and the SAAO in section 4, the \HI\ observations obtained with the 64m radio telescope of the Parkes Observatory of the Australian Telescope National Facility (ATNF) in section 5, and the X-ray data from ROSAT PSPC observations in section 6. The results are summarized and discussed in the last section. | We have observed the Seyfert galaxy WKK 6092 at different wavelengths. The resulting data are summarized in Table~\ref{seytab}. \begin{table} \caption{Observational parameters of WKK 6092} \label{seytab} \begin{tabbing} Apparent blue (IIIaJ)(m$_{B}$) magn \= \kill {\bf Coordinates:} \> \\ R.A. (1950) \> $16^{h} 07^{m} 32.7^{s}$ \\ DEC. (1950) \> $-60^{\circ} 30' 11''$ \\ Galactic Longitude \> $325.20^{\circ}$ \\ Galactic Latitude \> $-6.74^{\circ}$ \\ \vspace{5mm} {\bf Properties:} \> \\ Hubble Type \> SBa(r) \\ Dimensions (a x b) \> 56'' x 47'' \\ Ellipticity (1-b/a) \> 0.11 \\ Inclination \> 28$\deg$ \\ Position Angle \> 96$\deg$ \\ \vspace{5mm} {\bf Photometry:} \> \\ $B_{\rm J}$ (IIIaJ) \> 14.7 $\pm$ 0.5 mag \\ $B_{25}$ (CCD) \> 14.96 $\pm$ 0.09 mag \\ $B_{\rm T}$ (CCD) \> 14.88 $\pm$ 0.13 mag \\ $R_{24}$ (CCD) \> 13.38 $\pm$ 0.12 mag \\ $R_{\rm T}$ (CCD) \> 13.30 $\pm$ 0.14 mag \\ J$_{\rm c}$ \> 12.91 $\pm$ 0.03 mag \\ H \> 11.98 $\pm$ 0.03 mag \\ K \> 11.51 $\pm$ 0.03 mag \\ L \> 10.60 $\pm$ 0.20 mag \\ \HI\ flux \> 0.93 Jy \kms \\ X-Ray (0.5--2.0 keV): \> \\ \hspace{0.25cm} Flux \> $1.05 \pm 0.15 \cdot 10^{-12}$ erg s$^{-1}$ \\ \> \hspace{2.65cm} cm$^{-2}$ \\ \hspace{0.25cm} ${\cal L}_X$ \> $1.2 \pm 0.17 \cdot 10^{42}$ erg s$^{-1}$\\ \vspace{5mm} {\bf Galactic Extinction (A$_B$):} \> \\ from HI \> 1.5 mag \\ from Balmer decrement \> $\la$ 1.7 mag \\ from X-ray \> 1.6 mag \\ \vspace{5mm} {\bf Heliocentric velocity:} \> \\ MEFOS \> 4711 $\pm$ 30 \kms \\ S.A.A.O. \> 4688 $\pm$ 40 \kms \\ Parkes 64-m \> 5012 $\pm$ 5 \kms \\ \hspace{0.25cm} \DVF \> 88 \kms \\ \hspace{0.25cm} \DVT \> 97 \kms \\ \end{tabbing} \end{table} WKK 6092 and its neighbour are both members of ACO 3627. They have similar redshifts but show no indications of interaction. The morphology of both galaxies do not seem distorted (\cf\ Fig.~\ref{ccdim}). The Seyfert has a very blue nucleus, a distinct bar and a ring superimposed on an otherwise smooth disk. An upper limit for the Galactic foreground extinction in the line of sight of the Seyfert galaxy can be set at A$_{\rm B}$ = 1.6 mag. This was determined by three different methods, the Balmer decrement in the optical spectrum, the fitting of an absorbed power low to the X-ray spectrum and the Galactic \HI\ column densities. All give a consistent value of the foreground extinction. A minor fraction of the extinction is intrinsic to the galaxy. The extinction corrected near-infrared colours of WKK 6092 are typical of a Seyfert 1 and are in agreement with well known Seyfert 1's such as NGC 1566 (Glass and Moorwood 1985). The X-ray sepctrum is also consistent with the standard expectation for this object. At the adopted cluster distance of R = 93 $h_{50}^{-1}$ Mpc, the absolute magnitude (corrected for the galactic extincton) is M$_{B_T}^o = -21.52$. The \HI\ and total mass is $1.9 \cdot 10^{9} {\cal M}_{\odot}$ and $30 \cdot 10^{9} {\cal M}_{\odot}$, respectively. The Seyfert is at a projected distance of 0.8 $h_{50}^{-1}$ Mpc from the cluster centre and the \HI\ content of the galaxy might be influenced by interactions with the Inter Cluster Medium due to processes like ram pressure stripping (Cayatte \etal\ 1990). The \HI\ content is in fact quite low. This is, however, not inconsistent with the expectation for a barred early-type spiral. Despite the difficulties in analysing data of an object deep within the Milky Way, all data concerning the here investigated Seyfert galaxy WKK 6092 at 30 arcmin from the centre of the rich cluster ACO 3627 correspond to the standard characteristics of a Seyfert 1 galaxy. | 98 | 4 | astro-ph9804052_arXiv.txt |
9804 | astro-ph9804322_arXiv.txt | We report the {\em first} unambiguous detection of the host galaxy of a normal radio-quiet QSO at high-redshift in $K$-band. The luminosity of the host comprises about 35\%\ of the total $K$-band luminosity. Assuming the average colour of QSOs at $z\approx2$, the host would be about 5 to 6~mag brighter than an unevolved $L_*$ galaxy placed at $z\approx2$, and 3 to 4~mag brighter than a passively evolved $L_*$ galaxy at the same redshift. The luminosity of the host galaxy of the QSO would thus overlap with the highest found in radio-loud QSOs and radio-galaxies at the same redshift. | Recent evidence that the cosmological evolution of the density of star formation in the Universe (Madau et al. 1996) follows closely the QSO density evolution (Boyle \& Terlevich 1998) emphasizes the need to study the kinds of galaxies that host Active Galactic Nuclei in order to understand the link between star-formation and nuclear activity, and potentially the role of nuclear activity in galaxy formation. At the peak value of QSO density ($z\approx 2$ to 3), the few QSO host-galaxies detected so far present rest-frame UV fluxes that reach up to 20\%\ of the total QSO luminosity, indicating star-formation rates about 200 \Msun/yr and above for both radio-loud (Lehnert et al. 1992) and radio-quiet samples (Aretxaga, Boyle \& Terlevich 1995, Hutchings 1995). These values are almost an order of magnitude above those of field galaxies at similar redshifts selected through Lyman Break techniques (Steidel et al. 1996, Lowenthal et al. 1997). The properties of these QSO hosts are not unprecedented, since they follow very closely the luminosity--size relation of nearby star forming galaxies, overlapping with its high-luminosity end (Aretxaga, Terlevich \& Boyle 1998). However, the UV fluxes only carry information about the high-mass end of the stellar populations in the galaxies, and say little about the bulk of the stellar mass which is better characterized by optical to NIR observations. Although a few hosts of extreme radio-loud QSOs at $z\approx 2$ have been detected in NIR bands (Lehnert et al. 1992, Carballo et al. 1998), attempts to image the hosts of normal radio-quiet QSOs at the same redshifts have been unsuccessful to date (Lowenthal et al. 1995, Aretxaga et al. 1998). Imaging radio-quiet systems, which constitute more than 95\%\ of all QSOs, is important in order to characterize the bulk of the population. The observed optical sizes of FWHM$\approx 1$ arcsec (Aretxaga et al. 1995), clearly demand a technique which offers the highest available angular resolution. In this paper we focus our attention on the detection of the host of a normal radio-quiet $z\approx 2$ QSO with the Adaptive Optics System in operation at the ESO 3.6m telescope in La Silla. Preliminary results on similar programs to image the host-galaxies of QSOs at $z\approx 0.5 \hbox{ \ and \ } 1.7$ using Adaptive Optics have been presented in a recent conference devoted to quasar hosts (Bremer et al. 1997, Hutchings 1997). \ifoldfss | 98 | 4 | astro-ph9804322_arXiv.txt |
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9804 | astro-ph9804114_arXiv.txt | We obtain self-similar solutions that describe the gravitational collapse of nonrotating, isothermal, magnetic molecular cloud cores. We use simplifying assumptions but explicitly include the induction equation, and the semianalytic solutions we derive are the first to account for the effects of ambipolar diffusion following the formation of a central point mass. Our results demonstrate that, after the protostar first forms, ambipolar diffusion causes the magnetic flux to decouple in a growing region around the center. The decoupled field lines remain approximately stationary and drive a hydromagnetic C-shock that moves outward at a fraction of the speed of sound (typically a few tenths of a kilometer per second), reaching a distance of a few thousand AU at the end of the main accretion phase for a solar-mass star. We also show that, in the absence of field diffusivity, a contracting core will not give rise to a shock if, as is likely to be the case, the inflow speed near the origin is nonzero at the time of point-mass formation. Although the evolution of realistic molecular cloud cores will not be exactly self similar, our results reproduce the main qualitative features found in detailed core-collapse simulations (Ciolek \& K\"{o}nigl 1998). | Low-mass stars are generally believed to form as a result of the gravitational collapse of molecular cloud cores. The cores are initially supported by thermal and magnetic forces, but because of ambipolar diffusion (the drift of ions, to which the magnetic field lines are attached, relative to the dominant neutral gas component), they gradually lose their magnetic support and eventually collapse after becoming ``supercritical'' (see, e.g., Mouschovias 1987 for a review).\footnote{In this paper we reserve the term ``core'' for the high-density central region of a molecular cloud and do {\em not} apply it to the point mass that forms from the collapse of such a core.} The most detailed numerical treatments to date of the problem of the ambipolar diffusion-initiated formation of supercritical cores and the early stages (prior to point mass formation) of their subsequent dynamical collapse have been presented by Mouschovias and collaborators (Fiedler \& Mouschovias~1992, 1993; Ciolek \& Mouschovias~1993, 1994, 1995, hereafter CM93, CM94, CM95; Basu \& Mouschovias~1994, 1995a, 1995b, hereafter BM94, BM95a,b). Because the timescale for core formation is much longer than the timescale for dynamical collapse, special numerical techniques had to be employed in these calculations. The simulations were terminated when the central densities reached $\sim 10^{10}\ {\rm cm}^{-3}$ and the underlying assumptions of isothermality (e.g., Gaustad~1963) and flux freezing onto the ions (e.g., Pneuman \& Mitchell~1965) broke down. These calculations were nevertheless able to demonstrate that {\em supercritical cores begin to collapse dynamically before a point mass (i.e., a protostar) appears at the origin}. The dynamical evolution of supercritical cores after their formation has been studied by many researchers. Solutions exist for the collapse of nonrotating, self-gravitating spheres without thermal support (Henriksen~1994), self-gravitating spheres with thermal support (Penston~1969; Larson~1969; Shu~1977; Hunter~1977; Boss \& Black~1982; Whitworth \& Summers~1985; Foster \& Chevalier~1993) as well as with a combined thermal and isotropic magnetic pressure support (Chiueh \& Chou 1994), and self-gravitating disks with thermal support (Narita, Hayashi, \& Miyama~1984; Matsumoto, Hanawa, \& Nakamura~1997) and also with ordered, frozen-in magnetic fields (Nakamura, Hanawa \& Nakano~1995; Li \& Shu~1997, hereafter LS). In order to choose a particular solution for a given problem, one needs to know the properties of the supercritical core at the time of its formation. This information, however, can only be gleaned from a study of the preceding, quasi-static evolution of the core under the influence of ambipolar diffusion. Although different assumptions about the initial state of the core yield solutions that are qualitatively similar in their gross behavior (the core collapses with near free-fall speeds and a point mass eventualy forms at the center), the solutions do differ in such important details as the accretion rate onto the central point mass and the formation (or absence) of shocks. The well-known examples of the Larson-Penston (1969) and Shu~(1977) similarity solutions in fact represent two extremes of a whole continuum of self-similar collapse solutions specified by a cloud's initial configuration and the conditions at its boundary (Hunter 1977; Whitworth \& Summers 1985; see also Chiueh \& Chou 1994 for a generalization to the case of an isotropic internal magnetic pressure). The Larson-Penston (1969) solution is characterized by a spatially uniform, supersonic (at $\sim 3.3$ times the isothermal speed of sound $C$) infall speed and an inverse-square dependence of the density $\rho$ on the radius $r$ at the instant of point-mass formation (PMF); the mass accretion rate at the center is $\sim 29 \ C^3/G$ (where $G$ is the gravitational constant) at that instant and increases to $\sim 47 \ C^3/G$ immediately after PMF. Numerical simulations of the collapse of nonmagnetic isothermal spheres (Hunter 1977; Foster \& Chevalier 1993) have indicated that this solution provides a good approximation to the conditions near the center at the PMF epoch for clouds that are initially near a marginally stable equilibrium. The Shu (1977) solution strictly applies only to the post-PMF evolutionary phase: it consists of an inner free-fall region and a hydrostatic outer envelope that are separated by an outward-propagating (at a speed $C$ relative to the gas) expansion wave. The envelope corresponds to a singular isothermal sphere ($\rho \propto r^{-2}$) and the mass accretion rate onto the center is $\sim 1 \ C^3/G$. In applying this solution to real systems, it was proposed to identify the initial core configuration at the end of the quasi-static ambipolar-diffusion phase with a singular isothermal sphere (or, more generally, a toroid) at the instant of PMF (e.g., Shu, Adams, \& Lizano 1987; Li \& Shu 1996). However, as we noted above, the conclusion from detailed numerical simulations has been that the dynamical phase of core collapse generally commences well before the PMF epoch, so that the innermost region is not well represented by a quasi-static solution at the time of point-mass formation. Another interesting effect that depends on the specific choice of initial conditions and on the detailed physical properties of the collapsing core is the formation (or absence) of shocks (e.g., Tsai \& Hsu 1995; LS). For example, LS discovered that when, instead of a spherical core, one considers the collapse of a flattened disk, the expansion wave of Shu~(1977) becomes a shock. As we show in this paper, when one takes proper account of the fact that supercritical cores collapse dynamically before a point mass first forms at the origin, that shock disappears. Nevertheless, a physical basis for the formation of shocks in collapsing magnetized molecular cloud cores has been discussed by Li \& McKee (1996), who argued that a hydromagnetic C-shock will appear as a result of the outward diffusion of inwardly advected magnetic flux. The existence of such a shock has been confirmed in the numerical simulations of Ciolek \& K\"{o}nigl (1998, hereafter CK), and it is, in fact, a salient feature of the semianalytic solutions derived in this paper. The aim of the present work is to clarify the effects of ambipolar diffusion in dynamically collapsing supercritical cores. Toward this goal, we construct semianalytic, time-dependent similarity solutions of gravitationally contracting, magnetized, isothermal disks. Although the evolution of real molecular cloud cores is not expected to be exactly self similar, we demonstrate, through a comparison with the detailed numerical simulations of CK, that our solutions capture the main traits exhibited by the latter calculations. Based on an analogous comparison with the results of numerical simulations, Basu (1997) showed that a self-similar scaling describes the pre-PMF evolution in the innermost flux tubes of collapsing supercritical cores quite well. To complement his study, we concentrate in this paper on the post-PMF evolutionary phase. Our approach differs, however, from that of Basu (1997) in that we explicitly solve the induction equation, whereas he accounted for the effects of ambipolar diffusion only in a phenomenological manner.\footnote{Our work is thus also distinguished from that of Safier, McKee, \& Stahler (1997), who studied the effects of ambipolar diffusion in the spherically symmetric, quasi-static limit without explicitly solving the induction equation.} In fact, the solutions that we derive, while involving various simplifications, are nevertheless the first to consistently incorporate ambipolar diffusion into a self-similar representation of the collapse of a magnetized cloud core. \footnote{The effect of {\em weak} magnetic fields on a dynamically collapsing core in the presence of ambipolar diffusion was previously investigated by Galli \& Shu (1993a), who carried out a perturbation expansion of the (nonmagnetic) spherical similarity solution of Shu (1977). As was already noted and discussed by Li \& McKee (1996), the semianalytic solution derived in that paper, as well as the associated numerical calculation in Galli \& Shu (1993b), did not uncover the existence of a flux diffusion-driven shock.} We formulate the problem in \S 2, present our solutions in \S 3, and discuss the results in \S 4. Our conclusions are summarized in \S 5. | In this paper we have presented a self-similar solution of the collapse of a magnetized molecular cloud core (assumed to also be nonrotating and isothermal) that, for the first time, incorporated the effects of ambipolar diffusion in a self-consistent manner. We have focused on the post-PMF (point-mass formation) phase of the collapse of a disk-like core, noting that Basu (1997) had previously explored the self-similar nature of the collapse before a central mass (i.e., a protostar) first appears at the origin. We clarified the distinction between the ideal and nonideal MHD cases by plotting the singular lines in the position--velocity space and showing that they correspond to different critical speeds (the magnetosonic speed and thermal sound speed in the ideal and nonideal problems, respectively). We obtained a solution for the ideal (flux-frozen) case that exhibits a split-monopole field topology near the center. This solution differs from the one obtained by Li \& Shu (1997) in that it involves no shocks. We showed that the shock in the LS solution is a direct consequence of their assumption that the core at the time of PMF is described by a stationary density distribution (corresponding to a singular isothermal toroid), and we pointed out that a shock will generally {\em not} be present under the more realistic assumption of a nonzero inflow speed near the origin at that instant. We demonstrated, however, that a shock is a generic feature of the solution in the nonideal (ambipolar diffusion) case. This (C-type) shock is a direct consequence of the action of ambipolar diffusion in the central region of the core following PMF: the magnetic diffusivity decouples the field from the matter, causing the gas to free-fall to the center (where it accumulates in a point mass) and the field to stay behind and drive a shock outward. We have compared this solution with the results of the numerical simulations of Ciolek \& K\"onigl (1998) and confirmed that, while the more realistic numerical models are not strictly self-similar, our simplified solution nevertheless captures the main features of the core evolution after PMF. | 98 | 4 | astro-ph9804114_arXiv.txt |
9804 | astro-ph9804138_arXiv.txt | s{ We review the topic of Cosmic Microwave Background Anisotropy measurements carried out by means of balloon-borne telescopes. After a short description of the experimental methodology, we outline the peculiar problems of these experiments, and we describe the main results obtained and the perspects for future developments. } | Ballooning for CMB Anisotropy measurements is a very active field worlwide. The activity is growing, for two main scientific reasons: 1) High frequency ($>$90 GHz) and high angular resolution ($\sim$ 10 arcmin FWHM) measurements are possible and effective. 2) These measurements complement the forthcoming data from MAP and are a very important test-bed for Planck technologies. Moreover, good science is being produced, and promises for important results (like the $\ell$-space spectroscopy of the acoustic peaks, or detection/falsification of non-gaussian statistics for the CMB fluctuations) are quite convincing. As an example relevant for this conference, we can mention the fact that determination of several cosmological parameters is possible with very good accuracy from LDB experiments. For example, if all the systematics effects are properly removed, a single LDB experiment with 12 arcmin FWHM beams, 16 total power bolometric detectors with sensitivity of 80 $\mu K \sqrt{s}$, 10 days of observing time spent over a $50^o \times 50^o$ sky region, can measure the power spectrum of the CMB anisotropies with very good accouracy \cite{Silvia}. Fits can be done on the measured power spectrum \cite{KL}, allowing to recover $\Omega_{tot}$ with a 3$\%$ error, $\Omega_{\Lambda}$ with 6$\%$ error, $\Omega_B$ with $1\%$ error, $n$ scalar with 18$\%$ error, $H_o$ with 1$\%$ error, $Q_{rms,PS}$ with 4$\%$ error. Here the errors for any parameter make no assumptions about the value of the other parameters. These measurement errors can be significantly reduced if one or more of the cosmological parameters are constrained by other observations or fixed by assumptions. | 98 | 4 | astro-ph9804138_arXiv.txt |
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9804 | astro-ph9804281_arXiv.txt | We present a description of the observations and data reduction procedures for an extensive spectroscopic and multi-band photometric study of nine high redshift, optically-selected cluster candidates. The primary goal of the survey is to establish new constraints on cluster and galaxy evolution, with specific emphasis on the evolution of galaxy morphology and on the star-formation history of the galaxies within and around distant clusters. We have measured 892 new redshifts for galaxies with $R \le 23.3$. The data will also serve as deep probes of the foreground and background large-scale structures. The observations include broad band optical imaging and spectroscopy with the Low Resolution Imaging Spectrograph at the 10 meter W. M. Keck Observatory telescope; K-band imaging with IRIM at the 4 meter Kitt Peak National Observatory telescope; and deep, high angular resolution imaging with the WFPC2 onboard the Hubble Space Telescope. We also describe the procedures used to obtain morphological information. We have established that six of the nine cluster candidates are indeed real space density enhancements and are representative of those typically associated with clusters of galaxies. The remaining three candidates appear to be projections of several smaller groups at widely separated distances. This success rate is consistent with estimates of the false positive rate in 2D optical high-$z$ cluster searches. | Clusters of galaxies have historically provided an important tool for studying cosmology and the evolution of galaxies. Because of their high concentration of galaxies, clusters provide an environment in which to study large, statistical samples of galaxies. Therefore, examining clusters of galaxies from the local universe to those at high redshift allows us to probe galactic evolution to redshifts of the order of 1. Clusters of galaxies at redshifts of $z \simless 0.2$ have been extremely well cataloged in the optical regime (e.g.\ Abell 1958; Zwicky \etal 1968; Dressler 1980; Shectman 1985; Abell \etal 1989; Lumsden \etal 1992; Dalton \etal 1994). The analyses of local clusters indicate that these systems are dense (Abell richnesses of $\sim 30 - 300$ galaxies), massive ($M \sim 10^{14} - 2 \times 10^{15}~h^{-1}~{\rm M_{\odot}}$), and dominated by early-type galaxies ($\sim 50 - 80\%$ of the total galaxy population). These studies provide a strong basis on which to compare cluster properties at increasingly higher redshift. Detailed photometric, spectroscopic and morphological studies have been extended to clusters of galaxies at redshifts up to $z \sim 0.6$. The first substantial contribution at these redshifts came from Butcher \& Oemler (1984) who found a surprisingly large population of blue galaxies in conjunction with the expected red sequence of early-type cluster members. Further photometric and spectroscopic campaigns, including the ambitious CNOC and MORPHS surveys, have confirmed the progressive bluing of the cluster's galaxy population and have tracked the passive evolution of the early-type galaxies (Dressler \& Gunn 1992; Oke, Gunn \& Hoessel 1996; Yee, Ellingson \& Carlberg 1996; Ellingson \etal 1997; Ellis \etal 1997; Stanford et al.\ 1995, 1997; and references therein). The Hubble Space Telescope (HST) has enabled morphological classification of intermediate-redshift ($z \simless 1$) galaxies on scales which are comparable to the classifications made of their local counterparts. These high-resolution studies have revealed that there may be a substantial change in the morphological composition of the clusters (Smail \etal 1997; Dressler \etal 1997; however, see Stanford \etal 1997). All of these results imply that there is a significant amount of evolution occurring in the cluster environment between redshifts of $z \approx 0.5$ and $z = 0.0$. In order to understand this apparent change in the galaxy population, it is essential to probe in similar detail clusters of galaxies at even higher redshift where the effects of evolution and cosmology will be even greater. In light of this, we have undertaken an extensive survey of nine candidate clusters of galaxies at redshifts greater than 0.6. Only a few optical/near-IR surveys have attempted to detect systematically clusters at high redshift; therefore, we have chosen our cluster sample from the Gunn, Hoessel, \& Oke (1984) survey and the Palomar Distant Cluster Survey (Postman \etal\ 1996). The observational data compiled for this survey, to date, includes deep $BVRIK$ photometry, over 900 low-resolution spectra, and deep F606W/F702W/F814W imagery from HST. The large redshift database allows us to reliably distinguish between physically real clusters and chance line-of-sight projections. The full data allow us to measure the global properties of the clusters, such as profile shape and dynamics, as well as the individual properties of the cluster galaxies, such as color, star-formation rate, and morphology. In this introductory paper to our high-redshift cluster series, we describe in detail the observational and data reduction techniques of each aspect of this survey and present the redshift histograms for our nine fields. The subsequent papers in this series will present the specific analyses and scientific results of this survey. These papers include the second and third installments in this series which present a detailed photometric, spectroscopic, and morphological analyses of the first two clusters to be completed in this survey, CL0023+0423 and CL1604+4304 (Postman, Lubin \& Oke 1998; Lubin \etal 1998). | We have described the data acquisition and reduction procedures of our photometric and spectroscopic campaign to study nine candidate clusters of galaxies at redshifts of $z > 0.6$. The observational program consists of four main parts : \newcounter{discnt} \begin{list} {\arabic{discnt}.} {\usecounter{discnt} \setlength{\leftmargin 0.2in}{\itemsep 0in}{\topsep 0in}{\parskip 0in}} \item Spectra for approximately 80\% of all galaxies down to a Johnson--Cousins $R$ magnitude of $\sim 23.5$ within a fixed area of ${2}^{'}.2 \times {7}^{'}.6$ of each cluster field using the Low--Resolution Imaging Spectrograph (LRIS) at the Keck 10m telescope. We have obtained spectra covering the range 4400 \AA\ to 9500 \AA\ for $\sim 130 - 150$ galaxies per cluster field. Redshifts have been determined for 892 objects. \item Deep $BVRI$ imaging with Keck of all galaxies in the full LRIS field of $6^{'} \times 8^{'}$ centered on each cluster. The $5\sigma$ detection limits are $B = 25.1$, $V = 24.1$, $R = 23.5$, and $I = 21.7$ in our standard 3 arcsecond radius aperture. The photometric data are converted to absolute fluxes in order to obtain absolute spectral energy distributions. \item High angular resolution imagery with HST in order to provide morphological information on the galaxies in the WFPC2 field-of-view (${160}^{''} \times {160}^{''}$) centered on each cluster. Each of the cluster candidates has been or will be observed by HST in Cycles 5 and 6. \item High precision $K$ band photometry across the WFPC2 field--of--view with the KPNO 4m telescope for each cluster. The infrared survey reaches a limiting magnitude of $K^{'} = 20$ in the standard aperture. \end{list} We have presented the redshift histograms for the nine candidate clusters of galaxies in this survey. We find that six of the nine candidate clusters of galaxies are real density enhancements. They include CL0023+0423, CL0943+4804, CL1324+3011, CL1325+3009, CL1604+4304, and CL1604+4321. The remaining three candidates are of a more dubious nature. Their redshift distributions reveal no clear density enhancement but rather an apparent superposition of small groups of galaxies along the line-of-sight. This false positive rate is consistent with the estimate of $\sim 30$\% provided in Postman \etal (1996). At lower redshifts ($z \simless 0.5$) the spurious rate is about 20\% or less. This is based on spectroscopic follow-up of PDCS candidates being done by Holden \& Nichol (1998) at the KPNO 4m telescope. We conclude that optical detection of clusters remains a successful and important method for identifying such systems out to $z \sim 1$ and, further, will provide an important complement to cluster searches at other wavelengths. Results on the star-formation history, dynamics, and morphological properties of CL0023+0423 ($z = 0.84$) and CL1604+4304 ($z = 0.90$) are presented in Papers II and III. \vskip 1.0cm We thank Don Schneider and the anonomous referee for their invaluable comments on this manuscript. The W.M. Keck Observatory is operated as a scientific partnership between the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. It was made possible by the generous financial support of the W. M. Keck Foundation. LML graciously acknowledges support from a Carnegie Fellowship. Support for this work was also provided, in part, by NASA through grant number GO-06000.01-94A from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. \clearpage | 98 | 4 | astro-ph9804281_arXiv.txt |
9804 | astro-ph9804004_arXiv.txt | s{We review estimates of small scale fluctuations due to extragalactic point sources in the Planck Surveyor frequency bands. While our undestanding of the spectral and evolutionary properties of these sources is far from complete, conservative estimates allow us to confidently conclude that, in the frequency range 100--200 GHz, their contaminating effect is well below the expected anisotropy level of the cosmic microwave background (CMB), down to angular scales of at least $\simeq 10'$. Hence, an accurate subtraction of foreground fluctuations is not critical for the determination of the CMB power spectrum up to multipoles $\ell \simeq 1000$. In any case, Planck's wide frequency coverage will allow to carefully control foreground contributions. On the other hand, the all sky surveys at 9 frequencies, spanning the range 30--900 GHz, will be unique in providing complete samples comprising from several hundreds to many thousands of extragalactic sources, selected in an essentially unexplored frequency region. New classes of sources may be revealed in these data. The familiar ``flat''-spectrum radio sources should show spectral features carrying essential information on their physical properties. Crucial information will be provided to understand the nature of radio sources with strongly inverted spectra. Scenarios for the cosmological evolution of galaxies will be extensively tested.} | The multifrequency all-sky maps produced by the Planck Surveyor mission will comprise, in addition to anisotropies which are outgrowths of primordial fluctuations, and whose precision measurements are the main goal of the mission, astrophysical foregrounds, the most important of which, over the frequency range of interest, are those due to emissions in our own Galaxy and to extragalactic radio and mm/sub-mm sources. We deal here with extragalactic sources, which may be a major limiting factor for experiments, like Planck, aimed at accurately determining the cosmic microwave Background (CMB) power spectrum $C_\ell$ up to multipoles $\ell \sim 2000$, corresponding to angular scales $\theta \sim 5'$. In fact, a Poisson distribution of sources produce a white noise power spectrum with the same power in all multipoles~\cite{Tegmark}, so that their contribution to fluctuations in a unit logarithmic multipole interval increases with $\ell$ as $\ell(\ell +1)C_\ell \propto \ell^2$ (for large values of $\ell$), while, at least for the standard inflationary models, which are consistent with the available anisotropy detections, the function $\ell(\ell +1)C_\ell$ yielded by primordial CMB fluctuations is approximately constant for $\ell \lsim 100$, then oscillates and finally decreases quasi exponentially for $\ell \gsim 1000$ ($\theta \lsim 10'$). Hence confusion noise due to discrete sources will dominate at small enough angular scales. In \S$\,$2 we summarize the limitations set by fluctuations due to extragalactic sources on Planck measurements of primordial CMB anisotropies. On the other hand, the multifrequency all sky surveys carried out by the Planck Surveyor mission will provide a very rich database for astrophysical studies; their impact on investigations of physical and evolutionary properties of different classes of extragalactic sources is briefly outlined in \S$\,$3. Our main conclusions are presented in \S$\,$4. | Luckily enough, both for galaxies and active galactic nuclei, the crossover between the radio and the dust emission components, determining a minimum in the spectral energy distribution, is roughly coincident with the CMB intensity peak. The dust temperature tends to be higher for bright distant objects, moving the minimum to higher frequencies in the rest frame and thus partially compensating for the effect of redshift. This situation makes the mm region ideal for mapping primordial anisotropies. Although our understanding of foregrounds at Planck frequencies is far from complete, estimates using worst-case parameters in extrapolating existing measurements to Planck frequencies or angular scales, allow us to safely conclude that, in the frequency range 100-200 GHz, the foreground fluctuations, which are dominated, on small scales ($\theta \lsim 30'$), by extragalactic sources, are well below the expected amplitude of CMB anisotropies over much of the sky. Hence, the removal of foreground contamination is not critical for accurate determinations of the power spectrum of CMB anisotropies up to multipoles of at least $\ell \sim 1000$. On the other hand, while only a small fraction of high Galactic latitude pixels are strongly contaminated by astrophysical foregrounds, the Planck surveys at 9 frequencies will provide sufficiently rich complete samples for astrophysical studies. Spectral information will be provided for ``flat''-spectrum radio sources (compact radio galaxies, radio loud QSOs, BL Lacs, blazars) over a frequency region where spectral features carrying essential information on their physical conditions show up (breaks due to energy losses of relativistic electrons, self-absorption turnovers of flaring components, ...). Planck surveys will be unique in providing complete samples of bright radio sources with inverted spectra, essentially undetectable in radio-frequency surveys. Important classes of sources of this kind are GHz peaked spectrum sources, which may be the youngest stages of radio source evolution and may thus provide insight into the genesis of radio sources, and advection dominated sources, corresponding to final stages of accretion in giant elliptical galaxies hosting a massive black hole. The high frequency Planck channels will detect thousands of dusty galaxies, a large fraction of which at substantial redshifts, allowing to extensively test scenarios for galaxy evolution. The increasing evidence that a large, and perhaps dominant fraction, of star formation at high redshifts may be hidden by dust, makes far-IR to sub-mm surveys an essential complement to optical data. | 98 | 4 | astro-ph9804004_arXiv.txt |
9804 | astro-ph9804142_arXiv.txt | Using the IRAM 30-m telescope, we observed the supernova remnant 3C~391 (G31.9+0.0) and its surroundings in the \cotwo, \hcop, \cstwo, \csthree, and \csfive\ lines. The ambient molecular gas at the distance (9 kpc) of the remnant comprises a giant molecular cloud whose edge is closely parallel to a ridge of bright non-thermal radio continuum, which evidently delineates the blast-wave into the cloud. We found that in a small (0.6 pc) portion of the radio shell, the molecular line profiles consist of a narrow (2 \kms) component, plus a very wide ($> 20$ \kms) component. Both spectral components peak within $20^{\prime\prime}$ of a previously-detected OH 1720 MHz maser. We name this source 3C~391:BML (broad molecular line); it provides a new laboratory, similar to IC 443 but on a larger scale, to study shock interactions with dense molecular gas. The wide spectral component is relatively brighter in the higher-excitation lines. We interpret the wide spectral component as post-shock gas, either smoothly accelerated or partially dissociated and reformed behind the shock. The narrow component is either the pre-shock gas or cold gas reformed behind a fully dissociative shock. Using the 3 observed CS lines, we measured the temperature, CS column density, and H$_2$ volume density in a dense clump in the parent molecular cloud as well as the wide-line and narrow-line portions of the shocked clump. The physical conditions of the narrow-line gas are comparable to the highest-density clumps in the giant molecular cloud, while the wide-line gas is {\it both} warmer and denser. The mass of compressed gas in 3C~391:BML is high enough that its self-gravity is significant, and eventually it could form one or several stars. | Supernovae are thought to be the source of kinetic energy of the interstellar medium, keeping the gas in motion and returning material from dense molecular clouds into the more diffuse interstellar medium and the galactic halo. When a massive star ends its life in a supernova explosion, it often does so in the vicinity of the molecular cloud in which it was born, as is evidenced by the close correspondence of OB associations and giant H~II regions in spiral arms (\cite{elmlad77}). Despite the expected close association between Type II supernovae and molecular clouds, very few cases of supernova-molecular cloud (SN-MC) interaction are known or suspected. The blast wave from a supernova within or near the edge of a cloud will progress rapidly through the inter-clump medium and drive slower shocks into dense clumps. Multiple reflections of high-energy charged particles within the complicated magnetic field of an SN-MC interaction are a possible source of cosmic rays, which will permeate the entire region (\cite{chevalier77}; \cite{esposito96}). The thermal radiation from the remnant interior (mainly X-rays), cosmic rays and their secondary gamma rays, and direct impact of the blast wave onto clumps should visibly perturb the excitation, chemistry, and dynamics of the parent molecular cloud for at least the $\sim 10^5$ year period during which the SN blast wave is most powerful. So far, the only well-known case of an SN-MC interaction is IC 443, where molecular lines have been detected with FWHM $\sim$ 20 km s$^{-1}$ (much wider than the lines from nearby, un-shocked gas), and from energy levels far above the ground state (\cite{white87}; van Dishoeck, Jansen, \& Phillips 1993; \cite{wang92}). Other remnants, including W~28, CTB 109, Kes 79, and W~51C have been suggested as SN-MC interactions based on their proximity to molecular clouds, wide molecular lines, or both (\cite{woot77}; \cite{woot81}; \cite{tatematsu90}; \cite{green92}; \cite{koo97a}; \cite{koo97b}). In the case of W~28, W~44 and 3C~391, 1720 MHz OH emission has been detected from many small spots, with brightness temperatures so high that they must be masers; these masers are thought to be collisionally excited and they strongly suggest the presence of SN-MC interactions (Frail, Goss, \& Slysh 1994; \cite{frail96}). 3C~391 is one of the brightest radio supernova remnants, and high-resolution radio images suggest a `break-out' morphology due to an explosion near the edge of a molecular cloud (\cite{rm93}). The X-ray emission from 3C~391 peaks in its interior and has a thermal spectrum (\cite{rp96}), characteristic of a newly-defined class of supernova remnants, called `mixed-morphology' remnants, whose nature has been linked to interaction with a strongly inhomogeneous pre-shock interstellar medium (\cite{rp98}). A recent map of the \coone\ emission in the vicinity of 3C~391 revealed a giant molecular cloud that is precisely parallel to the bright ridge of radio emission, confirming that its `break-out' radio morphology is indeed due to the strong density contrast between the molecular cloud to the northwest and the relatively empty regions elsewhere (Wilner, Reynolds, \& Moffett 1998). The work described in this paper is part of our recently-initiated campaign to search for and characterize SN-MC interactions in the mixed-morphology supernova remnants. Our first result was the detection of bright [O~I] 63 $\mu$m and dust emission from 3C~391, showing that the blast-wave into the molecular gas is radiative, and the SN-MC interaction is a significant energy loss for the remnant, although it remains globally adiabatic (\cite{reach96}). In this paper, we present new observations of molecular emissions from 3C~391, designed to search for the effects of the SN-MC interaction on the molecular cloud, using millimeter-wave observations at high angular resolution and several transitions requiring a range of physical conditions for excitation. Throughout this paper, we assume a distance to 3C~391 of 9~kpc, which is based on the comparison of H~I 21-cm emission and absorption line profiles (\cite{radakrish}) and the H$_2$CO absorption line at 96 \kms\ (\cite{downes}); our adopted distance is consistent with that adopted by others for this remnant (\cite{rm93}). | We observed the entire supernova remnant 3C~391 in millimeter lines of CS, CO, and HCO$^+$. The lower-excitation lines reveal a giant molecular cloud to the northwest of the remnant, explaining the `break-out' morphology of the radio emission. The interactions between the blast wave and a very dense molecular clump was found within a small ($50^{\prime\prime}$) region that we call 3C~391:BML. A wide component (FWHM 25 \kms) and a narrow component (FWHM 2 \kms) both peak at 3C~391:BML, which is coincident with an OH 1720 MHz maser. The excitation of the wide molecular lines requires both high gas temperature ($> 100$ K) and density ($\sim 3\times 10^5$ cm$^{-3}$). The narrow-line region require somewhat lower density and are consistent with much lower ($\sim 20$ K) temperatures. We identified a clump in the parent molecular cloud with properties similar to the narrow-line region. Therefore, the 3C~391:BML clump was similar to the highest-density clumps in the parent molecular cloud, and it is currently being shocked. The brightness of the wide \cotwo\ line from 3C~391:BML is consistent with C-type molecular shocks with $10^4 < n_0 < 10^5$ cm$^{-3}$ and $10 < v_s < 50$ \kms, or J-type shocks with $n_0\sim 10^3$ and $v_s \sim 100$ \kms. The pressure in the shocked clump is much higher than the estimated ram pressure of the remnant, possibly because of its self-gravity. If so, this clump is a likely site of triggered star formation. A widespread interaction of 3C~391 with molecular gas is evidenced by the distribution of CO and CS lines with central velocities offset from that of the parent cloud by 10 to 15 \kms; these velocities correspond to those of the two OH masers. The interaction comprises nearly an entire hemisphere of the remnant, making 3C~391 a `CO shell' remnant. {\bf Acknowledgment} We thank Hans Ungerechts at IRAM for helping us get started on the IRAM telescope and David Wilner and Steve Reynolds for sharing early results of their observations. WTR thanks the Commissariat d'Energie Atomique, in Saclay, France for hospitality and computing power during part of the data analysis. We thank Bon-Chul Koo for his comments and support. The research described in this paper was carried out in part by the California Institute of Technology, under a contract with the National Aeronautics and Space Administration. \clearpage \begin{deluxetable}{lllllll} \footnotesize \tablecaption{Observed spectral lines and Telescope parameters\label{tab:telparams}} \tablewidth{0pt} \tablehead{ \colhead{transition} & \colhead{frequency} & \colhead{$T_{sys}$\tablenotemark{a}} & \colhead{$\eta_{mb}$} & \colhead{beam} & \colhead{$\delta v$\tablenotemark{b}} & \colhead{$\Delta v$\tablenotemark{c}} \\ & \colhead{(GHz)} & \colhead{(K)} & & \colhead{($^{\prime\prime}$)} & \colhead{(\kms)} & \colhead{(\kms)} } \startdata HCO$^+$($1\rightarrow 0$) & 89.1885 & 310 & 0.82 & 27 & 0.22 & 430\nl CS($2\rightarrow 1$) & 98.9798 & 250 & 0.76 & 24 & 0.24 & 430\nl CS($3\rightarrow 2$) & 146.9690 & 240 & 0.58 & 16 & 0.65 & 290\nl CS($5\rightarrow 4$) & 244.9356 & 490 & 0.43 & 10 & 1.2 & 310\nl CO($2\rightarrow 1$) & 230.5380 & 720 & 0.45 & 10 & 1.3 & 330\nl \enddata \tablenotetext{a}{typical system temperature for observations presented in this paper} \tablenotetext{b}{velocity resolution} \tablenotetext{c}{velocity coverage} \end{deluxetable} \clearpage \begin{deluxetable}{llll} \tablecolumns{4} \footnotesize \tablecaption{Measured properties of spectral lines\tablenotemark{a}\label{tab:spectab}} \tablewidth{0pt} \tablehead{ \colhead{transition} & \colhead{$T_{mb}$} & \colhead{$V_{LSR}$} & \colhead{$\Delta V$ (FWHM)} \\ & \colhead{(K)} & \colhead{(\kms)} & \colhead{(\kms)} } \startdata \cutinhead{wide-line position in shocked clump $(-40^{\prime\prime},-50^{\prime\prime})$\tablenotemark{b}} HCO$^+$($1\rightarrow 0$) & 0.91 & 111.4 & 25.5 \nl CS($2\rightarrow 1$) & 0.32 & 108.9 & 19.0 \nl & 0.24 & 104.2 & 1.4 \nl CS($3\rightarrow 2$) & 0.37 & 108.9 & 20.0 \nl CS($5\rightarrow 4$) & 0.35 & 108.5 & 15.6 \nl CO($2\rightarrow 1$) & 17.6 & 111.1 & 22.5 \nl & 15.1 & 104.2 & 2.6 \nl \cutinhead{narrow-line position in shocked clump $(-10^{\prime\prime},-85^{\prime\prime})$\tablenotemark{b} } HCO$^+$($1\rightarrow 0$) & 1.0 & 105.3 & 1.7 \nl & 0.38 & 102.6 & 14.4 \nl CS($2\rightarrow 1$) & 0.90 & 105.5 & 1.0 \nl & 0.23: & 103.7 & 6.0 \nl CS($3\rightarrow 2$) & 0.55 & 105.4 & 1.2 \nl & 0.23 & 103.7 & 6.0 \nl CS($5\rightarrow 4$) & $<0.13$ & & \nl CO($2\rightarrow 1$) & 11.6 & 105.5 & 1.2 \nl & 9.7 & 104.8 & 3.7 \nl \cutinhead{radio ridge $(-130^{\prime\prime},+90^{\prime\prime})$\tablenotemark{c} } HCO$^+$($1\rightarrow 0$) & 0.22: & 96.4 & 2.4 \nl CS($2\rightarrow 1$) & $<0.03$ & & \nl CS($3\rightarrow 2$) & $<0.04$ & & \nl CS($5\rightarrow 4$) & $<0.13$ & & \nl CO($2\rightarrow 1$) & 8.9 & 96.5 & 4.0 \nl & 3.3 & 102.0 & 2.1 \nl & 3.6 & 107.0 & 3.8 \nl \cutinhead{clump in parent cloud $(-130^{\prime\prime},+240^{\prime\prime})$\tablenotemark{c} } HCO$^+$($1\rightarrow 0$) & 0.32 & 96.1 & 7.3: \nl CS($2\rightarrow 1$) & 0.55 & 96.7 & 3.8 \nl CS($3\rightarrow 2$) & 0.27 & 96.5 & 3.6 \nl CS($5\rightarrow 4$) & $<0.13$ & & \nl CO($2\rightarrow 1$) & 12.0 & 96.0 & 5.4 \nl & 8.4 & 107.9 & 4.0 \nl \enddata \tablenotetext{a}{only components between 90 and 140 \kms\ are listed} \tablenotetext{b}{spectra averaged within $11^{\prime\prime}$ radius} \tablenotetext{c}{spectra averaged within $21^{\prime\prime}$ radius} \end{deluxetable} \clearpage | 98 | 4 | astro-ph9804142_arXiv.txt |
9804 | astro-ph9804156_arXiv.txt | The scattering diameters of \sgra\ and several nearby OH masers ($\approx 1\arcsec$ at 1~GHz) indicate that a region of enhanced scattering is along the line of sight to the Galactic center. The scattering diameter of an extragalactic source seen through this scattering region will be larger by the ratio of the Sun-GC distance to the GC-scattering region separation. This ratio could be a factor of a few, if the scattering region is far from the GC and only a random superposition with it, to more than 100, if the scattering region is within the \hbox{GC}. We have used the VLA to survey 10 (11) fields at 20~cm (6~cm) that are between 7\arcmin\ and 137\arcmin\ from \sgra. Our objective was to identify extragalactic sources and measure their scattering diameters so as to constrain the GC-scattering region separation. In order to find sources within these fields, we have employed pdf\clean, a source detection algorithm in which sources are identified in an image by comparing the intensity histogram of the image to that expected from a noise-only image. We found over 100 sources, with the faintest sources being approximately 3~mJy. The average number of sources per field is approximately 10, though fields close to \sgra\ tend to contain fewer sources. In a companion paper we combine our survey with previous observations of the GC, and we assess the likelihood that the scattering region is so close to the GC that the resulting scattering diameters cause extragalactic sources to be resolved out by our observations. A number of Galactic sources is included in our source catalog. We discuss the double-lobed source 1LC~359.872$+$0.178, potentially an X-ray quiet version of 1E~1740.7$-$2942, a shell-like structure with a central point source, and a possible radio transient. | \label{sec:gc.intro} If viewed through a plasma containing density fluctuations, an otherwise unresolved source will have a visibility, as measured by an interferometer of baseline length~$b$, of \begin{equation} V(b) = \exp\left[-\frac{1}{2}D_\phi(b)\right]. \label{eqn:visibility} \end{equation} The phase structure function, $D_\phi(b) \equiv \langle[\phi(0)-\phi(b)]^2\rangle$, is a measure of the phase perturbations, on a length scale~$b$, imposed on a propagating electromagnetic wave by fluctuations in the electron density. For a plane wave impinging on this scattering region \begin{equation} D_\phi(b) = 8\pi^2r_{\mathrm{e}}^2\lambda^2\int_0^D dz\,\int_0^\infty dq\,q[1 - J_0(bq)]P_{\delne}(q, z), \label{eqn:structurefunction1} \end{equation} where $r_{\mathrm{e}}$ is the classical electron radius, $J_0(x)$ is the zeroth-order Bessel function, $P_{\delne}$ is the spatial spectrum of the density fluctuations, and the integral over $z$ is taken \emph{from the source to the observer}. If the source of radiation is close to or embedded within the scattering medium, so that the medium is illuminated by spherical wavefronts, the argument of the Bessel function is $bq(z/D)$ (\cite{i78}); the factor~$z/D$ accounts for the divergence of spherical waves. The apparent angular diameter of the source is determined by the width of the visibility function, and, hence, by how quickly $D_\phi(b)$ decreases as a function of~$b$. Since $z/D < 1$, the difference in the form of the phase structure function for plane and spherical wavefronts means that sources close to the medium will show smaller angular diameters than those far from it. Hence, by comparing the scattering diameters of Galactic and extragalactic sources along similar lines of sight, one can constrain the \emph{radial} location of the scattering material. Toward the Galactic center, the observed diameter of \sgra\ scales as $\lambda^2$ over the wavelength range 30~cm to~3~mm (Davies, Walsh, \& Booth~1976; \cite{rogersetal94}), as expected if very strong interstellar scattering from microstructure in the electron density determines the observed diameter. Maser spots in OH/IR stars within 25\arcmin\ of \sgra\ also show enhanced angular broadening (\cite{vfcd92}; \cite{fdcv94}). The scattering disks of \sgra\ and many of the OH masers are observed to be anisotropic as well (\cite{vfcd92}; \cite{bzkrml93}; \cite{krichbaumetal93}; \cite{fdcv94}; \cite{y-zcwmr94}); in the case of \sgra, its scattering disk is anisotropic at least over the wavelength range 21~cm to 7~mm. These observations indicate that a region of enhanced scattering with an angular extent of at least 25\arcmin\ (60~pc at 8.5~kpc) is along the line of sight to \sgra. At 1~GHz the level of angular broadening produced by this scattering region is roughly 10 times greater than that predicted by a recent model for the distribution of free electrons in the Galaxy (Taylor \& Cordes~1993, hereinafter \cite{tc93}), even though this model includes a general enhancement of scattering toward the inner Galaxy. Because all of the sources observed through this region have thus far been Galactic sources, with (presumably) approximately the same location (i.e., in the Galactic center), the radial location of the scattering region is unconstrained. The scattering region could be local to the Galactic center, within approximately 100~pc from the Galactic center---which we refer to as the GC model---or the region could be a random superposition and more than 1~kpc from the GC---which we refer to as the RS model. In the GC model, the region would be a site of excess scattering, and presumably arises from processes unique to the GC; in the RS model, the level of scattering in the region would be high, but not unusually so. Previous estimates for the location of the scattering region have ranged from 10~pc to 3~kpc. Ozernoi \& Shisov~(1977) concluded that an ``unrealistic'' level of turbulence is implied unless the region is within 10~pc of the \hbox{GC}. The level of turbulence they considered unrealistic, however, namely $\sqrt{\langle n_{\mathrm{e}}^2\rangle}/\langle n_{\mathrm{e}}\rangle \sim 1$, does appear to occur elsewhere in the interstellar medium (\cite{s91}). Further, van~Langevelde et al.~(1992) showed that the free-free absorption toward \sgra\ would be excessive unless the scattering region was at least 0.85~kpc from the GC, though suitable adjustment of free parameters (outer scale and electron temperature) can decrease the limit to 0.03~kpc. With the free-free absorption they also placed an upper limit on the region's distance from the GC of 3~kpc. Although the GC model is attractive for phenomenological reasons, other sites of enhanced interstellar scattering are found throughout the Galaxy (e.g., NGC~6634, \cite{mrgb90}; Cyg~X-3, \cite{mmrj95}) and the mean free path for encountering such a region is approximately 8~kpc (\cite{cwfsr91}). Identifying the location of the scattering region may provide clues to the origin of the scattering. The density fluctuations responsible for interstellar scattering are believed to be generated by velocity or magnetic field fluctuations (\cite{h84}, 1986; Montgomery, Brown, \& Matthaeus~1987; \cite{s91}; \cite{sg94}; \cite{gs95}). Velocity or magnetic field fluctuations are also a natural means for inducing anisotropy in the density fluctuations and thereby in the scattering disks. If this supposition is correct, the amplitude of the density fluctuations may provide a measure of the coupling between the density and velocity or magnetic field fluctuations or, more generally, provide information about the small-scale velocity or magnetic field in the scattering region. However, because the radial location of the scattering region is unconstrained, relevant quantities, e.g., the rms density, are uncertain by a factor of $\delgc/\dgc$, where $\delgc$ is the GC-scattering region separation and $\dgc$ is the GC-Sun distance. Observations of extragalactic sources viewed through the scattering region could constrain $\delgc$; however, few extragalactic sources have been identified toward the \hbox{GC}. The two sources closest to \sgra\ are B1739$-$298 (\cite{dkvgh83}) and GPSR~0.539$+$0.263 (Bartel~1994, private communication), which are 48\arcmin\ and 40\arcmin\ from \sgra, respectively. Neither of these is within the region of enhanced scattering defined by the OH masers. This paper reports VLA and VLBA observations of potential extragalactic sources seen through the \hbox{GC}. Section~\ref{sec:gc.observe} describes the observations and data reduction, Section~\ref{sec:catalog} discusses the identification of potential extragalactic sources and presents the catalog of sources, Section~\ref{sec:sources} discusses certain Galactic sources found in our VLA survey, and Section~\ref{sec:gc.conclude} discusses our results and presents our conclusions. A companion paper (Lazio \& Cordes~1998, hereinafter \cite{lc98}) combines the results of this paper with the previous observations of OH and \hoh\ masers and free-free emission in a likelihood analysis that constrains the angular extent and radial location of the scattering region. \cite{lc98} also discusses the physical conditions inside the scattering region. | \label{sec:gc.conclude} This paper has reported the results of a program to identify and obtain scattering diameters for extragalactic sources seen through the Galactic center scattering region. Because they are located far behind the GC, the scattering diameters of extragalactic sources, when compared to the scattering diameter of GC sources such as \sgra, can constrain the \emph{radial} location of the scattering region, viz.\ equation~(\ref{eqn:xgalsize}) and Figure~\ref{fig:xgalsize}. Using the VLA we observed 10 (11) fields at~20~cm (6~cm) containing 15 suspected extragalactic sources. We increased our catalog of sources to well over 100 through the use of pdf\clean: The intensity histogram of the primary beam was used to identify positive brightness image pixels that produced deviations from the shape of the expected noise-only histogram. We found approximately 10 sources per field. Follow-up VLBI observations on a subset of these sources have determined the scattering diameters for two heavily scattered extragalactic sources. Their diameters are too small, by factors of 4--10, for them to be seen through the scattering region in front of \sgra. However, they can be used, in conjunction with the heavily scattered masers, to set constraints on the angular extent of the region. Our fields show a paucity of sources near \sgra; a previous survey with more uniform sky coverage, but at a lower sensitivity also shows a paucity. Such a deficit could arise if the scattering toward the GC is so severe that our (and previous) observations resolve out extragalactic sources. The sources reported here are combined with angular broadening measurements of \sgra\ and OH masers and free-free emission and absorption measurements from the literature. These data are then used in a likelihood analysis to determine the model parameters of the GC scattering region (\cite{lc98}). | 98 | 4 | astro-ph9804156_arXiv.txt |
9804 | astro-ph9804010_arXiv.txt | In this paper, the third of a series dedicated to the investigation of the nuclear properties of spiral galaxies, we have {\it (i)} modelled the {\tt WFPC2} F606W nuclear surface brightness profiles of 41 spiral galaxies presented in Carollo et al.\ 1997c, 1998 with the analytical law introduced by Lauer et al.\ 1995, and {\it (ii)} deprojected these surface brightness profiles and their analytical fits, so as to estimate the nuclear stellar densities of bulges of spiral galaxies. We find that the nuclear stellar cusps (quantified by the average logarithmic slope of the surface brightness profiles within 0.1$''$-0.5$''$) are significantly different for $R^{1/4}$-law and exponential bulges. The former have nuclear properties similar to those of early-type galaxies, i.e. similar values of nuclear cusps for comparable luminosities, and increasingly steeper stellar cusps with decreasing luminosity. By contrast, exponential bulges have (underlying the light contribution from photometrically distinct, central compact sources) comparative shallower stellar cusps, and likely lower nuclear densities, than $R^{1/4}$-law bulges. | The galactic nuclei are the repositories of low angular momentum material sunk to the centers over the lifetime of the parent systems. Therefore, they are likely to hold answers to important questions related with the origin of structure in the parent galaxies. In this perspective, establishing the demographics of galactic nuclei along the entire Hubble sequence lies at the heart of our understanding of the complex process of galaxy formation and evolution. Observations of nearby ellipticals and lenticulars with the Faint Object Camera ({\tt FOC}), the Wide Field Planetary Camera ({\tt WF/PC}) and the Wide Field Planetary Camera-2 ({\tt WFPC2}) aboard the {\it Hubble Space Telescope} (HST) have revealed that the nuclei of these galaxies are complex environments (e.g., Crane et al.\ 1993; Jaffe et al.\ 1994; Lauer et al.\ 1995, hereafter L95; Forbes et al.\ 1995; Carollo et al.\ 1997a, hereafter C97a; Carollo et al.\ 1997b, hereafter C97b; Faber et al.\ 1997, hereafter F97). They show surface brightness profiles that increase down to the innermost point measurable at HST resolution, i.e., $I(r) \propto r^{-\gamma}$ as $r \rightarrow 0$ (where $I(r)$ is the surface brightness at the radius $r$, and $\gamma > 0$); furthermore, several galaxies host stellar and gaseous disks, unresolved nuclear spikes, double nuclei. These inner features might possibly be related to the presence of massive black holes (e.g., Lauer et al.\ 1996). By contrast, much is still to be learned about the nuclear properties of nearby spiral galaxies at HST resolution. F97 found that the surface brightness profiles of the three Sa-Sb bulges present in their sample show a behaviour similar to that of early-type spheroidals of comparable luminosity. The same result was found by Phillips et al.\ (1996) for the three spirals of type earlier than Sc contained in their {\tt WF/PC} F555W sample of 20 disk galaxies. Furthermore, Phillips et al.\ found that later type spirals show instead (almost) flat nuclear profiles, and suggested that the nuclear properties of disk galaxies are more closely related to those of nucleated dwarf galaxies than to those of elliptical galaxies. Further exploration is necessary to assess how the nuclear properties scale with the properties of the spheroidal component. This is likely to provide feedback information about the epoch and processes of nucleus, bulge, and, ultimately, galaxy formation. In order to address these issues, we have performed a {\tt WFPC2} snapshot survey in the F606W filter of the nuclei of 107 (mostly Sa to Sc) disk galaxies. In paper I (Carollo et al.\ 1997c) and paper II (Carollo et al.\ 1998) we have presented the 75 targets imaged so far within our program. Our analysis shows that bulge-like structures are present in most of the galaxies. While in some cases these are ``classical'', smooth, featureless $R^{1/4}$-law bulges, in others they are better fitted by an exponential profile (see also Courteau, de Jong \& Broeils 1996, and references therein, for similar results). The exponential bulges include two classes of objects: {\it (i)} dwarf-looking systems, whose surface brightness profiles within $\approx 15''$ are well fitted by a single exponential. These galaxies are strongly bulge-dominated; their surrounding, faint regions (``disk/halo'') show no signs of spiral arms, and have typically a quiescent, i.e. non star forming, appearance. {\it (ii)} Small exponential bulges embedded in dominant (spiral-armed/star-forming) surrounding disks, i.e., the inner exponential structures of double-exponential fits to the surface brightness profiles within $\approx 15''$. The exponential bulges as-a-class are statistically fainter than the featurless, smooth $R^{1/4}$ bulges, for constant disk luminosity and Hubble type. Resolved, central compact sources are found in most of the exponential bulges; the hosts of central compact sources often contain a barred structure. The nature of these compact source, and in particular their relation with e.g., star clusters and Seyfert 2 nuclei, is discussed in Carollo (1998). In this paper, the third of the series, we investigate the relation between the nuclear structure of spiral galaxies and the physical properties of inner disks and/or bulges. In particular, we {\it (i)} present the results of modeling the nuclear surface brightness profiles with the analytical law introduced by L95 (for the 43 galaxies of paper I and II for which we could perform the measurements), {\it (ii)} deconvolve the surface brightness profiles and their analytical fits in order to estimate the nuclear stellar densities, {\it (iii)} study the nuclear properties as a function of the global properties discussed in papers I and II (e.g., $R^{1/4}$ against exponential bulges), and {\it (iv)} compare the nuclear properties of our sample with those observed in early-type galaxies. The paper is organized as follows. In section 2 we briefly summarize the properties of the sample, the data used in our investigation, the procedure adopted for the data reduction, and the steps performed to derive the surface brightness profiles. In section 3 we present the results of the analytical fits applied to the surface brightness profiles, and of deconvolving data and models in spherical symmetry. In section 4 we investigate the dependence of the nuclear properties on global galactic properties. We conclude in section 5. | In this paper we have investigated the relation between the nuclear structure of spiral galaxies and the physical properties of their bulges. In particular, we have {\it (i)} modelled the {\tt WFPC2} F606W nuclear surface brightness profiles of 41 spiral galaxies with the analytical law introduced by L95 (data from papers I and II), and {\it (ii)} deconvolved the surface brightness profiles and their analytical fits in order to estimate the nuclear stellar densities of disk galaxies. Our main result is that $R^{1/4}$-law bulges and exponential bulges have significantly different nuclear stellar cusps and densities. Specifically, $R^{1/4}$-law bulges have steep stellar cusps which steepen with decreasing luminosity; furthermore, their stellar cusp slopes and densities are similar in values to those of early-type systems of comparable luminosity. By contrast, in exponential bulges, the inward extrapolations underlying the light from the compact sources which sit in their very centers imply rather shallow stellar cusps and, very likely, relatively low nuclear stellar densities. \bigskip \bigskip \noindent {\bf Acknowledgements} We heartly thank Tim Heckman and Colin Norman for helpful discussions, and the anonymous referee for constructive comments to an earlier version of this paper. CMC is supported by NASA through the grant HF-1079.01-96a awarded by the Space Telescope Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA under contract NAS 5-26555. This research has been partially funded by grant GO-06359.01-95A awarded by STScI, and has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, Caltech, under contract with NASA. \bigskip \bigskip \noindent \begin{center}{{\bf Appendix A.} Details on the Analytical Fits}\end{center} \medskip We used the same fitting procedure described in C97a; we addess to this reference for further information. This was carried out in two steps: (1) The first step isolated primary from secondary minima of $\chi^2$. The $\chi^2$ values were computed on a grid of points uniformly distributed on a wide hypercube in parameter space (with dimension equal to the number of free parameters). Once a minimum value was found, a new, smaller, hypercube was placed on that location, and the procedure iterated. (2) The minimum value found on the hypercube was then used as starting point to initialize a downhill simplex minimization. We tested the procedure on simulated data, and verified that it recovered the initial values with high accuracy. We accepted as final the fits associated with the absolute minimum of $\chi^2$. Deconvolutions of {\tt WF/PC} data have been proven to be very reliable (e.g., L95). However, in our analysis, we chose not to apply any deconvolution to the post-refurbishment {\tt WFPC2} images, and to correct for PSF-blurring while modeling the light profiles. Therefore, the models were convolved with the appropriate PSFs before being compared to the data. Since pointlike sources with adequate S/N located near to the nuclei were not available for most of the galaxies, we computed the PSFs by running Tinytim (Krist 1992). Focus drifts and breathing modify the PSF profile and affect the flux within a 1 PC pixel radius up to 10\% (and within 5 PC pixels up to 5\%; Suchkov \& Casertano 1997). Therefore, the simulated PSFs obtained by construction at the nominal focus position are in principle of similar quality than PSFs derived from archival stars. Furthermore, our approach of convolving the models rather than deconvolving the data minimizes the effects of using a possibly non-perfect PSF. \bigskip \bigskip \noindent \begin{center}{{\bf Appendix B.} Classifying $R^{1/4}$-law or exponential bulges outside $1''$}\end{center} \medskip In order to ensure that the classification of a bulge as an $R^{1/4}$-law or an exponential structure is valid on a radial range entirely different from that used in the derivation of $\langle \gamma \rangle$ (equal to 0.1$''$-0.5$''$), we performed as a test the $R^{1/4}$-law and exponential fits after excluding the data inside the innermost $1''$. This value is a compromise between a radius large enough to exclude the range where $\langle \gamma \rangle$ is computed, and small enough to still allow the detection of the small, disk-embedded exponential bulges (pentagons in the figures). As an example, the results of the test are illustrated in Figure 8 for the same four galaxies presented in Figure 6. The solid lines represent either a single exponential (left panels) or a double exponential (right panels) fit; the dashed lines represent either a single $R^{1/4}$-law (left panels) or an $R^{1/4}$-law plus exponential (right panels) fit. Two different scales are used for the abscissa for the galaxies in the left and right panels, consistently with the different scales of their bulge components. An offset of two magnitudes has been applied to ESO482G17 for plotting purposes. There are two important points to note: {\it (i)} the kind of profile which provides the bulge classification given in papers I and II, i.e. exponential or $R^{1/4}$-law, still provides a better fit to the inner galactic regions, even when the innermost $1''$ is excluded from the fits; {\it (ii)} the alternative profile with respect to the one that provides the classification given in papers I and II generally provides (not only a worse fit but also) physically meaningless best fit parameters (e.g., for ESO482G17, the $R^{1/4}$-law best fit of Figure 8 has an $R_e\sim385''$). We conclude that the distinction between $R^{1/4}$-law and exponential bulges holds in a radial range which excludes the one used to derive $\langle \gamma \rangle$, and that the difference in nuclear cusp slopes $\langle \gamma \rangle$ between $R^{1/4}$-law and exponential bulges has a physical origin. We retained the bulge parameters presented in papers I and II for our discussion, since those fits gave an overall better description of the profiles. \bigskip \bigskip | 98 | 4 | astro-ph9804010_arXiv.txt |
9804 | astro-ph9804226_arXiv.txt | An object discovered during an infrared survey of the field near the quasar B2 0149$+$33 has an emission line at 2.25\,$\mu$m that we interpret as H$\alpha$ at a redshift of 2.43. The K-band image shows two compact components 10\,kpc apart surrounded by more extended emission over $\sim 20$\,kpc. The H$\alpha$ emission appears to be extended over $\sim 15$\,kpc (2$^{\prime\prime}$) in a coarsely sampled (0\farcs8/pixel) image. The star formation rate may be as high as 250 -- 1000\,M$_\odot$\,yr$^{-1}$, depending on the extinction. Alternatively, the line may be powered by an active nucleus, although the probability of serendipitously discovering an AGN in the survey volume is only $\sim 0.02$. The increasing number of similar objects reported in the literature indicate that they may be an important, unstudied population in the high redshift universe. | Discovering the properties of galaxies at redshifts greater than one requires techniques that can readily distinguish the high redshift objects from those at lower redshift that predominate in any deep image. Multi-wavelength approaches are especially fruitful and have been dominated to date by optical and radio surveys. The use of photometric redshifts based on the strong Lyman break redshifted into the optical band has been one of the most successful of these techniques (Steidel et al. 1996, and references therein). Methods using infrared images are beginning to uncover objects not easily discovered with optical or radio methods that may, nevertheless, constitute a significant fraction of the high redshift population. These include objects distinguished by unusually red colors (Elston, Rieke, \& Rieke 1988; Soifer et al. 1994; Graham et al. 1994; Hu \& Ridgway 1994; Cowie et al. 1994; Dey, Spinrad, \& Dickinson 1995) and objects with emission lines redshifted to infrared wavelengths (Songaila et al. 1994; Thompson, Mannucci, \& Beckwith 1996, hereafter TMB96; Malkan, Teplitz, \& McLean 1995; Bechtold et al. 1997). These objects have been interpreted as elliptical galaxies (Graham et al. 1994; Hu \& Ridgway 1994; Dunlop et al. 1996), young ``protogalaxies'' undergoing bursts of star formation (Eisenhardt \& Dickinson 1992; Graham \& Dey 1996; Malkan, Teplitz, \& McLean 1996; Yee et al. 1996; Bechtold et al. 1997), or active galactic nuclei (Cowie et al. 1994; Dey, Spinrad, \& Dickinson 1995). All of these populations could be significant for cosmology, since the first case implies massive galaxy formation at redshifts greater than $\sim$3, the second indicates a substantial population of young galaxies that can be discovered only in infrared surveys, and the third implies a population of infrared bright active galactic nuclei (AGN) comparable in number density to the populations of AGN discovered by more traditional methods. Only a few such objects have been studied in enough detail to reveal redshifts, source morphologies, and colors. In the course of our survey for emission line galaxies (TMB96), an object was discovered near the quasar B2 0149$+$33 that had an emission line at 2.25\,$\mu$m, the same wavelength as the quasar's H$\alpha$ line at a redshift of 2.43; we call this object TMB 0149 - cK39, or simply cK39. A spectrum between 1.5 and 2.4 $\mu$m, presented and discussed here, confirmed the presence of an emission line nearly coincident in wavelength with that of the quasar. The object is very red, making it unusual among known, distant galaxies. This paper describes the results of these observations and suggests that such objects may be common but previously unobservable owing to the lack of optical emission. | The emission line object, TMB 0149 $-$ cK39, appears to be a pair of galaxies undergoing a merger at redshift of 2.4. If cK39 derives a substantial part of its luminosity from star formation, the formation rate is as high as 1000 M$_\odot$ yr$^{-1}$, an exceptionally large value that is rarely seen in other starforming galaxies. On the other hand, the system could contain at least one active nucleus, perhaps with some contribution from star formation. The individual components appear to be 6 and 7 kpc in extent with the centers separated by 10 kpc, consistent with a merger-induced fuelling of the nuclear activity. An extinction of A$_{\rm V} \sim 1.7^{\rm m}$ is required to produce the observed R-K color of 5.5, requiring the presence of a significant amount of dust in order to suppress the ultraviolet light and redden the galaxies. Although gravitational lensing by a foreground galaxy or cluster of galaxies could enhance the brightness of the emission, there is little evidence of such a galaxy or cluster along the line of sight. The growing number of very red galaxies at high redshift indicates that there are new populations uncovered only in infrared surveys (e.g. Graham \& Dey 1996; Malkan, Teplitz, \& McLean 1996; and references therein). In a separate paper (Thompson et al. in preparation), we present statistics that show extremely red objects (R-K$^\prime > 6$) have a sky density of order 500 deg$^{-2}$ for K$^\prime$ $\leq$19.75, so cK39 may, indeed, be part of a larger population. These results underscore the importance of employing a number of different techniques for exploring the epoch of early star formation and demonstrate that a better understanding of these objects is important to an understanding of galaxy formation in the early universe. | 98 | 4 | astro-ph9804226_arXiv.txt |
9804 | astro-ph9804295_arXiv.txt | s{Direct and indirect detection rates of relic neutralinos are reviewed in the framework of the Minimal Supersymmetric Standard Model. The theoretical estimates are compared with the most recent experimental limits from low--background detectors and neutrino telescopes. The properties of neutralino under the hypothesis that preliminary experimental results of the DAMA/NaI Collaboration may be indicative of a yearly modulation effect are examined.} \normalsize\baselineskip=15pt \vspace{-5pt} | \vspace{-5pt} Supersymmetric theories predict a large number of particles in excess to the Standard Model ones. If the R--parity is conserved, the lightest among all the supersymmetric particles (LSP) must be stable. This feature makes the LSP a dark matter candidate, since this particle can be present today as a relic from the early stages of the evolution of the Universe. Different candidates have been proposed in the framework of supersymmetric theories: the neutralino or the sneutrino\cite{sneutrino} in gravity mediated models, the gravitino\cite{gravitino} or some messenger fields in gauge mediated theories\cite{messenger}, the axino\cite{axino}, stable non--topological solitons (Q--balls)\cite{Qballs} or others. In this paper we will focus on the most promising among all the different candidates, the neutralino, which is defined as the lowest mass linear superposition of photino ($\tilde \gamma$), zino ($\tilde Z$) and the two higgsino fields ($\tilde H_1^{\circ}$, $\tilde H_2^{\circ}$), i.e. $\chi \equiv a_1 \tilde \gamma + a_2 \tilde Z + a_3 \tilde H_1^{\circ} + a_4 \tilde H_2^{\circ}$. The aim of this review is at providing the latest results on the calculation of different kinds of detection rates of relic neutralinos, in the framework of the Minimal Supersymmetric extension of the Standard Model (MSSM), constrained by the most recent experimental data coming from accelerator physics. We do not discuss here the details of the model, for which we refer to Refs.\cite{pinning,extending} and to the references quoted therein. We only recall the standard assumptions employed here: i) all trilinear parameters are set to zero except those of the third family, which are unified to a common value $A$; ii) all squarks and sleptons soft--mass parameters are taken as degenerate: $m_{\tilde l_i} = m_{\tilde q_i} \equiv m_0$; iii) the gaugino masses are assumed to unify at $M_{GUT}$, and this implies $M_1= (5/3) \tan^2 \theta_W M_2$ at the electroweak scale. After these conditions are applied, the free parameters are: $M_2, \mu, \tan\beta, m_A, m_0, A$. The parameters are varied in the following ranges: $10\;\mbox{GeV} \leq M_2 \leq 500\;\mbox{GeV},\; 10\;\mbox{GeV} \leq |\mu| \leq 500\;\mbox{GeV},\; 65\;\mbox{GeV} \leq m_A \leq 500\;\mbox{GeV},\; 100\;\mbox{GeV} \leq m_0 \leq 500\;\mbox{GeV},\; -3 \leq {\rm A} \leq +3,\; 1.01 \leq \tan \beta \leq 50$. In our analysis the supersymmetric parameter space is constrained by all the experimental limits on Higgs, neutralino, chargino and sfermion searches at accelerators. Moreover, the constraints due to the $b \rightarrow s + \gamma$ process\cite{LEP} are satisfied. In addition to the experimental limits, we require that the neutralino is the lightest supersymmetric particle. Finally, the regions of the parameter space where the neutralino relic abundance exceeds the cosmological bound, i.e. $\Omega_{\chi}h^2 > 1$, are also excluded. \vspace{-7pt} | \vspace{-10pt} In this paper we have reported the most recent calculations of the direct and indirect detection rates of relic neutralinos in the framework of the Minimal Supersymmetric Standard Model. We have shown that the theoretical estimates of the detection rates may be at the level of the present experimental sensitivities of low--background detectors and neutrino telescopes. For many supersymmetric configurations, and for median values of the astrophysical parameters which enter in the calculations of the detection rates, the predicted signals may already exceed the present experimental bounds. This shows that the different experimental efforts to search for relic particles are potentially able to deeply investigate the possibility that neutralino is a component of the dark matter of the Universe. An interesting preliminary analysis of the DAMA/NaI Collaboration has shown that their data are compatible, at 90\% C.L., with a modulation signal of the direct detection rate. The features of a neutralino able to satisfy the prerequisites of this signal have been analyzed and it has been shown that many configurations are compatible with a dark matter scenario where the neutralino is the major component, both on galactic and cosmological scales. However, we have to remind here that the occurrence of a possible modulation effect will necessarily require further investigations with much higher statistics. This project is currently under way. \vspace{-7pt} | 98 | 4 | astro-ph9804295_arXiv.txt |
9804 | astro-ph9804341_arXiv.txt | Primordial Black Holes (PBH) may have formed from the collapse of cosmic string loops. The spectral shape of the PBH mass spectrum can be determined by the scaling argument for string networks. Limits on the spectral amplitude derived from extragalactic $\gamma$-ray and galactic $\gamma$-ray and cosmic ray flux observations as well as constraints from the possible formation of stable black holes remnants are reanalyzed. The new constraints are remarkably close to those derived from the normalization of the cosmic string model to the cosmic microwave background anisotropies. | Cosmic strings (CS) are linear topological defects that are believed to originate during phase transitions in the very early Universe \cite{VSrev,HKrev,RBrev}. Here, we consider the ``standard" CS model \cite{Zel,V81}, according to which the network of linear defects quickly reaches a ``scaling" solution characterized by having the statistical properties of the string distribution independent of time if all lengths are scaled to the Hubble radius ($R_H = c t$, where $c$ is the speed of light). Cosmic string loops (CSL) are continually formed by the intersection and self-intersection of long CS (infinite CS or CSL with radius of curvature larger than $R_H$). After formation, a loop oscillates due its own tension and slowly decays by emitting gravitational radiation. The initial length of a CSL is $l(t) = \alpha R_H$, where $\alpha$ is expected to be $\sim G \mu/c^2$. The mass of a CSL is $m(t) = l(t) \mu$, where $\mu$ is the mass per unit of length of the string Since CS also lead to cosmic microwave background (CMB) anisotropies, the string model can be normalized by the recent COBE observations giving the constraint \cite{LP93,ACSSV} \be \label{cmbnorm} G\mu / c^2 \leq 1.7(\pm0.7) \times 10^{-6} \ee Our assumption is that a distribution of PBH was formed by the collapse of a fraction $f$ of CSL \cite{SH,AP}. Hence, from the observational consequences of a present surviving distribution of PBH we can derive updated constraints on the CS scenario \cite{jru}. These constraints are important because: {\it i-)} They may indicate new ways to search for direct signatures from CS; {\it ii-)} They may provide constraints on CS models with symmetry breaking scale $\mu^{1/2}$ smaller than $10^{16}$ GeV which are not constrained by CMB and large-scale structure data; and {\it iii-)} They may provide tighter limits than the CMB on CS models with $G \mu / c^2 \sim 10^{-6}$. Because CS do not dominate the energy density of the Universe, the CS network must lose energy. We derive the rate of CSL production ${{dn_l} \over {dt}}$ from the conservation of string energy in the ``scaling" scenario, \[ \label{consmass} {\dot \rho_{\infty}} - 2 H \rho_{\infty} = - {{dn_l} \over {dt}} \alpha \mu t \, , \] where $\rho_{\infty} = \nu \mu c^{-3} t^{-2}$ is the energy density in long strings and $\nu$ is proportional to the average number of long strings crossing each Hubble volume. Hawking \cite{SH} and Polnarev and Zembowicz \cite{AP} first postulated that a fraction $f$ of the CSL could collapse within its Schwarzchild radius and then form a BH. More recently, Caldwell and Casper \cite{CC} have performed numerical simulations to determine $f$ and found \be \label{fvalue} f = 10^{4.9 \pm 0.2} (G \mu / c^2)^{4.1 \pm 0.1} \, . \ee The BH are sufficiently small that they lose mass due to the Hawking evaporation process. The fraction of the critical density of the Universe in PBH today due to collapsing CSL is (see \cite{jru} and references quoted therein) \be \label{omegaPBH} \Omega_{PBH}(t_o) = \frac{1}{\rho_{crit}(t_o)} \int_{t_*}^{t_o} dt' \frac{dn_{BH}}{dt'} m(t',t_o) \, , \ee where $t_o$ is the present age of the Universe; $t_*$ is formation time for a PBH with mass $M_* = 4.4 \times 10^{14} h^{-0.3} \; \mbox{g}$, which is expiring today; $m(t',t_0)$ is the mass of a PBH formed at a time $t'$ at a later time $t$; and $h$ is the Hubble parameter in units of $100 \mbox{km} \mbox{s}^{-1} \mbox{Mpc}^{-1}$. PBH formed at times $t < t_* \; (M < M_*)$ do not contribute to this integral because they will have evaporated by today. If we assume for simplicity that PBH with mass $M > M_*$ will have evaporated little by the present time, we can approximate $m(t',t_0)$ by $\alpha \mu c t'$. | We have taken advantage of the recent numerical simulations to better understand PBH formation. The observational consequences of a PBH distribution were used to constrain the CS scenario. We have found that the limits on $G \mu/c^{2}$ are comparable to those stemming from other criteria. Unless the mass of the BH remnants is larger than $10^3 m_{pl}$, these remnants will contribute negligibly to the dark matter of the Universe, even if the BH formation rate has the maximal value allowed by the $\gamma$-ray flux constraints. A remnant mass of $10^3 m_{pl}$, however, can arise naturally in some models \cite{CPW} of BH evaporation. In this case, cosmic strings could consistently provide an explanation for the origin of cosmological structure, for the dark matter, and for the origin of the extragalactic $\gamma$-ray and Galactic cosmic ray backgrounds around $100 MeV$. | 98 | 4 | astro-ph9804341_arXiv.txt |
9804 | astro-ph9804207_arXiv.txt | We present new elements in the identification of the lens-system producing the 4 images of the BAL quasar H1413+117, based on the recent HST/NICMOS-2/F160W observations. After a careful PSF subtraction of the 4 images of the quasar, the residual H$_{F160W}$ image reveals the presence of a faint object ($H\sim 20.5$) within the region enclosed by the 4 quasar images. This object corresponds to a single galaxy: the primary lens of the lens-system. We also identify the galaxies around the Cloverleaf which had been proposed to belong to a galaxy cluster/group at high redshift (Kneib et al 1998): the other component in the lens-system that provides the additional ``external" shear. For these galaxies, we have derived a likely redshift based upon their R$_{F702W}$, I$_{F814W}$ and H$_{F160W}$ magnitudes. We find that most of them are consistent with belonging to a galaxy cluster/group with mean redshift $\overline{z}=0.9 \pm 0.1$. Furthermore we detect 2 very red objects ($I-H\sim 4$): the faintest one has no observed optical (R$_{F702W}$ and I$_{F814W}$) counterpart, while the brightest has a predicted redshift around $z\sim 2$, and may be identified with one of the Cloverleaf absorbers. This gravitational-lens system constitutes an excellent target for IR imaging/spectroscopy with the new generation of 8m ground-based telescopes. | The excellent quality and broad wavelength coverage of current observations of gravitational lensing systems have allowed to unveil part of their mysteries. The primary lens is often detected. The immediate surrounding of the multiply imaged quasars/galaxies is studied in great detail and generally shows some galaxy clustering ({\it e.g.} Tonry 1998, Hjorth \& Kneib 1998) or even sometimes X-ray cluster emissions (Hattori et al 1997, Chartas et al 1998). Furthermore, the measure of a time-delay in Q0957+561/PG1115 (Kundic et al 1997, Schechter et al 1997) has strengthened the interest of a detailed study of these gravitational lens systems in order to use them as a cosmological tool. Since the identification by Magain et al (1988) of the quadrupole lens-system called the Cloverleaf, 4 images of the BAL quasar H1413+117 at z=2.558, many efforts have been dedicated to a direct search of the lens or of elements of the lens-system. Early models of the lens-system have involved one or two galaxy-lenses very close to the line-of-sight toward the quasar (Kayser et al, 1990). A more recent analysis showed that an external shear was needed to model this system correctly (Keeton, Kochaneck \& Seljak 1997), and indeed it is probably related to the existence of an overdensity of galaxies nearby, as detected by Kneib et al (1998). The lensing geometry, amplification and time-delays are quite sensitive to the parameters of the galaxy-lens expected to be located amid the 4 images of the quasar. A positive detection of the galaxy-lens would bring stringent constraints in the modeling of the lens-system. However, despite relatively deep searches, either in R and I imaging with the HST (Turnshek et al, 1997; Kneib et al, 1998) or K imaging with the Keck telescope (Lawrence et al 1996), the lensing galaxy has not been detected. The galaxy cluster recently revealed near the Cloverleaf (Kneib et al, 1998) was assumed to be at $z\sim 1.7$ as this corresponds to the mean value of the narrow absorption line-systems observed in the quasar spectra (z=1.44, 1.66, 1.87, 2.07 and 2.09: Turnshek et al, 1988; Magain et al, 1988, Monier et al 1998). Combining the IR image with the R$_{F702W}$ and I$_{F814W}$ WFPC-2 images of this system can allow to estimate the likely redshift for the faint galaxies surrounding the Cloverleaf. The recently acquired HST/NICMOS-2/F160W observation consists of a unique dataset to help answer both the existence of the lensing galaxy and to constrain the distance of the nearby galaxies. The HST/NICMOS-2 data are presented in Sect.2, while in Sect. 3 we discuss the identification of the lensing galaxy after the PSF subtraction of the 4 quasar images. The redshift estimates (derived from photometry) of the surrounding galaxies are explained in Sect.4. The discussion and concluding remarks are provided in Sect.5. Throughout this paper we use $H_0=$50 km/s/Mpc and $\Omega_0=$1. | Two main results have been obtained from the NICMOS-2 data. For the first time the galaxy-lens, H1, close to the line of sight toward the quasar, has been identified. Its position with respect to the quasar line-of-sight is found to be similar (within the uncertainties) to the one derived in the various Cloverleaf gravitational lens models ({\it e.g.} Kneib et al, 1998). It remains difficult to estimate the redshift of the galaxy-lens H1 because the PSF subtraction leaves an increased background noise in the region amid the 4 quasar images. Yet, a redshift estimate around 1.0 or higher is consistent with the H$_{F160W}$ magnitude and the I$_{F814W}$ lower limit magnitude we have derived for H1. Clearly deep spectroscopic data are needed to solve for the determination of its redshift. We find also that there is a unique galaxy-lens, in contradiction to some early models in which two galaxy-lenses had been envisaged (Kayser et al, 1990). Assuming that H1 is around $z\sim 1$, and has similar colors and absolute magnitude than the nearby galaxies gives a Mass-to-light ratio of M($<5.1$kpc)/L$_B$ $\sim$ 25 M/L$_{B\odot}$. With regard to the Cloverleaf environment, we show that 8 nearby galaxies have a most probable redshift around 0.9, giving credit to the presence of a galaxy cluster/group along the line of sight to the Cloverleaf. In our previous modelling (Kneib et al, 1998), we assumed for this galaxy cluster/group a redshift of 1.7, as a mean of the redshifts of the 4 absorbers silhouetted on the quasar spectrum. This value should be revised. The location of the galaxy-lens is now known from the NICMOS-2 observations and will be implemented in a new model of the lens-system. One of the faint galaxies surrounding the Cloverleaf appears to be at a larger redshift, around 2, and might be related with the absorber at z$=$2.07 or 2.09 (Monier et al 1998). Further IR imaging/spectroscopy of these galaxies should remove the remaining uncertainties of the Cloverleaf lens-system. | 98 | 4 | astro-ph9804207_arXiv.txt |
9804 | astro-ph9804031_arXiv.txt | We present data for 18 blazars observed with the X--ray satellite {\sl ASCA}, half of which were also observed contemporaneously with the EGRET instrument onboard {\sl Compton Gamma-ray Observatory} as parts of multi-wavelength campaigns. The observations show a clear difference in the spectra between three subclasses of blazars, namely the High-energy peaked BL Lac objects (HBLs), Low-energy peaked BL Lac objects (LBLs), and quasar-hosted blazars (QHBs). The \asca X--ray spectra of HBLs are the softest, with the power law energy index $\alpha \sim 1 - 2$, and they form the highest observable energy tail of the low energy (LE, synchrotron) component. The X--ray spectra of the QHBs are the hardest ($\alpha \sim 0.6$) and are consistent with the lowest observable energy end of the high energy (HE, Compton) component. For LBLs, the X--ray spectra are intermediate. We find that the radiation process responsible for the HE peak for HBLs {\sl can} be explained solely by Doppler-boosted Synchrotron-Self-Compton (SSC) emission, with the Doppler factor $\delta$ consistent with the VLBI and/or $\gamma$--ray variability data. For many QHBs, on the other hand, the $\gamma$--rays {\sl cannot} be solely due to the SSC mechanism unless $\delta$ is significantly in excess of that inferred from VLBI data. We consider an alternative scenario consistent with the measured values of $\delta$, where the SSC component is still present in QHBs and it dominates in the X--ray band, but it is below the observed $\gamma$--ray spectrum. With an assumption that the peak of the SSC emission is on the extrapolation of the X--ray spectrum, and adopting $\delta$ of 10, we infer the magnetic field $B$ to be 0.1 -- 1 Gauss, and Lorentz factors $\gamma_{b}$ of electrons radiating at the peak of the $\nu F(\nu)$ spectrum of $\sim 10^{3}$ for QHBs; this is much lower than $\gamma_{b} \sim 10^{5}$ for HBLs, even though the values of $B$ are comparable in the two sub-classes. This difference of $\gamma_{b}$ is most likely due to the large photon density expected in QHBs (e.g. from thermal components visible in these objects) as compared with that of HBLs; Compton upscattering of these photons may well provide the observed GeV flux. | The overall electromagnetic spectra of blazars -- a class of active galactic nuclei (AGNs) that includes BL Lac objects and Optically Violently Variable (OVV) quasars -- are believed to be dominated by Doppler-boosted radiation from relativistic jets pointing closely to our line of sight (\cite{blandford78}; \cite{blandford79}; \cite{urry95} for a review of radio loud AGNs). The VLBI studies of these objects show compact radio cores on milli-arcsecond angular scale with jet-like structures which often show superluminal motion, with apparent speeds $v/c \sim 5 - 10$ (e.g., \cite{vermeulen}). Apparent variability time scale and luminosity amplification depend on various powers of the ``beaming'' (Doppler) factor $\delta$ (e.g., \cite{lind}), defined via $\delta = \Gamma_{j}^{-1} (1 - \beta \cos \theta)^{-1}$, where $\Gamma_{j}$ is the bulk Lorentz factor of the emitting matter, $\beta = v/c$, and $\theta$ is the angle of motion with respect to the line of sight. Blazars are commonly detected as $\gamma$--ray sources. The EGRET instrument onboard the {\sl Compton Gamma-Ray Observatory} (\cgro) has so far detected emission in the GeV range from $\sim$ 50 blazars (\cite{fichtel94}; \cite{thompson95}; \cite{mattox97}; \cite{mukherjee97}); $\gamma$--ray emission has been detected up to the TeV range from the nearby BL Lac objects Mkn~421 and Mkn~501 with ground-based Cherenkov telescopes (\cite{punch92}; \cite{petry93}; \cite{quinn96}; \cite{bradbury97}). As these sources show large luminosity and compact emission regions in the spectral range where the opacity to pair production via $\gamma\gamma\rightarrow e^{+}e^{-}$ is large, it is generally accepted that the $\gamma$--ray emission is anisotropic and Doppler-boosted as well (\cite{maraschi92}; \cite{mattox93}; \cite{dondi95}; \cite{buckley96}), suggesting that the {\sl entire} observed electromagnetic emission arises in the jet. As these objects emit in practically every observable waveband, any study of the structure and physical conditions in the jets requires broad-band spectral observations, which, given the rapid large amplitude flux variability, must be conducted simultaneously. The overall spectra of blazars have two pronounced components: one peaking at low energies (LE), $10^{13}-10^{17}$ Hz (e.g., \cite{sambruna96}), and another peaking at high energies (HE), in the $\gamma$--rays (e.g., \cite{montigny95}). For the blazars that are hosted in quasars (QHBs), and for BL Lac objects discovered via radio-selection techniques (the so-called ``Low-energy peaked BL Lacs'' or LBLs), the LE component peaks in the infrared. For the the majority of BL Lac objects -- those found as a result of their X--ray emission -- it peaks in the ultraviolet or even in the soft X--rays (\cite{giommi95}; \cite{sambruna96}; \cite{padovani96}; \cite{fossati97a}), and thus they are named ``High-energy peaked BL Lacs (HBLs)'' (\cite{padovani95}). The local power-law shape, the smooth connection of the entire radio - to - UV (and, for the HBLs, soft X--ray) continuum, as well as the relatively high level of polarization observed from radio to the UV, imply that the emission from the LE component is most likely produced via the synchrotron process of relativistic particles radiating in magnetic field. This view is strongly supported by spectral variability observed in X--rays in a number of HBLs, where the variability at lower X--ray energies lags behind the more energetic X--rays (\cite{tashiro92}; \cite{sembay93}; \cite{kohmura94}; \cite{tashiro95}; \cite{takahashi96}) The HE component, on the other hand, peaks in the $\gamma$--ray band, in the MeV - to - GeV range, and, in the case of a few HBLs, it extends to the TeV range; it is believed to be produced via Comptonization by the same particles that radiate the LE component. The source of the ``seed'' photons, can either be the synchrotron radiation, internal to the jet -- as in the Synchrotron-Self-Compton (SSC) models (\cite{rees67}; \cite{jones74}; \cite{konigl81}; \cite{ghisellini85}; \cite{band85}; \cite{ghisellini89}; \cite{maraschi92}; \cite{bloom96}; \cite{mastichiadis97}). Alternatively, these can be external to the jet, as in the External Radiation Compton (ERC) models: either the UV accretion disk photons (\cite{dermer92}; \cite{dermer97}), or these UV photons reprocessed by the emission line clouds and/or intercloud medium (\cite{sikora94}; \cite{blandford95}), or else, IR radiation ambient to the host galaxy (\cite{sikora94}). The ratio of the power in the HE to the LE components is systematically larger for QHBs than for BL Lac objects (\cite{maraschi94a}; \cite{dondi95}; \cite{sambruna97a}; \cite{ulrich97}; \cite{fossati97b}). If we assume that the LE component is due to the synchrotron radiation, its peak frequency is determined by the intensity of magnetic field and the distribution function of electron energies, while the location of the HE peak is related to the distribution functions of electron and target photon energies. The ratio of the luminosity of these components ($L_{HE}/L_{LE}$), in the context of this synchrotron plus Compton model, is expected to reflect the ratio of energy densities of photon and magnetic fields in the jet. This paper reports the X--ray spectra of 18 blazars measured by \asca in the context of their multi-band emission. The \asca observations and results are described in \S2, followed by multi-band analysis and discussion in \S3. Summary of this paper is presented in \S4. Throughout this paper we use $H_0$=75 km s$^{-1}$ Mpc$^{-1}$, $q_0$=0.5. | As we mentioned previously, the two leading models of the high energy emission in blazars invoke Comptonization, of internal (SSC) or external (ERC) seed photons. In the following analysis, we assume that {\sl both} SSC and ERC processes may operate in blazars. We then estimate the contribution of the SSC emission in the HE component. In order to calculate the predicted luminosity due to the SSC emission, we assume a simple homogeneous model, in which photons are produced in a region of radius $R$ and with a constant magnetic field $B$. We considered the radiation by a single population of relativistic electrons, with a broken power law distribution of Lorentz factors $\gamma_{el}$, and a break point at $\gamma_{b}$ (similar to e.g. \cite{sambruna96}). We also assume that the radiation spectrum of the LE component peaks at a frequency corresponding to that radiated by the electrons with $\gamma_{b}$. The peak frequency of the synchrotron component in the observer frame, $\nu_{sync}$, is then given as, when pitch angle is $\pi$/2: \begin{equation} \nu_{sync} = 1.2 \times 10^{6} \gamma_{b}^{2} B \frac{\delta}{(1+z)} \quad {\rm Hz} \label{eqn1} \end{equation} where B is in Gauss. If the electron energy is still in the Thomson regime, ($\gamma_{el} \times h\nu_{sync} << m_ec^2$), the expected peak of the SSC component in the observer frame ($\nu_{SSC}$) is $\nu_{SSC} = 4 \gamma_{b}^{2} \nu_{sync}/3$. The ratio of the observed luminosity of the SSC component $L_{SSC}$ to the observed synchrotron luminosity $L_{sync}$ is: \begin{equation} \frac{L_{SSC}}{L_{sync}}=\frac{u_{sync}}{u_{B}} \label{eqn2} \end{equation} where the $u_{sync}=L_{sync}/(4\pi R^{2}c\delta^{4})$ is the rest-frame energy density of the synchrotron photons, and $u_{B}=B^{2}/(8\pi)$ is the magnetic field energy density. To check the validity of the assumption that the observed HE component is solely due to the SSC emission, we calculated the beaming factor ($\delta$), which is given from above equations: \begin{equation} \delta^2 = 1.6\times 10^{12} \frac{L_{sync}}{c R^2} \left(\frac{L_{sync}}{L_{SSC}}\right) \frac{\nu_{SSC}^2}{\nu_{sync}^4} \frac{1}{(1+z)^2} \label{eqn3} \end{equation} where $L$ is in erg~s$^{-1}$, $\nu$ in Hz, $c$ in cm s$^{-1}$, and $R$ in cm. We estimate $R$ from the shortest observed variability (doubling) time scale $\Delta t$ observed in any wavelength, as given in Table 2. Assuming that $R \lesssim c \delta \Delta t / (1+z)$, then Eq. 3 can be rewritten as: \begin{equation} \delta^4 \gtrsim 1.6\times 10^{12} \frac{L_{sync}}{c^3 \Delta t^2} \left(\frac{L_{sync}}{L_{SSC}}\right) \frac{\nu_{SSC}^2}{\nu_{sync}^4} \label{eqn4} \end{equation} where $\Delta t$ is in s, and other quantities are as in Eq. 3. An application of this equation to the data in Table 2 assuming $L_{sync}=L_{LE}$, $L_{SSC}=L_{HE}$, $\nu_{sync}=\nu_{LE}$, $\nu_{SSC}=\nu_{HE}$ implies that the lower limits of $\delta$ for HBLs are $\sim $3 or less, which is consistent with the VLBI results (cf. \cite{vermeulen}), and the limits obtained from the arguments of the $\gamma$--ray opacity (cf. \cite{dondi95}). However, for 4 QHBs, where the $\gamma$--ray flux severely dominates the radiative output, we derive values of $\delta$ that are much larger than the VLBI results (see Table 2). This suggests that an additional emission mechanism -- such as the ERC process -- may contribute significantly in the $\gamma$--ray regime, dominating over the SSC flux, and the values of $\nu_{SSC}$ and $L_{SSC}$ are {\sl very} different than $\nu_{HE}$ and $L_{HE}$, with the SSC component ``hidden'' well below the ERC component. However, the fact that the QHBs have X--ray spectra which are hard, with $\alpha \sim 0.6$, and which are {\sl not} located on the extrapolation of the synchrotron optical / UV spectra, implies that the X--rays observed in QHBs are due to a separate emission process than synchrotron. The fact that for most of QHBs the $\gamma$--ray spectra are above the extrapolation of X--ray spectra (Fig. 2) suggests that the dominant process is different for X--rays than it is for $\gamma$--rays. One explanation is that the SSC process dominates in the X--ray range, while the ERC process dominates in $\gamma$--rays (\cite{inoue}). With the assumption that SSC process is dominant in X--rays for QHBs, we estimate the location of the ($\nu_{SSC}$, $L_{SSC}$) point in the log($\nu$) -- log($\nu F (\nu)$) space by the following method. We assume that it lays on or below the extrapolation of the {\sl ASCA} spectrum (line (a) in Fig. 3), but above the highest value of $\nu F (\nu)$ measured by {\sl ASCA}. Since the spectra of QHBs generally have $\alpha < 1$ and thus $\nu F (\nu)$ is the highest at the end of the \asca bandpass (10 keV $\simeq$ 2$\times$10$^{18}$Hz), this second limit is equivalent to $\nu_{SSC} L_{SSC} > 2\times 10^{18}$ Hz $L_{10 \rm keV}$ (line (b) in Fig. 3). We further constrain $L_{SSC}$ using Eq. 3; once we assume a given $\delta$, there is a unique relationship between $L_{SSC}$ and $\nu_{SSC}$ described as: \begin{equation} L_{SSC} = 1.6 \times 10^{12} \left(\frac{L_{sync}^{2}}{cR^{2}\nu_{sync}^{4} \delta^{2} (1+z)^2} \right) \nu_{SSC}^{2} \quad {\rm erg~s^{-1}} \label{eqn5} \end{equation} where $L$, $R$, $c$, and $\nu$ are in the same units as in Eqs 3 \& 4. The VLBI data and $\gamma$-ray opacity argument suggest that $5 < \delta < 20$ for most blazars (e.g.,\cite{vermeulen}; \cite{dondi95}). This corresponds to the lines (c) and (d) in Fig. 3, respectively for $\delta$=5 and 20. The above four constraints correspond to the shaded area of Fig. 3, where for illustration, we use the overall spectral energy distribution for the QHB CTA~102. Since we have to use a unique value in calculating the physical parameters, we use $\delta$ = 10 as a geometrical mean between 5 and 20. The $L_{SSC}$ - $\nu_{SSC}$ line calculated from Eq. 5 corresponding to $\delta$ = 10 intersects both the extrapolation of the {\sl ASCA} spectrum and the highest {\sl ASCA} value, and the intersections yield the lower and upper values for both $L_{SSC}$ and $\nu_{SSC}$. We adopt a mean of these values, which are given in Table 2, and plotted in Figure 4c. We used \ginga data instead of \asca data for 3C~279 because a simultaneous campaign from radio to $\gamma$--ray bands was conducted during \ginga observation. The values for the other blazars where $\delta$ derived from Eq. \ref{eqn4} is $<20$, are calculated by assuming $L_{SSC}$ = $L_{HE}$. For LBL AO0235+164 where $\delta$ derived from Eq. 4 is $>20$, we assume $L_{SSC}$ = $L_{LE}$ because the \asca spectra of AO0235+164 is thought to be mixture of the LE and HE component, as discussed by \cite{madejski96} based on the \rosat and \asca spectra, so that the above method may be inappropriate. For two QHBs (3C~273, PKS~0208-512) the lower limits of $\delta$ are $\sim$5. The fact that observed $\gamma$-ray flux of PKS~0208-512 is much higher than the extrapolation of \asca spectrum implies the $\gamma$-ray peak is not solely due to SSC mechanism. Therefore we applied the above method to this source. On the other hand, since the $\gamma$-ray spectrum of 3C~273 is below the extrapolation of \asca spectrum, the $\gamma$-ray emission is assumed to be due to SSC mechanism so there, we assume $L_{SSC}$ = $L_{HE}$. It is important to note, however, that 3C~273 is unique as compared to other blazars considered here in that the ``blue bump'' is very pronounced, and thus it is unlikely that the jet dominates the entire electromagnetic emission, and therefore, a more complex analysis is necessary (see, e.g., \cite{montigny97} for further discussion). Once we obtain $L_{SSC}$, and $\nu_{SSC}$, we can calculate the strength of the magnetic field and the electron Lorentz factor $\gamma_{b}$ from Eq. (\ref{eqn1}), (\ref{eqn2}) and those are given as follows: \begin{equation} B = 0.27 \left(\frac{R_{\mbox{pc}}}{10^{-2}}\right)^{-1} \left(\frac{\delta}{10}\right)^{-2} \sqrt{\left(\frac{L_{sync}}{10^{46}}\right) \left(\frac{L_{sync}}{L_{SSC}}\right)} \quad {\rm Gauss} \end{equation} \begin{equation} \gamma_{b} = 1.8\times 10^3 \left(\frac{R_{\mbox{pc}}}{10^{-2}}\right)^{1/2} \left(\frac{\delta}{10}\right)^{1/2} \left(\frac{\nu_{sync}(1+z)}{10^{13}}\right)^{1/2} \left[\left(\frac{L_{sync}}{10^{46}}\right) \left(\frac{L_{sync}}{L_{ssc}}\right)\right]^{-1/4} \end{equation} where $R_{\mbox{pc}}$ is size of emission region in parsecs, and other quantities are as in Eqs 3, 4, \& 5. As before, the upper limit of the size $R$ can be estimated from the observed time variability ($\Delta t$) from Table 2, given by $R \lesssim c\Delta t \delta / (1+z)$. Our calculated values of $B$ and $\gamma_{b}$ are plotted respectively in Figures 4d and 4e. In these Figures, we also plot the values calculated with $R = 0.01$ pc, which would correspond to an observed variability time scale of $\sim 1$ day. From our analysis, the magnetic field for blazars observed with \asca is inferred to be 0.1 -- 1 Gauss. The value of $B$ is comparable between the different subclasses of blazars, although $B$ is somewhat lower in HBLs than in QHBs. With these values of $B$, we estimate $\gamma_{b}$ to be $10^{3} - 10^{4}$ for QHBs, and $10^{5}$ for HBLs. The differences of $\gamma_{b}$ between different sub-classes of blazars imply that the relativistic electrons are accelerated to higher energies in HBLs than in QHBs. Alternatively, higher $\gamma_{b}$ in HBLs might be obtained by increasing $\delta$. However, in those objects, we believe that there is no contribution to the $\gamma$--ray production from other mechanisms besides SSC, and thus the observed $L_{HE}$ is $L_{SSC}$. In such case, $\gamma_{b}$ depends on $\delta$ only linearily (cf. Eq. 7 and $R\lesssim c\Delta t \delta/(1+z)$), and thus varying $\delta$ to be 5 or 20 respectively decreases or increases our derived $\gamma_{b}$ only by a factor of two, which is small when compared to the large difference of $\gamma_{b}$ calculated by us (cf. Fig. 4e). In QHBs the strong optical and UV line emission implies a presence of dense external radiation fields. This means that in the frame of reference of the jet, these can easily dominate over the internal synchrotron radiation, resulting in the ERC emission dominating over the SSC emission in $\gamma$--ray band (e.g., \cite{sikora97}). It is likely that the difference of $\gamma_{b}$ is most likely due to the large photon density in QHBs as compared with that of HBLs. It should be noted that TeV photons have been observed only from HBLs, where we calculate higher values of $\gamma_{b}$. | 98 | 4 | astro-ph9804031_arXiv.txt |
9804 | astro-ph9804096_arXiv.txt | In this paper we present a new method that can be used for analysis of time of arrival of a pulsar pulses (TOAs). It is designated especially to detect quasi-periodic variations of TOAs. We apply our method to timing observations of PSR B1257+12 and demonstrate that using it it is possible to detect not only first harmonics of a periodic variations, but also the presence of a resonance effect. The resonance effect detected, independently of its physical origin, can appear only when there is a non-linear interaction between two periodic modes. The explanation of TOAs variations as an effect of the existence of planets is, till now, the only known and well justified. In this context, the existence of the resonance frequency in TOAs is the most significant signature of the gravitational interaction of planets. | The first extra-solar planetary system was discovered by \cite{Wolszczan:92::} around a millisecond radio pulsar, PSR B1257+12. The three planets orbiting the pulsar have been indirectly deduced from the analysis of quasi-periodic changes in the times of arrival (TOAs) of pulses caused by the pulsar's reflex motion around the center of mass of the system. In the analyses of this kind, it is particularly important to establish a reliable method of distinguishing planetary signatures from possible TOA variations of physically different origin. In the case of PSR B1257+12, it was possible to make this distinction and confirm the pulsar planets through the detection of mutual gravitational perturbations between planets B and C \cite[]{Wolszczan:94::}, following predictions of the existence of this effect by \cite{Rasio:92::}, \cite{Malhotra:92::} (see also \cite{Malhotra:93::}, \cite{Rasio:93::} and \cite{Peale:93::}). Practical methods of detection of the TOA variations caused by orbiting planets include direct fits of Keplerian orbits to the TOA or TOA residual data \cite[e.g.]{Thorsett:92::,Lazio:95::} and model--independent frequency domain approaches based on Fourier transform techniques \cite[]{Konacki:96::,Bell:97::}. In fact, it appears that it is best to search for periodicities in TOAs (or post-fit TOA residuals left over from fits of the standard timing models) by examining periodograms of the data, and then refine the search by fitting orbits in the time domain using initial orbital parameters derived from a frequency domain analysis. The presence of planets around a pulsar causes pulse TOA variations which, for planets moving in orbits with small eccentricities, have a quasi-periodic character and generate predictable, orbital element-dependent features in the spectra of TOA residuals. This has led \cite[]{Konacki:96::} to devising a method of TOA residual analysis based on the idea of a successive elimination of periodic terms applied by Laskar (1992) in his frequency analysis of chaos in dynamical systems. The frequency analysis provides an efficient way to decompose a signal representing the TOA residual variations into its harmonic components and study them in an entirely model-independent manner. As shown in \cite[]{Konacki:96::} this method works perfectly well under idealized conditions in which covariances among different parameters of the timing model are negligible. In this paper, we present an improved scheme for the frequency analysis of pulsar timing observations in which a successive elimination of periodicities in TOAs is incorporated in the modelling process rather than being applied to the post-fit residuals. This makes the results obtained with our method less sensitive to the effect of significant covariances which may exist between various timing model parameters. We apply the frequency analysis to TOA measurements of the planet pulsar, PSR B1257+12, using a computer code developed to fit spectral timing models to data. We show that our method allows an easy detection of the fundamental orbital frequencies of planets A, B and C in the pulsar system and the first harmonics of the frequencies of planets B and C generated by eccentricities of the planetary orbits. Furthermore, by detecting the effect of perturbations between planets B and C, we demonstrate that the frequency analysis method represents a sensitive, model-independent tool to analyze nonlinear interactions between periodic modes of processes of various physical origins. | 98 | 4 | astro-ph9804096_arXiv.txt |
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9804 | astro-ph9804269_arXiv.txt | We have developed in detail the theory of X-ray line and continuum production due to atomic interactions of accelerated ions, incorporating in our calculations information from a broad range of laboratory measurements. We applied our calculations to the Orion region from which nuclear gamma-ray lines were observed with the COMPTEL instrument on {\it CGRO}. The accelerated particles which produce this gamma-ray emission via nuclear reactions also produce X-ray lines via atomic interactions. We predict strong line emission in the range from 0.5 to 1 keV, mainly due to de-excitations in fast O ions. While much of the diffuse X-ray emission observed with ROSAT from Orion could be due to accelerated ions, the current X-ray data do not provide unambiguous signatures for such an origin. If future observations with high spectral resolution would confirm the predicted X-rays, the combined analysis of the X-ray and gamma-ray data will set important constraints on the origin of the accelerated particles and their interaction model. | Strong gamma-ray emission in the 3-7 MeV range has been detected from the Orion molecular cloud complex with the COMPTEL instrument on the {\it Compton Gamma Ray Observatory (CGRO}; Bloemen et al. 1994, 1997). As the observed spectrum exhibits characteristic structures (Bloemen et al. 1997), this emission is most likely due to the superposition of nuclear gamma-ray lines, mainly the 4.44 MeV line from $^{12}$C and the 6.13, 6.92 and 7.12 MeV lines from $^{16}$O. Such line emission can only be produced by accelerated particle interactions. Gamma-ray emission at photon energies $>$30 MeV was also observed from Orion, with the EGRET instrument on {\it CGRO} (Digel, Hunter, \& Mukherjee 1995). This gamma-ray emission is consistent with pion production and bremsstrahlung due to irradiation by standard Galactic cosmic rays (Digel et al. 1995). As such cosmic rays underproduce the observed line emission by at least three orders of magnitude, the gamma-ray line production in Orion must predominantly be a low energy cosmic ray phenomenon. Information on the spatial distribution of the gamma-ray line emission in Orion has come from both the COMPTEL and {\it CGRO}/OSSE observations. The extended nature of the emission seen in the COMPTEL map of Orion (Bloemen et al. 1997) could provide an explanation for the fact that so far it was not possible to confirm the COMPTEL results with OSSE (Murphy et al. 1996; Harris et al. 1998). Based on the observed line widths, Bloemen et al. (1994) first suggested that the line emission is produced by accelerated C and O ions interacting with ambient H and He, rather than by accelerated protons and $\alpha$-particles interacting with ambient C and O. More detailed analyses of the initial COMPTEL data have shown that a mix of the two processes could not be ruled out (Ramaty, Kozlovsky, \& Lingenfelter 1995; Cowsik \& Friedlander 1995). But, as the emission peaks in the more recent COMPTEL data do not appear at the line center energies for $^{12}$C and $^{16}$O de-excitations (Bloemen et al. 1997), a significant narrow-line contribution from accelerated proton and $\alpha$-particle interactions seems to be excluded (Kozlovsky, Ramaty, \& Lingenfelter 1997). This conclusion is also supported by energetic arguments, as the very large power deposited by the accelerated particles into the ambient medium in Orion is lowered by enhancing the C-to-proton and O-to proton abundance ratios (Ramaty et al. 1995; Ramaty, Kozlovsky \& Lingenfelter 1996). Apart from the observed emission in the 3-7 MeV band, the COMPTEL observations revealed only upper limits at other gamma-ray energies (Bloemen et al. 1994, 1997). In particular, the upper limit on the 1-3 MeV emission sets constraints on the accelerated Ne-Fe abundances relative to those of C and O. The suppression of both the Ne-Fe and proton and $\alpha$-particle abundances relative to C and O could be understood if the seed particles injected into an as-yet unknown particle accelerator (see Nath \& Biermann 1994; Bykov \& Bloemen 1994) come from the winds of massive stars or the ejecta of supernovae resulting from massive progenitors (Bykov \& Bloemen 1994; Ramaty et al. 1995; Cass\'e, Lehoucq, \& Vangioni-Flam 1995; Ramaty et al. 1996; Parizot, Cass\'e, \& Vangioni-Flam 1997a). Ip (1995) and Ramaty et al. (1996) have also considered the possible acceleration of ions resulting from the breakup of interstellar dust. The gamma-ray line production in Orion should be accompanied by a large ionization rate of the ambient medium which could exceed the observed infrared luminosity (Cowsik \& Friedlander 1995). This problem is alleviated if the gamma-rays are produced at cloud boundaries, but not in their interiors. The accelerated particles could have ionized $\sim$2$\times$10$^4$M$_\odot$ in 10$^5$ years (Ramaty 1996), a small fraction of the total available mass. It is thus possible that a large fraction of the power that accompanies the gamma-ray production is deposited in an ionized gas. While the X-ray emission produced by low energy particle interactions is potentially a promising tracer of low energy cosmic rays in the Galaxy (e.g. Hayakawa \& Matsuoka 1964), there are as-yet no astrophysical X-ray observations that unambiguously indicate the presence of such cosmic rays. The Orion region, however, has become an interesting target owing to the COMPTEL discovery of the nuclear gamma-ray line emission. A variety of processes lead to X-ray production by low energy ion interactions. Inverse bremsstrahlung (Boldt \& Serlemitsos 1969) results from the interactions of fast ions and ambient electrons; secondary electron bremsstrahlung is produced by knock-on electrons accelerated in fast ion interactions (Hayakawa \& Matsuoka 1964). Both of these processes lead to continuum X-ray emission. X-ray line emission results from atomic de-excitations in the fast ions following electron capture (Silk \& Steigman 1969; Watson 1976; Pravdo \& Boldt 1975; Bussard, Ramaty, \& Omidvar 1978) and in ambient ions following inner-shell vacancy creation. The latter process has not yet been applied to astrophysics. Dogiel et al. (1997) have recently considered the X-ray emission that should accompany the gamma-ray line production in Orion. They have only considered the secondary electron bremsstrahlung and concluded that the 0.5-2 keV emission that accompanies the observed gamma-ray line emission from Orion will exceed the upper limits that they derived using ROSAT observations. We have subsequently taken into account both continuum processes and line emission from de-excitations in fast O (Ramaty, Kozlovsky, \& Tatischeff 1997a) and showed that, even though the inverse bremsstrahlung is more important than the secondary electron bremsstrahlung, the total X-ray continuum emission from Orion is not inconsistent with the Dogiel et al. (1997) derived ROSAT upper limit. On the other hand, we showed that a conflict may exist between that ROSAT upper limit and the X-ray line emission following electron capture onto fast O nuclei. However, as we suggested, this conflict could be resolved if the X-ray and gamma-ray lines are produced in an ionized medium or if the current epoch accelerated particle spectrum is suppressed at low energies, for example by energy losses. In this paper we present detailed calculations of X-ray continuum and line production by accelerated particle interactions. The bulk of our treatment is for a steady state, thick target model with a neutral ambient medium. This is the standard model in which most of the gamma-ray calculations have been carried out (e.g. Ramaty et al. 1996). But we have also investigated the effects of an ionized ambient medium and a time-dependent model, as these modifications could have important consequences on the predicted X-ray to gamma-ray production ratio. In our treatment of the continuum, we have supplied the details of the calculations and we have improved the employed cross sections, thereby confirming our previous preliminary results (Ramaty et al. 1997a). We have greatly expanded our treatment of X-ray line emission. We have investigated in detail the atomic physics relevant to line emission from de-excitations in fast O, checking our theoretical calculations against laboratory data whenever available. We then expanded the treatment to the other abundant accelerated ions (C, N, Ne, Mg, Si, S and Fe), and we have also calculated the X-ray line emission produced in ambient ions following inner-shell vacancy creation by the accelerated particles. We have used the ROSAT all-sky survey (Snowden et al. 1995) to derive the X-ray count rates from the Orion region that could be associated with accelerated particle interactions; the implied fluxes are quite different from the upper limit given by Dogiel et al. (1997). The unambiguous future detection of the predicted X-rays produced by accelerated particles in Orion, and potentially elsewhere in the Galaxy, should provide important new insights into the origin of the low energy cosmic rays whose presence in Orion is revealed by the COMPTEL gamma-ray line observations. | We have investigated all the processes that lead to X-ray production by low energy cosmic rays for a variety of accelerated particle compositions and energy spectra. We demonstrated that the dominant continuum producing process is inverse bremsstrahlung produced by fast ions interacting with ambient electrons. In addition, there is also a significant contribution from the bremsstrahlung produced by secondary knock-on electrons. However, below a few keV the total X-ray emission produced by accelerated ions is dominated by relatively broad line emission (line widths $\delta E/E$$\simeq$0.1) resulting from de-excitations in the fast ions following electron captures and excitations. In addition, accelerated particle interactions also produce much narrower X-ray lines, due to inner-shell vacancy creation. The most prominent of such line is that at 6.4 keV from ambient Fe. We have calculated the X-ray line and continuum emission produced by the accelerated particles in Orion which are thought to be responsible for the nuclear gamma-ray line emission observed with COMPTEL (Bloemen et al. 1994, 1997). By first comparing the results with the extragalactic diffuse X-ray background, we found that while the continuum is generally below this background, the line emission from about 0.5 to 1.5 keV exceeds the background for all the combination of parameters that we considered. We wish to point out that there could be a significant contribution to the $\sim$0.5-1.5 keV diffuse X-ray background from as-yet unknown sources within our Galaxy (e.g. Park et al. 1997), leaving the possibility that a substantial fraction of the observed X-ray intensity in this energy range results from low energy cosmic ray interactions. We have also compared our results with ROSAT observations of Orion in the 0.47 to 1.2 keV energy band, again normalizing the X-ray emission to the observed gamma-ray emission. We found that there is no conflict between the predicted total X-ray emission (lines and continuum) and the data for a broad range of parameters if the gamma-ray line emission is uniformly distributed over the entire molecular cloud complex. This conclusion differs from our previous one (Ramaty et al. 1997a) because of a lower predicted X-ray line emission, resulting from improved atomic physics input, and because our estimated ROSAT flux from Orion is higher than the upper limit given by Dogiel et al. (1997). However, the COMPTEL data show significant spatial structure. We found that for the most prominent hot spot in the COMPTEL map, the standard thick target, steady state interaction model, with a neutral ambient medium, predicts X-ray fluxes which exceed the ROSAT data for a broad range of parameters. But the calculations could be consistent with the data for any one, or a combination of the following possibilities: a very hard accelerated particle spectrum; a partially ionized ambient medium; and a time-dependent accelerated particle energy spectrum resulting from essentially instantaneous acceleration some tens of thousand of years ago. There are as-yet no astrophysical X-ray observations that would unambiguously indicate an origin resulting from low energy, accelerated ion interactions. Our calculations show that the most promising signatures are the relatively broad lines between 0.5 and 1.5 keV, mainly the lines from fast O, and that a promising target is the Orion region where the presence of such accelerated particles is known from gamma-ray line observations. We acknowledge K. Omidvar for discussions on the atomic processes leading to X-ray line production. V. T. acknowledges an NRC-NASA/GSFC Research Associateship. \clearpage | 98 | 4 | astro-ph9804269_arXiv.txt |
9804 | astro-ph9804025_arXiv.txt | We report the serendipitous discovery of a 7-s X-ray pulsar using data acquired with the {\it Advanced Satellite for Cosmology and Astrophysics} (\asca ). The pulsar is detected as an unresolved source located towards a region of the Galactic plane ($l,b \simeq 29.5, 0.08$) that coincides with an overdensity of star-formation tracers. The signal suffers tremendous foreground absorption, equivalent to $N_H \simeq 10^{23}$ cm$^{-2}$; the absorption correlates well with a line-of-sight that is tangential to the inner spiral arms and the 4-kpc molecular ring. The pulsar is not associated with any known supernova remnants or other cataloged objects in that direction. The near sinusoidal pulse (period $P \simeq 6.9712$) is modulated at 35\% pulsed amplitude, and the steep spectrum is characteristic of hot black-body emission with temperature $kT \sim 0.65$ keV. We characterize the source as an anomalous X-ray pulsar (AXP). | A canonical young pulsar (period $\sim 100$ ms and stellar dipole field $\sim 3\times10^{12}$ G) is a rapidly rotating neutron star, created as the stellar remnant during a Type II (or Ib) supernova explosion of a massive star. The birthrate of pulsars is known to be close to $1 - 3$ per century, and it is estimated that there are about $10^5$ active and $10^8$ defunct neutron stars in the Galaxy (see Lorimer et al. 1993 and refs. therein). In the last few years, there has been growing recognition of a population of ultra-magnetized neutron stars, or ``magnetars'' (Thompson \& Duncan 1993). The mostly circumstantial evidence comes from investigations of the following categories of objects: the soft gamma-ray repeaters (Thompson \& Duncan 1995; Frail et al. 1997), long period pulsars in supernova remnants (Vasisht \& Gotthelf 1997 and refs. therein), other seemingly isolated, young, long period pulsars (Thompson \& Duncan 1996) nowadays referred to as the anomalous X-ray pulsars (AXP; van Paradijs et al. 1995), and perhaps their older variants (Kulkarni \& van Kerkwijk 1998). These objects share some common properties; they are steady, bright X-ray sources ($L_X \simgt 10^{35}$ erg s$^{-1}$) which show no signs for an accompanying companion, those with known periods are found to be spinning down, and all are relatively young ($ \simlt 10^{5}$ yr-old). The evolutionary consequences of such large dipole fields are reflected in the properties listed above. Most importantly, large braking torques acting on the star cause it to spin-down rapidly, and the magnetic free energy quickly dominates over the rotation energy, i.e., within several hundred years. For the above sources, the rotation rates lie between 6 - 12 s, with ages $\simlt 10^5$ yr (for the SGRs the evidence for periods is indirect, however, their ages are well constrained due to their association with supernova remnants). It is believed that field decay, which is expected for ultramagnetized neutron stars (Thompson \& Duncan 1996; also Goldreich \& Reisenegger 1992), influences the thermal evolution and powers the large X-ray luminosities observed for the purported magnetars, $L_X \simgt 10^{35}$ erg s$^{-1}$. If magnetars represent the tail-end of the magnetic field distribution of neutron stars, then they are bound to be rare. Assume that their birthrate is 10\% the birthrate of neutron stars (some justification for this comes from the estimated birthrates of SGRs; Kulkarni et al. 1994), and that they have active X-ray lifetimes of $\sim 10^{5}$ yr. These assumptions imply that at present there are only $\sim 100$ active magnetars in the Galaxy, a conclusion that is borne out by the observations of the aforementioned objects. The fact that we observe the five known AXPs through large column densities in the Galaxy, $N_H \simgt 10^{22}$ cm$^{-2}$, suggests that they are indeed that rare, and the fact that they are often associated with supernova remnants or lie near star-formation regions (in spite of the large random velocities usually attributed to neutron stars; Lyne \& Lorimer 1994) suggests that they are young. Similarly, only two Galactic SGRs are known, and it has been suggested that the SGR population census is nearly complete (Kouveliotou 1995). In summary, AXPs have long rotation periods, hot blackbody-like spectra ($kT \sim$ 0.5 keV) with $10^{35-36}$ erg s$^{-1}$ steady luminosities, and have thus far only been observed at X-ray wavelengths. A search for new AXPs in the ASCA database has turned up another candidate, which we refer to as \psr. | On the basis of its long rotation period, steady X-ray flux, steep spectral characteristics, and location in the Galactic plane ($|b| \le 0.5$), we classify \psr\ as an anomalous X-ray pulsar (see Table 1). The high foreground absorption suggests that the pulsar is distant, and its line of sight along the tangent to the Sagittarius-Carina and Scutum-Crux spiral arms, and the 4-kpc molecular ring justifies its enormous foreground absorption (see figure 4). Its $N_H$ is roughly twice that of the nearby remnant Kes 73 for which quoted distances lie between $10 - 20$ kpc (Blanton \& Helfand 1996). However, the disparity in $N_H$ does not in itself imply a great dissimilarity in distances. For instance, at 10 kpc the lines of sight vectors to these two objects are already separated by $\sim 100$ pc, the typical sizes and scale heights of dense giant molecular clouds which are likely to be responsible for most of the absorbing gas. For the purposes of this article we assume the distance to be 15 kpc, an estimate likely to be accurate to within a factor of two. The steep X-ray spectrum is characteristic of the Wien tail of a blackbody radiator. Using the best fit blackbody parameters the isotropic X-ray luminosity is $L_X \simeq 2.5\times 10^{35}d_{15}^2$ erg s$^{-1}$, the distance being 15$d_{15}$ kpc. In effect, the X-ray pulsations can be ascribed to the viewing of a rotating stellar hotspot of area $0.15 A_sd_{15}^2$, where $A_s$ is the area of a neutron star of radius 10 km; this estimate ignores any relativistic corrections to the inferred area. Note that the spectrum is unlike that of any accreting high-mass neutron star binary. Although such binaries have periods in the range 0.07 - 900 s, and sometimes go into low luminosity states with $L_X \sim 10^{35}$ erg s$^{-1}$, they generally display very hard spectra ($0.8 < \Gamma < 1.5$), and show stochastic variability on all time-scales (Nagase 1989; Koyama et al. 1989) as is generally seen in accretion powered sources. We find no evidence for such variability in our data. With an AXP classification in hand we can compare the properties of \psr\ with those of five other members of the AXP family in Table~1. A few years ago Schwentker (1994) reported weak 5-s pulsations from RX J1838.4$-$0301 which have not yet been confirmed, Mereghetti et al. (1997) have argued that this X-ray source might be due to coronal emission from a late type star. We, therefore, exclude this source from our list. Although only a future $\dot P$ measurement can help determine the linear spindown age of \psr\ (an estimator for the age of an isolated neutron star), its location in the Galactic plane suggests that it is young, $\tau < 10^5$ yr. We consider it extremely likely that the pulsar is associated with one of the several star-formation complexes expected to lie along this line-of-sight (see Fig 4), out to a distance of 20 kpc. The pulsar lies along a rich region of the Galaxy; there are 10 supernova remnants, several radio pulsars, and the $\gamma$-ray source GRO J1838$-$04, all within a $3\times3$ deg$^2$ patch of sky surrounding \psr. However, to the best of our knowledge no cataloged sources are associated with it. We now describe other models that address the unique properties of AXPs. It is often noted in the literature, that the inferred accretion rate for the pulsars in Table~1 are close to those expected for accretion powered pulsars with field strengths $B \sim 10^{11-12}$~G, spinning at their equilibrium periods $P_{eq}$ (see Bhattacharya \& van den Heuvel 1991). This motivated Mereghetti \& Stella (1995) to suggest that these pulsars are members of a subclass of low mass X-ray binaries (LMXBs) in equilibrium rotation, with the stellar magnetic field of order $B_s \sim 10^{11}$ G. In contrast, van Paradijs et al. (1995) argue that these objects are isolated neutron stars accreting from a fossil disk, while Ghosh et al. (1997) suggest that AXPs are formed as the result of a Thorne-$\dot{\rm{Z}}$ytkov phase of a high mass X-ray binary with strong spherical accretion leading to the soft X-ray spectra with high foreground absorption. It is worth mentioning that accretion scenarios would be hard-pressed to explain the spin-down age of at least one member of Table 1, the $\sim 2000$ yr of the pulsar in Kes~73 (see Table 1; Gotthelf \& Vasisht 1997). First, it is difficult for accretion torques to spin-down a pulsar to 12-s in $\sim 10^3$ yr from initial periods $P_i \simlt 10^2$ ms unless, of course, the pulsar were born a very slow rotator, which is quite interesting in its own right. Secondly, if the pulsar were rotating near its equilibrium period, as in the Ghosh and Lamb (1979) scenario, the spin-down time of ${P/ 2\dot P} \sim 3900$ yr is inconsistent with the luminosity implied accretion rate, $\dot M \simeq 10^{-11}$ M$_\odot$ yr$^{-1}$ (assuming the pulsar has a standard dipolar field $\simeq 10^{12}$ G); these usually lie in range $10^4 - 10^5$ yr. In conclusion, further X-ray observations are required to secure the classification of the 7-s pulsar to the growing family of AXPs - by measuring the long term stability of the X-ray flux, secular trends in the pulse period including Doppler modulation, the lack of which will firm up the likelihood against an accreting binary hypothesis. Indeed, if AXPs are akin to SGRs then they could display sporadic hard X-ray transients, although such behavior is yet to be observed. Infrared observations to search for a possible counterpart to \psr\ could be carried out in spite the somewhat crude localization; we mention that past optical/IR searches for counterparts have been unsuccessful. Furthermore, radio observations to uncover an associated supernova remnant could be pursued. The high foreground absorption could easily cloak the soft X-ray emission of an aged, few$\times 10^4$ yr-old, remnant. We find it remarkable that the period of \psr\ agrees so well with that of 1E~2259+586, although at present we believe this is sheer coincidence. | 98 | 4 | astro-ph9804025_arXiv.txt |
9804 | astro-ph9804213_arXiv.txt | In efforts to demonstrate the linear Hubble law $ v = H r $ from galaxy observations, the underlying simplicity is often obscured by complexities arising from magnitude-limited data. In this paper we point out a simple but previously unremarked fact: that the shapes and orientations of structures in redshift space contain in themselves independent information about the cosmological redshift-distance relation. The orientations of voids in the CfA slice support the Hubble law, giving a redshift-distance power index $ p = 0.83 \pm 0.36 $ (void data from Slezak, de Lapparent, \& Bijoui 1993) or $ p = 0.99 \pm 0.38 $ (void data from Malik \& Subramanian 1997). | Hubble's (1929) observation that redshift increases linearly with distance for nearby galaxies has now been known for almost seventy years, and it is likely that its validity is not doubted by many. Yet, any attempt to demonstrate this simple law from galaxy observations soon descends into the complexities of magnitude-limited observations, the broad galaxy luminosity function, Malmquist bias, and, at larger redshifts, $K$-corrections, evolution, etc. This has permitted an often-ignored but persistent challenge to a linear redshift law from those preferring a quadratic relation, with both challengers and supporters citing data ranging from relatively nearby, bright optical galaxies (Soneira 1979, Segal 1980) through the {\it IRAS} 1.2~Jy redshift catalog (Segal et al. 1993; Koranyi \& Strauss 1997). One complication in understanding the classical Hubble diagram has been inhomogeneity, appearing as a dependence of density on radius or redshift or in clustering and peculiar velocities. With the advent of more galaxy redshift catalogs covering larger amounts of the sky, we obtain ever clearer pictures of the universe. The CfA ``slice'' (de Lapparent, Geller, \& Huchra 1986), in particular, first revealed dramatic structures, large voids separated by well-defined walls, extending to a significant fraction of the survey volume, perhaps even calling into question whether a survey to this depth represents a statistically fair sample of the universe. In other ways, however, this inhomogeneity itself can be useful. In this paper we use the shapes and orientations of structures to obtain information about the background cosmology in which they are embedded. In Section 2 below we discuss how shapes and orientations of structures in redshift space depend on the redshift-distance relation, and in Section 3 we apply these considerations to the CfA slice (de Lapparent, Geller, \& Huchra 1986) and the Las Campanas redshift survey (Shectman et al 1996, LCRS). Section 4 contains a final discussion. | We have shown in this paper how the apparent shapes and orientations of objects in redshift space can be used to determine the redshift-distance relation. When redshift is not linear in distance, objects such as voids that are intrinsically round in space when viewed in redshift space are stretched along the line of sight; and an initially isotropic distribution of orientations becomes distorted in the radial direction, with measurable effect, as in \eq{mu2d}. Application to voids in the CfA slice gives redshift power index $ p = 0.83 \pm 0.36 $ from Slezak et al. (1993) and $ p = 0.99 \pm 0.38 $ from Malik \& Subramanian (1997), both a modest preference for $ p = 1 $ over $ p = 2 $. To obtain these results in redshift space with a true $ p = 2 $, voids in the CfA slice in space would have to be flattened and preferentially aligned transverse to the line of sight. The main complication to the simple interpretation is likely to come from peculiar velocities; an expansion velocity will introduce an apparent line-of-sight elongation. However, this effect is expected to be small: in the most extreme case, of an uncompensated, completely empty, isolated void in an $ \Omega_0 = 1 $ universe, expansion would make $ p = 1 $ appear to be $ p = \case43 $. In any case, unless voids are contracting, a condition difficult to make physical sense of, peculiar velocities will only increase, not decrease, the apparent value of $p$, and will not distort a quadratic redshift law to appear to prefer $ p = 1 $. We do not expect void orientation statistics to replace classical methods for demonstrating the Hubble law. The most precise confirmation of the linearity of the Hubble law, $ p/5 = 0.2010 \pm 0.0035 $, has been obtained recently from the Hubble diagram of type Ia supernovae (Riess, Press, \& Kirshner 1996). The luminosity function as a function of redshift and the radial dependence of galaxy density in the {\it IRAS} 1.2-Jy catalog (Koranyi \& Strauss 1997) also support the linear Hubble law. However, our method offers an independent verification of the Hubble law in which magnitude-limited data, Malmquist bias, and evolution do not present problems. Further examination of the effect in real data, especially including in detail the peculiar velocities of the void walls, will undoubtedly introduce complications to the simple relations in equations (\ref{mu3d}) and (\ref{mu2d}). Whatever these complications, they will be different from those that enter arguments over the redshift-magnitude relation. | 98 | 4 | astro-ph9804213_arXiv.txt |
9804 | astro-ph9804163_arXiv.txt | We present radio observations of the gamma-ray burster \grb\ made with the Very Large Array (VLA) and the Owens Valley Radio Observatory (OVRO) spanning a range of postburst timescales from one to 300 days. A search for a time-variable radio source was conducted covering an area which included a fading X-ray source and an optical transient, both of which are thought to be the long wavelength counterparts to the gamma-ray burst. At the position of the optical transient sensitive limits between 10 $\mu$Jy and 1 mJy can be placed on the absence of a radio counterpart to \grb\ between 1.4 and 240 GHz. We apply a simple formulation of a fireball model which has been used with some success to reproduce the behavior of the optical and X-ray light curves. Using this model we conclude that the radio non-detections are consistent with the peak flux density of the afterglow lying between 20-40 $\mu$Jy and it requires that the optical flux peaked between 4 and 16 hours after the burst. | The gamma-ray burst of 28 February 1997 was a turning point in our understanding of these enigmatic objects, resulting in the first-ever discovery of X-ray and optical counterparts to a burst. Within the original 3-arcminute localization provided by the Wide Field Cameras (WFC) on board the BeppoSAX satellite, a previously unknown X-ray source 1SAX\ts{J0501.7+1146} was detected by the Narrow Field Instruments (NFI). Eight hours after the burst the flux of 1SAX\ts{J0501.7+1146} was 2.8$\times{10}^{-12}$ erg cm$^{-2}$ s$^{-1}$ (2-10 keV), but three days later its flux had dropped by a factor of 20 (Costa et al. 1997). A comparison of V- and I-band images taken on 28 February and 8 March revealed an optical transient within a reduced error box, defined by the intersection of the WFC circle, the $\pm$50\arcsec\ circle of the NFI, and the Interplanetary Network (IPN) annulus (van Paradijs et al. 1997, Hurley et al. 1997). Predictions of long-lived afterglows from gamma-ray bursts at X-ray, optical and radio wavelengths have been made for some time (e.g. Paczy\'nski \& Rhoads 1993, M\'esz\'aros, Rees, \& Papathanassiou 1994, Katz 1994, M\'esz\'aros \& Rees 1997). In particular, a power-law decay in the observed long-wavelength flux with time is a generic consequence of a class of models known as ``fireballs''. The gamma-rays are thought to be produced when the relativistically expanding blast wave (aka fireball) is slowed down by the ambient gas or is self-shocked by its own ejecta. The fireball accelerates particles in a shock which then radiate via the synchrotron process (M\'esz\'aros et al. 1994, Waxman 1997a, 1997b, Sari, Piran, \& Narayan 1998). As the fireball expands it cools, shifting the peak in the spectrum to lower energies and resulting in delayed emission at longer wavelengths. The optical and X-ray decay from \grb, as well as that from several other subsequent GRBs, is consistent at least to first order with one of the simpler formulations of these models (M\'esz\'aros \& Rees 1997). Costa et al. (1997) fit a t$^\delta$ decay to the X-ray data with $\delta\simeq-1.33\pm{0.12}$, while global fits to the X-ray, optical and infrared data give $\delta=-1.2$ (Wijers, Rees and M\'esz\'aros 1998) and $\delta=-1.09\pm0.23$ (Reichart 1997). More recent optical fits, aided by a long time-baseline, yield $\delta=-1.12\pm0.08$ (Garcia et al. 1997), $\delta=-1.21\pm0.02$ (Masetti et al. (1997), and $\delta=-1.10\pm0.04$ (Galama et al. 1998). The character of this decay agrees well with a fireball produced by a one-time impulsive injection of energy in which only the forward blast wave efficiently accelerates particles (i.e. the adiabatic piston model). The adiabatic model predicts a simple relation between the slope of the temporal decay $\delta$ and the flux spectrum $\beta$=${2\over3}\delta$. Observationally the spectrum is not well determined, with $\beta=-0.7\pm0.1$ optically (van Paradijs et al. 1997) and X-ray values of $\beta\simeq-0.9$ with large scatter (Frontera et al. 1998). The slope of the particle spectrum $p$ (where $\beta=(p-1)/2)$ inferred from these values of temporal decay is $p\sim-2.6$, not an unreasonable value for a relativistic shock (Blandford and Eichler 1987). Given the early success of the fireball model in predicting the gross properties of the optical and X-ray behavior of \grb, it is reasonable to look for delayed radio emission from this burst. This is particularly relevant in the light of the discovery of the radio afterglow from GRB\thinspace{970508} (Frail et al. 1997a) which exhibited temporal and spectral behavior consistent with this model (Waxman, Kulkarni \& Frail 1997). The properties of the fireball for \grb\ are well constrained by the optical and X-ray data. Thus the presence or absence of radio emission at late times has a bearing on the validity of this model. With this in mind, we began a radio search centered on the optical transient detected by van Paradijs et al. (1997). This {\it Letter} is a summary of our monitoring program for the first 300 days. | The detection of an optical transient, thought to be the afterglow from \grb, has made it possible to perform a search for time-variable radio emission at the optical position. Delayed emission at longer wavelengths (X-ray, optical and radio) is a generic prediction of all fireball models. VLA and OVRO observations have been made that span a range of postburst timescales from one to 300 days putting limits on the absence of a radio counterpart to \grb\ between 10 $\mu$Jy and 1 mJy. Applying a simple version of the fireball model which has been used successfully to fit the temporal behavior of the decaying X-ray and optical emission from \grb\ suggests that the radio afterglow has yet to peak. If this is correct then continued deep imaging in the coming months offers the possibility for the detection of a weak but increasing radio signal. Detecting this emission would be a powerful verification of the fireball model as described by M\'esz\'aros \& Rees (1997). Indeed, as the recent detection of GRB\thinspace{970508} has taught us, radio afterglows yield unique GRB diagnostics that are not obtainable by any other means (Frail et al. 1997b). Unlike optical or X-ray wavelengths one is presented with the possibility of following the {\it full} evolution of the fireball emission through all its different stages; first while it is optically thick, then as it slowly rises to a peak flux density and thereafter decays, making a transition from an ultra-relativistic to sub-relativistic shock. Furthermore, both the scintillation of the radio source (Goodman 1997) and its flux density, when it is synchrotron self-absorbed (Katz 1994), allow a determination of the size and expansion of the fireball. If the radio afterglow is going to be detected from \grb\ and others like it, then continued long-term monitoring is going to be required at the microJansky level. | 98 | 4 | astro-ph9804163_arXiv.txt |
9804 | astro-ph9804040_arXiv.txt | Few properties of astronomical masers are determined directly by observations, most are inferred indirectly. Foremost among the latter is the maser saturation stage. Saturation has a significant impact on maser growth, so determining whether a maser is saturated ($J > J_s$, where $J$ is the angle-averaged intensity and $J_s$ is the saturation intensity) or not is usually a precondition for analysis of the observations. Unfortunately, this crucial issue is not convincingly settled. Strong masers are generally believed to be saturated, but in most cases the evidence is less than compelling as it relies primarily on plausibility arguments rather than quantitative tests (see e.g.\ \[Book], \S 8.6). This unsatisfactory situation reflects a fundamental difficulty --- neither $J$ nor $J_s$ is directly measurable. The saturation parameter $J_s$ is a theoretical quantity, determined only within the context of a given pumping scheme. And because maser radiation is highly beamed, $J = I\Omega/4\pi$ so this quantity, too, cannot be directly measured; the intensity $I$ is measurable when the maser is resolved, but the beaming angle $\Omega$ is unobservable. Similarly, the amplification optical depth has never been directly determined for any maser that amplifies its own radiation. Recent VLA observations of OH 1720 MHz masers near the Galactic center by \[YZ96] open up new possibilities for direct determination of some maser properties. Significant circular polarization (upward of 20\%) is detected in various spectral features, and the right- and left-hand components coincide on the sky, as expected from the Zeeman effect. Furthermore, the spectral shape of the Stokes parameter $V$ follows an antisymmetric S-curve with sharp reversal at line center, the typical profile for Zeeman shift \DnuB\ much smaller than the Doppler linewidth \DnuD. Similar results were previously reported for H$_2$O masers in star-forming regions by \[FbG] and for OH 1612 MHz masers in OH/IR stars by \[ZFix], but the polarization was lower and the quality of the data not nearly as high. The general maser polarization solution was recently derived for arbitrary values of \eq{ \xb = {\DnuB \over \DnuD} } (Elitzur 1996, hereafter \[E96]) and the solution properties at $\xb \ll 1$ closely match the observed circular polarization. Here I show that a comparative analysis of the spectral profiles of $I$ and $V$, two measurable independent maser intensities, offers direct determination of various maser properties, in particular the saturation stage. The analysis is readily performed with the aid of the ratio profile \eq{\label{R0} \R(\nu) = {V(\nu) \over I'(\nu)} = {v(\nu)\over I'(\nu)/I(\nu)} } where the prime denotes derivative with respect to frequency and $v = V/I$ is the fractional circular polarization. When $\xb \ll 1$, spectral analysis of \R\ offers intrinsic sensitivity of order \xb\ and has long been an important tool in studies of the Zeeman effect of thermal radiation (see e.g.\ \[Trol]). In that case \R\ is constant across the spectral line and its magnitude determines the magnetic field along the line of sight. This constancy of \R\ follows from some simple, general symmetry arguments as shown by \[Cru] (see also \S3 below). However, maser exponential amplification during unsaturated growth destroys both the underlying symmetry and the constancy of \R, the saturation process restores both. The key to the different behavior, and \R-profiles, in the two regimes is the narrowing of the maser line during unsaturated amplification and its rebroadening during saturation. The important differences between thermal and maser polarization are discussed in detail below. For completeness, some basic elements of the polarization theory developed in \[E96] are reproduced in \S2. The \R\ profile is discussed in \S3 for thermal radiation and in \S4 for maser radiation when $\xb \ll 1$. In \S5, circular polarization for fully resolved Zeeman patterns, $\xb > 1$, is discussed. The implications for observations are discussed in detail in \S6. | 98 | 4 | astro-ph9804040_arXiv.txt |
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9804 | astro-ph9804089_arXiv.txt | Stellar dynamics is almost unreasonably well suited for an implementation in terms of special-purpose hardware. Unlike the case of molecular dynamics, stellar dynamics deals exclusively with a long-range force, gravity, which leads to a computational cost scaling as the square of the number of stars involved. While special tricks can lead to a reduction of this cost from $\sim N^2$ to $\sim N\log N$ in the case of very large particle numbers, such tricks are not suitable for all areas within stellar dynamics. When a stellar system is close to equilibrium, and has a very high density, it still pays to compute all interactions on a star by star basis, even for $N=10^5$. Any $cN\log N$ approach would either gloss over the subtle net effects of near-canceling interactions, driving the evolution of such a system, or would carry a prohibitively large coefficient $c$. This paper presents a brief introduction to the stellar dynamics of dense stellar systems, aimed at researchers using special purpose computers in other branches of physics. | Stellar dynamics is the branch of astrophysics that studies the structure and evolution of collections of stars, from small groups to larger star clusters to entire galaxies and clusters of galaxies. The interactions between the individual stars can be modeled to a high degree of accuracy as Newtonian gravitational interactions between point masses. Only under extremely high densities, such as occurs in the nuclei of galaxies and the centers of the densest star clusters, do stars have a reasonable chance to undergo a physical collision during their life time. In contrast, a typical star, such as our own Sun, has a probability of only 1 in $10^{8}$ to undergo a collision with a neighboring star, during the remaining $5\times10^9$ years of its life. The most spectacular example of a dense stellar system within our own galaxy is the agglomeration of roughly a million stars within the inner parsec from the center (a parsec is a unit of length, equal to a few light years, and corresponds to a typical distance between stars in the solar neighborhood). These stars describe orbits around the black hole that resides in the very center of our galaxy. The black hole itself has a mass that is a few million times larger than the mass of the Sun. The density of stars around the black hole is a million times larger than the stellar density in a typical part of the galaxy, such as where we reside. A detailed stellar dynamical modeling of the center of our galaxy is still difficult, partly because the observations of this heavily obscured area have only recently become accurate enough to tell us the physical characteristics of the system, partly because of the interference of other physical effects, such as the presence of gas clouds and ongoing star formation. Before tackling the stellar dynamics of the nucleus of our galaxy, it is therefore prudent to start our attempts with a simpler system, such as is provided by the core of a dense globular cluster. While most of the stars in and around our galaxy are spread out throughout the disk, and to a lesser extent through the halo, there are more than one hundred isolated star clusters circling the galaxy, each containing of order $10^6$ stars. In a dozen or so of those globular clusters, as they are called because of their appearance, the central densities rival that of the density in the nucleus of our galaxy. However, the absence of a large black hole, as well as gas clouds and concomitant star formation, makes it far easier to study and model globular cluster cores in detail. In addition, recent observations, notably with the Hubble Space Telescope ({\it cf.} \cite{Guh96}), have resolved those cores into individual stars, something that has not been possible with ground-based observations. This paper sketches some of the progress made in the study of globular cluster cores, emphasizing the role of special-purpose computers. | 98 | 4 | astro-ph9804089_arXiv.txt |
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9804 | astro-ph9804276_arXiv.txt | Modelling gravity is a fundamental problem that must be tackled in $N$-body simulations of stellar systems, and satisfactory solutions require a deep understanding of the dynamical effects of softening. In a previous paper (Romeo 1997), we have devised a method for exploring such effects, and we have focused on two applications that reveal the dynamical differences between the most representative types of softened gravity. In the present paper we show that our method can be applied in another, more fruitful, way: for developing new ideas about softening. Indeed, it opens a {\it direct\/} route to the discovery of optimal types of softened gravity for given dynamical requirements, and thus to the accomplishment of a physically consistent modelling of disc galaxies, even in the presence of a cold interstellar gaseous component and in situations that demand anisotropic resolution. | $N$-body simulations of disc galaxies rely on the use of softening. This artifice removes the short-range singularity of the gravitational interaction, which is dynamically unimportant and computationally troublesome, whereas it leaves the long-range behaviour of gravity unchanged. But softening is also a critical factor in simulations. It controls their quality and can affect their result on scales much larger than the softening length. Its dynamical effects are further exacerbated in the presence of a cold interstellar gaseous component and in situations that demand anisotropic resolution. Thus softening poses a dynamical problem of special concern, which should be probed carefully and in detail (e.g., Hernquist \& Barnes 1990; Pfenniger \& Friedli 1993; Romeo 1994, hereafter Paper I; Romeo 1997, hereafter Paper II% \footnote{Sections and equations of that paper are denoted by the prefix II.}; and references therein). In Paper I, we have investigated how faithful simulations are. In particular, we have concluded that the standard way of introducing softening in the presence of stars and cold interstellar gas is definitely unsatisfactory in several regimes of astrophysical interest. It is so because important small-scale instabilities of the gaseous component, e.g.\ those peculiar to star-formation processes, are suppressed just as unphysical noise of the stellar component. Faithfulness requires an appropriate introduction of two softening lengths, one for each component, and also a rigorous specification of the star-gas gravitational interaction. In Paper II, we have devised a method for exploring the dynamical effects of softening. As a major result, we have shown how to choose the softening length for optimizing the faithfulness of simulations to the Newtonian dynamics. Then we have focused on two applications that reveal the dynamical differences between the most representative types of softened gravity. In particular, we have concluded that it is desirable to improve the current way of introducing anisotropic softening. We need a clearer decoupling of the resolution parallel and perpendicular to the plane, and also more natural planar and vertical softening lengths. In the present paper, which completes our planned research work about softening, we propose an {\it innovative\/} solution to the problem. The understanding of galactic and extragalactic astrophysics is at a crucial stage. Unsolved problems are viewed in new perspectives, which promise major revisions of knowledge (see, e.g., Blitz \& Teuben 1996; Block \& Greenberg 1996). Recent investigations suggest, for instance, a more enigmatic interplay between stellar disc and bulge/halo (e.g., Lequeux et al.\ 1995), a clearer relation between cold gas and dark matter in spiral galaxies (e.g., Pfenniger et al.\ 1994; Pfenniger \& Combes 1994; Combes \& Pfenniger 1997), and a closer connection between the fractal structures of the interstellar medium and of the universe (e.g., de Vega et al.\ 1996, 1998). The implications are clear: modelling gravity in $N$-body simulations of disc galaxies should offer a flexible interface with such a progress. Our solution is to optimize the fidelity of simulations to given dynamical requirements. How do we apply this idea in practice? \begin{enumerate} \item We impose the requirements in the wavenumber space since this is the natural dynamical domain of gravity, as Pfenniger \& Friedli (1993) have previously emphasized. \item We identify the softening length with the characteristic dynamical scale length. \item Then we invert part of the method of Paper II, and the result is the optimal type of softened gravity that satisfies those dynamical requirements. \end{enumerate} Our application covers both 2-D and 3-D modelling. The basic cases are extended to more complex situations through recipes for implementing star-gas and anisotropic softening, which have already been motivated (cf.\ discussions of Papers I and II). Last but not least, each description is complemented by an example that leaves room for creativity. The present paper is organized as follows. The application is shown in Sects.\ 2 and 3 (see also Appendix A), and proceeds as in the previous discussion. Comments on related works concerning softening are made in Sect.\ 4. The conclusions and perspectives are drawn in Sect.\ 5, where we present our three papers about softening in a more unified view and emphasize their potentially strong impact on galactic dynamics. \begin{figure*} \vbox{\vspace{.1cm} \hbox{\hspace{-.25cm} \psfig{figure=romeof1.ps,width=18.9cm,height=13.275cm,angle=-90}} \vspace{-2.9cm}} \hfill\parbox[b]{5.7cm}{\caption[]{Examples of 2-D modelling: {\bf a} one-com\-pon\-ent case (cf.\ Sect.\ 2.1), {\bf b} two-com\-pon\-ent case (cf.\ Sect.\ 2.2). The abbreviations N, T and P mean Newtonian gravity, thickness and Plummer softening, respectively}} \end{figure*} \begin{figure*} \vbox{\vspace{.1cm} \hbox{\hspace{-.25cm} \psfig{figure=romeof2.ps,width=18.9cm,height=13.275cm,angle=-90}} \vspace{-2.9cm}} \hfill\parbox[b]{5.7cm}{\caption[]{Examples of 3-D modelling: {\bf a} iso\-tropic case (cf.\ Sect.\ 3.1), {\bf b} an\-iso\-tropic case (cf.\ Sect.\ 3.2). The abbreviations N, T and P mean Newtonian gravity, thickness and Plummer softening, respectively}} \end{figure*} | The importance of computer simulations in astrophysics is analogous to that of experiments in other branches of physics. They also serve as a welcome bridge between theories, often restricted to idealized situations, and observations, revealing instead the complexity of nature. Major present objectives are to construct physically consistent $N$-body models of disc galaxies and to simulate their dynamical evolution, especially in regimes of spiral structure in which a fruitful comparison between theories and simulations can be made (e.g., Pfenniger \& Friedli 1991; Junqueira \& Combes 1996; Zhang 1996; Bottema \& Gerritsen 1997; Fuchs \& von Linden 1998; von Linden et al.\ 1998; Zhang 1998a, b). The construction of such models is indeed a difficult task which has not yet been fully accomplished, and which should eventually provide clues of vital importance to a number of open questions posed by both theories and observations. Our involvement has been threefold. In Paper I, we have recognized a fundamental problem posed by this research programme (for a concrete use of that analysis and for interesting remarks see, e.g., Junqueira \& Combes 1996). In Paper II, we have devised a method for solving this problem. In the present paper, we apply this method and solve the problem, thus laying the foundations of such a plan. The {\it major result\/} is that gravity can be modelled so as to optimize the fidelity of simulations, and the procedure is practicable. The following conclusions point up the whys and wherefores: \begin{enumerate} \item Optimization is performed with respect to arbitrary dynamical requirements and, in specific examples, with respect to the Newtonian dynamics. This enriches the modelling with an {\it unprecedented\/} degree of freedom, which has clear epistemological motivations (cf.\ Sect.\ 1, discussion of the present paper). \item Optimization is performed in the wavenumber space. This is the {\it appropriate\/} domain for imposing dynamical requirements on the modelling. \item Optimization concerns {\it both\/} the softening length {\it and\/} the type of softened gravity. \item Softening is conceived as a {\it double\/} artifice. The softened gravity and finite-sized particle conceptions are equivalent in the basic cases. Concerning more complex situations, the latter is particularly useful for implementing star-gas softening, whereas the former is particularly useful for implementing anisotropic softening. Thus both conceptions contribute towards the accomplishment of a physically consistent modelling. \end{enumerate} Our application is ready for a concrete use. An attractive idea is to employ a particle-particle code together with MD-GRAPE, a highly parallelized special-purpose computer for many-body simulations with an arbitrary central force (Fukushige et al.\ 1996). We can also employ a classical particle-mesh code. Then the dynamical effects of the grid are known and factorize as those of softening (e.g., Bouchet et al.\ 1985; Efstathiou et al.\ 1985; for a review see, e.g., Hockney \& Eastwood 1988). So essentially the application proceeds as in the present paper, but it may be useful to act directly on the wavenumber space (e.g., Tormen \& Bertschinger 1996). A more complex problem concerns tree codes, which have hierarchical structure and adaptive resolution over multiple scales (e.g., Hernquist 1987; for a review see, e.g., Pfalzner \& Gibbon 1996). The solution to that problem would need a more advanced analysis (cf.\ following discussion). Welcome suggestions about the choice of the code can come from cosmological simulations (e.g., Splinter et al.\ 1998). Finally, what about the future? Our approach is connected with the technique of filtering in spectral domain used in the context of digital image processing. This is a rapidly evolving field with growing applications in science and engineering, which can promote further substantial advances in $N$-body modelling of disc galaxies. For instance, wavelets are ideal for resolving multi-scale problems in space and/or time, such as those concerning turbulence, bifurcations, fractals and many others (see, e.g., Kaiser 1994; Holschneider 1995; Bowman \& Newell 1998; for an alternative analysis tool see, e.g., Stutzki et al.\ 1998). Speculating further, wavelets might be used for speeding up simulations through fast solution of linear systems (cf.\ Press et al.\ 1992, pp.\ 597--599 and 782). These are the merits of our contribution. We hope that the trilogy (Papers I--III) and further reflections (Romeo 1998) will encourage $N$-body experimenters to model gravity so as to optimize the fidelity of their simulations, and that the result will be a stronger interdisciplinary connection with theories and observations. | 98 | 4 | astro-ph9804276_arXiv.txt |
9804 | astro-ph9804106_arXiv.txt | We propose that the majority of quasars at redshift $z\sim 1 - 5$ formed in the environment of new born collapsed halos with 1-D velocity dispersion $\sigma_v^{1d} \sim 400 \kms$. The harboring coefficient $f$ of quasars per halo and the lifetime of quasars depend only on local process, not modulated by the density inhomogeneities on scales larger than the size of the halos. Thus, the bias of quasars on scale larger than the size of these halos is mainly determined by the parameter $\sigma_v$ used for quasar environment identification. With this model, the popular structure formation models, like SCDM and LCDM, can be fairly well reconciled with the data of quasars, including a. observed feature of the environment for quasars; b. redshift evolution of quasar abundance; c. the two-point correlation functions of quasars. This bias model predicts that the correlation function of quasars doesn't significantly evolve, or only slightly increases with redshift. | \bigskip Mass distribution at high redshifts is being a hot subject of the large scale structure study. Data of various objects at moderate and high redshifts, in particular, clusters of galaxies, are becoming available for probing the formation and evolution of structures and for discriminating among popular dark matter models (e.g. Jing \& Fang 1994; Eke, Cole \& Frenk 1996; Bahcall, Fan \& Cen 1997; Kitayama \& Suto 1997). Considering that quasars are the most distant among various luminous objects, they have also been applied in this approach (e.g. Bi \& Fang 1997). However, as a mass tracer of the cosmic matter field, quasars are still playing a role different from clusters of galaxies. The problem stems from so-called ``bias''. Clusters are a biased tracer of the mass distribution. The correlation amplitude of clusters is believed to be much higher than that of the underlying matter and increases strongly with the cluster richness (Bahcall \& Soneira 1983). This bias is plausibly explained by the mechanism that the observed clusters are identified as massive collapsed halos of the density field (Kaiser 1984). That is, the bias of clusters is completely determined by the gravitational parameters, like mass $M$ and virial radius $r_{vir}$ used to identify the halos. With this approach, a detailed confrontation can be made between theories and the observations of clusters. Quasars may also be biased tracer of the mass distribution. Recent observations indicate that the correlation amplitude of quasars may also be different from the underlying dark matter (Mo \& Fang 1993; Komberg, Kravtsov \& Lukash 1994; Croom \& Shanks 1996; Franca, Andreani \& Cristiani, 1997). However, so far no detailed model is available for the bias of quasars, though their high clustering strength and environment imply that quasars are hosted by massive halos (see below). Because of the lack of such a model, one cannot confront the data of quasars with theoretical models in the way as for clusters. For instance, the abundance of quasars can only be used as an upper or lower limit to certain massive halos; no detailed comparison between the number densities of quasars and of halos is allowed. Obviously, it is very important to understand what kind of mass halos are associated with the majority of quasars. Such a knowledge will not only enable the observational data of quasars to be powerful tests for theoretical models of galaxy formation but also tell that what type of local environments is responsible for intriguing the nuclear activities of quasars. Like clusters and groups of galaxies, it is generally believed that quasars should be associated with certain type of collapsed dark matter halos. Yet, different from identification of clusters, the environment suitable for forming quasars is not merely determined by gravitational parameters, as the hydro processes are also involved. Therefore, the identification of quasar-harboring halos should be given by both gravitational and hydro parameters. In other words, not all halos with certain $M$ and/or velocity dispersion $\sigma_v$ harbor quasars, because certain hydro conditions must be satisfied. However, considering the hydro processes are local, it is reasonable to assume that the hydro conditions may not be modulated by the density inhomogeneities on scales much larger than the size of the halo $l$. In this case, the probability for a halo to have a quasar should be the same for all halos of the same kind, without depending on structures larger than $l$. Thus, the relative fractions of quasars with respect to the certain collapsed halos should be the same for all volumes larger than $l^3$. Consequently, when averaged on scale larger than $l$, the distribution of quasars $n_{qso}({\bf r},z)$ at redshift $z$ should be proportional to that of the considered halos, $n_{halo}({\bf r},z)$. Thus, all effects of the hydro processes can be absorbed into a ``normalization factor" $A$, i.e. $n_{qso}({\bf r},z)=An_{halo}({\bf r},z)$, and $A$ is less dependent on $z$ than $n_{halo}({\bf r},z)$. The bias of quasar distribution with respect to the mass distribution is then dominated by the bias of the selected halos with respect to the mass. Based on this analysis, quasar bias, at least on large scales, may also be only gravitational, depending on the gravitational parameters used for selecting the quasar-suitable halos. Accordingly, a possible model for quasar bias should at least satisfy the three conditions. 1. Gravitational environment given by the identified halos is consistent with the observed environment around quasars; 2. The abundance of quasars, $n_{qso}(z)$, at redshift $z$ is proportional to the number density of the identified halos, $n_{halo}(z)$ in a redshift-independent way, i.e. $n_{qso}(z)=An_{halo}$ where $A$ is a {\it z-independent} constant, 3. The amplitude and $z$-evolution of the halo-halo correlation function are consistent with the observed correlation function of quasars. In this letter, we will show within the framework of the CDM cosmogonic theories that such a bias model can indeed be settled following the above-mentioned points. The details of the points 1, 2 and 3 will be discussed in the \S 2, 3 and 4, respectively. | We showed that velocity-dispersion-selected halos are a possible mechanism for the bias of quasars. The majority of quasars at redshift $z\sim 1 - 5$ formed in the environment of new born collapsed halos with 1-D velocity dispersion $\sigma^{1d}_v \sim 400 \kms$. Both the harboring coefficient $f$ per halo and the lifetime of quasars are $z$-independent. With this bias model, the popular structure formation models, like SCDM and LCDM, can be fairly well reconciled with data of the abundance and correlations of quasars at $z \geq 0.5$. It is interesting to point out that the velocity dispersion identified halos generally don't have the same mass. Eq.(2) shows that for a given $\sigma^{1d}_v$, the redshifts the higher, the mass of the halos the smaller. This result has already been recognized in an earlier study, which shows that in order to fit with quasar abundance at high redshift, the mass of the halos has to be smaller than at the lower redshift (Bi \& Fang 1997). With this model, one can predict that 1. The environment for quasars at redshifts from $z \sim 1$ to 5 should be characterized by a velocity dispersion, $\sigma^{1d}_v \sim 400 \kms$; 2. The amplitudes of quasar two-point correlation function at high redshifts don't significantly evolve with redshifts. In the paper, only the models of the SCDM and LCDM are considered. We can expected that with better data of quasars becoming available, the bias model of quasars will play more important role for discriminating among models of structure formations. | 98 | 4 | astro-ph9804106_arXiv.txt |
9804 | astro-ph9804330_arXiv.txt | s{We review present understanding of Galactic free--free emission and its possible importance to CMB fluctuation measurements. Current results, from both ``direct'' observations in the microwave band and from H$\alpha$ studies, suggest that this foreground does not represent a serious obstacle to mapping the CMB; however, this is based on limited information and we emphasize the need for more exhaustive studies. We also present some preliminary results based on our recent H$% \alpha$ observations near the South Pole CMB data sets. The fluctuation amplitude seen in H$\alpha$ indicates that the detected CMB fluctuations are not significantly contaminated by free--free emission, at least if the diffuse gas is at a temperature of $T\sim 10^4$ K.} | To summarize the current status of our understanding of Galactic free--free emission, we would say that although there is {\em no indication} of fluctuations large enough to pose serious difficulties for CMB observations, the observational constraints remain weak. A critical interpretation of the results in the tables would be that the limit on large scale free--free fluctuations is the same order as the CMB amplitude on these same scales (at $~40$ GHz). On smaller scales, observations in $H\alpha$ have not turned up any signs of large amplitude variations, but those based on high resolution spectrographs are few and cover only a small percentage of the sky. There does appear to be a dust/free--free correlation on all angular scales, but there is room, and perhaps tentative indications of, an equally important non--correlated component (question \#1 posed in the introduction). And then there is the puzzeling result from Leitch et al. (1997), perhaps pointing to a hot phase of the ISM which could, due to lack of sensitivity, escape many of the present $H\alpha$ limits (question \#2 posed in the introduction). Obviously, CMB observations at higher frequencies, where much of the effort is now being concentrated, will suffer less from any possible free--free contamination, and the many of the next generation CMB experiments have a wide spectral coverage to aid the removal of foregrounds. Even given the above critical viewpoint, it would be a surprise to discover free--free emission presenting an important difficulty for all planned CMB experiments, at least in terms of measuring the variance, or power spectrum, from the early Universe. Foregrounds will, however, be much more important for higher order statistics looking for non--gaussian signatures. In such cases, the non--gaussian foregrounds will have to be removed to high precision. A sensitive, high spectral resolution H$\alpha$ survey of the entire sky would be of great value in the context of the above considerations. We have also reported some results from our recent H$\alpha$ observations taken with a Fabry--Perot system at La Silla. The data cover two bands along which the South Pole data sets show significant fluctuations in CMB brightness. The observed H$\alpha$ fluctuations indicate that the CMB results in this region are not significantly contaminated by free--free emission (assuming a gas with $T_4=1$). | 98 | 4 | astro-ph9804330_arXiv.txt |
|
9804 | astro-ph9804218_arXiv.txt | We present new spectroscopic and photometric time series observations of the $\delta$ Scuti star FG~Vir. We detect the oscillations via changes in the equivalent widths of hydrogen and metal absorption lines. {}From the ratios between spectroscopic and photometric amplitudes, we assign $\ell$ values to the eight strongest oscillation modes. In particular, we identify two radial modes ($\ell =0$) and find that the main pulsation mode (147~$\mu$Hz) has $\ell =1$. One of the radial modes (at 140$\mu$Hz) is the fundamental, implying that two modes with lower frequencies are {\it g}-modes. For the radial modes, we compare frequencies with those calculated from a scaled $\delta$~Scuti star model and derive a density $0.1645\pm 0.0005\,\rho_{\odot}$. We then obtain a distance of $84\pm 3$\,pc, in excellent agreement with the Hipparcos value. Finally, we suggest that a 3.5-day variability in all observables (equivalent widths and intensity) is caused by stellar rotation. | \label{intro} Oscillations in multi-periodic variables such as $\delta$ Scuti, roAp and $\beta$ Cephei stars have been observed extensively during the past 20 years. But even with high-quality data, it is still extremely difficult to identify which modes are being detected. Kjeldsen et al.\ (\cite{kbvf95}) used a new technique to detect solar-like oscillations in the bright G sub-giant $\eta$ Boo through their effect on the equivalent widths of the Balmer lines. A subsequent discussion by Bedding et al.\ (\cite{bkrb96}) of the sensitivity of different observables to modes with different degree $\ell$ suggested that one can determine the $\ell$-value of a given mode by combining measurements of absorption-line equivalent widths with simultaneous photometric observations. To test this idea, we chose the bright and well-studied $\delta$~Scuti star FG~Vir. FG Vir (HD 106384; $V=6.57$) is a multi-periodic $\delta$~Scuti star. It has a main pulsation period close to 1.9 hours and shows a fairly complex oscillation spectrum. This star has been studied extensively during the last few years, resulting in the detection of at least 24 well-determined frequencies between 100 and 400~$\mu$Hz, with amplitudes from 0.8 to 22 milli-magnitudes (mmag; Breger et al.\ \cite{bhn95}, \cite{br98}). Because of its slow rotation (which reduces the complicating effects of rotational splitting) and the large number of detected frequencies (some of which are probably {\it g} modes), FG Vir is one of the most promising candidates for performing asteroseismology on a $\delta$ Scuti variable. Observations and models of this star have been presented by Dawson et al.\ (\cite{dbl95}), Breger et al.\ (\cite{bhn95}) and Guzik \& Bradley (\cite{gb95}). By choosing a star like FG Vir we have the advantage of knowing the frequencies in advance. We are therefore able to determine the oscillation amplitudes and phases with high precision. The aim of this paper is to identify the $\ell$ values of the observed modes and to compare the oscillation frequencies with a pulsation model. A preliminary analysis of the observations presented in this paper was given by Viskum~et~al.~(\cite{vdk97}), while results on radial velocity measurements were given by Viskum et al.\ (\cite{vb97}). | We have investigated a new technique to measure the oscillations in $\delta$~Scuti stars via changes in the equivalent widths of absorption lines. An important advantage of this new technique is that only medium-dispersion spectra are needed, which makes the method suitable for small and medium-sized telescopes and for multi-site campaigns. Our main results are summarized below. \begin{itemize} \item Our detection of oscillations in FG~Vir from equivalent-width measurements of \Ha, \Hb\ and Fe\,{\sc i} lines is an important confirmation of the method developed by Kjeldsen et al.\ (\cite{kbvf95}) to search for solar-like oscillations. \item {}From the ratios between oscillation amplitudes measured in EW and the four Str\"omgren filters ({\it uvby}), we have identified $\ell$-values for eight modes in FG~Vir. \item We suggest that the two lowest-frequency modes are {\it g}-modes, while the strongest mode (147.2$\mu$Hz) is a dipole mode. \item {}By comparing the frequencies of radial modes with model calculations, we obtained a precise density and derived a distance that is in excellent agreement with the Hipparcos value. \item We detected a long-period variation in the time series with a period of 3.5 days, which we propose is caused by rotation of the star. \end{itemize} | 98 | 4 | astro-ph9804218_arXiv.txt |
9804 | astro-ph9804112_arXiv.txt | We discuss a technique for mapping the synchrotron turnover frequency distribution using nearly simultaneous, multi--frequency VLBI observations. The limitations of the technique arising from limited spatial sampl\-ing and frequency coverage are investigated. The errors caused by uneven spatial sampl\-ing of typical multi--frequency VLBA datasets are estimated through numerical simulations, and are shown to be of the order of 10\%, for pixels with the deconvolution ${\rm SNR} \sim 7$. The fitted spectral parameters are corrected for the errors due to limited frequency coverage of VLBI data. First results from mapping the turnover frequency distribution in \object{3C\,345} are presented. \keywords {methods: data analysis -- methods: observational -- quasars: individual: \object{3C\,345}} | } Information obtained with Very Long Baseline Interferometry (VLBI) about radio spectra of parsec--scale jets and their evolution can be crucial for distinguishing between various jet models. However, there are several aspects of VLBI which impede spectral studies of parsec--scale regions. The reliability of spectral information extracted from VLBI data depends on many factors including sampl\-ing functions at different frequencies, alignment of the images, calibration and self--calibration errors, a narrow range of observing frequencies, and source variability. The influences of all these factors must be understood and, if possible, corrected for, in order to reconstruct the spectral properties of parsec--scale jets consistently. Radio emission from the parsec--scale jets is commonly described by the synchrotron radiation from a relativistic plasma (e.g. Pacholczyk 1970). The corresponding spectral shape, $S(\nu) \propto \nu^{\alpha}$, is characterized by the location of spectral maximum ($S_{\rm m}$, $\nu_{\rm m}$) also called the turnover point, and by the two spectral indices, $\alpha_{\rm thick}$ (for frequencies $\nu \ll \nu_{\rm m}$) and $\alpha_{\rm thin}$ (for $\nu \gg \nu_{\rm m}$). In many kiloparsec--scale objects, spectral index distributions have been mapped, using observations made with scaled arrays. In such observations, the antenna configurations are selected at each frequency in a specific way such that the spatial sampl\-ings of the resulting interferometric measurements are identical at all frequencies used for the observations. It is virtually impossible to use the scaled array technique for VLBI observations of parsec--scale jets made at different frequencies. The uneven spatial sampl\-ings of VLBI data at different frequencies result in differences of the corresponding synthesized beams, and can ultimately lead to confusion and spurious features appearing in spectral index maps. In spectral index maps, the only available kind of information is the spectral slope between the two frequencies. While sufficient for many purposes, this information can be misleading in the situation when the frequency of thespectral maximum lies between the frequencies used for spectral index mapping. In the ranges of frequencies between 1.4 and 43\,GHz, frequently used for VLBI observations, such a situation can be quite common. Using observations at three or more frequencies, it is possible to estimate the shape of the synchrotron spectrum, and derive the turnover frequency (frequency of spectral maximum). Information about the turnover frequency can help to avoid the confusion which is likely to occur in spectral index maps. The turnover frequency is sensitive to changes of physical conditions in the jet such as velocity, particle density, and magnetic field strength. This makes it an excellent tool for probing the physics of the jet in more detail than is allowed by analysis of the flux and spectral index properties of the jet. In this paper, we present a technique suitable for determining the turnover frequency distribution from multi--frequency VLBI data, and investigate its limitations and ranges of applicability. We discuss the advantages of using the Very Long Baseline Array\footnote{The Very Long Baseline Array is operated by the National Radio Astronomy Observatory (NRAO)} (VLBA) for spectral imaging. A general approach to imaging of VLBA data from nearly simultaneous, snapshot--type observations at different frequencies is outlined in section~\ref{sc:imaging}. The effects of limited sampl\-ing and uneven {\it uv}--coverages are discussed in section~\ref{sc:spsens}. We provide analytical estimates of the sensitivity decrease, and use numerical simulations to evaluate the effect the uneven spatial sampl\-ings have on the outcome of a comparison of VLBI images at different frequencies. Alignment of VLBI images is reviewed in section~\ref{sc:imalign}. A method used for spectral fitting and determining the turnover frequency is described in section~\ref{sc:fitting}. Spectral fitting in the case of limited frequency coverage is discussed in section~\ref{sc:frcoverage}. The first results from the turnover frequency mapping are presented in section~\ref{sc:algorythm}. | } In this paper, we have covered several methodological and scientific aspects of studying synchrotron spectrum of the parsec--scale regions in AGN. The main conclusions can be stated as follows: 1)~We have discussed a technique that can be used for mapping the turnover frequency distribution and obtaining spectral information from multi--frequency VLBA data. A feasibility study shows that multi--frequency VLBA observations can be used for spectral imaging and continuous spectral fitting. 2)~Multi--frequency VLBA observations made with up to 10 minute separations between the scans at each frequency can provide a satisfactory spatial sampl\-ing and image sensitivity for sufficiently bright sources with intermediate ($\sim 10$--15\,mas) structures. The fractional errors from comparing the data at different frequencies should not exceed 10\% for emission with SNR$\ge 7$, in this case. 4)~A procedure for broadband synchrotron spectrum fitting has been introduced for mapping the distribution of spectral parameters of radio emission from parsec--scale jets. Corrections based on the local curvature of the fitted spectra are introduced, in order to compensate for the incomplete frequency coverage in cases where the true turnover frequency is outside of the range of observing frequencies. 5)~From a 4--frequency VLBA observation of \object{3C\,345}, the first map of the turnover frequency distribution are produced. The maps indicate possible locations of the relativistic channel and strong shock fronts inside the jet. The magnetic field distribution derived from the turnover frequency and flux distributions is consistent with the plane shocks existing in the immediate vicinity of the source core. The extended emission appears to have a very low turnover frequency for which the existing data do not warrant a good estimate, limiting the conclusions to deducing certain information from the gradients of the turnover frequency which are visible in the extended jet. The observed gradients are consistent with the patterns of velocity distribution and density gradients typical for Kelvin--Helmholtz instabilities propagating in a relativistic jet. A more detailed study, with observations made at lower frequencies, is required for making conclusive statements about the nature of the observed gradients of the turnover frequency. | 98 | 4 | astro-ph9804112_arXiv.txt |
9804 | astro-ph9804262_arXiv.txt | The EGRET experiment on board CGRO revealed the existence of a diffuse gamma-ray background (hereafter DGRB) at the level of $I_{DGRB} = 9.6 \cdot 10^{-7} E_{GeV}^{-2.11 \pm 0.05} ~ cm^{-2} s^{-1} sr^{-1} GeV^{-1}$ \cite{owz94} in the energy range $0.03 \div 10$ GeV. However, a recent reanalysis of the EGRET data \cite{swz} found that the level of the DGRB is systematically lower by a factor $\sim 20 \%$ in the energy range $\sim 0.1 \div 4$ GeV. The DGRB is observed at high galactic latitudes $b > 10$ deg and such an evidence suggested an extragalactic origin for this diffuse background. Nonetheless, the specific origin of the DGRB is still under debate. In fact, the EGRET experiment \cite{kan} has a poor angular resolution ($\theta_{min} \sim 1$ deg) so that it is hard to discriminate among different origins of this extragalactic background. Specifically, it is still difficult to discriminate between a purely diffuse nature of the DGRB (see \eg \cite{ca}) and the option of a DGRB made by a superposition of unresolved, discrete sources. The large number of identified AGNs and flat spectrum radio quasars (hereafter FSRQ) in the EGRET sky (\cite{fi96} \cite{has} \cite{mat}) suggested that most of the DGRB can be produced by a non-resolved population of AGNs, the actual fraction of the DGRB produced by FSRQ and BLLacs being in the range $\sim 40 \div 95 \%$ (see \eg \cite{com}). Separately, it has been evaluated that a fraction $\sim 42 \div 97 \%$ of the DGRB could be ascribed to blazars \cite{pad}. However, the flatness of the spectrum of the DGRB seems to favour the possibility that BLLacs could be the major contributors to the DGRB of extragalactic origin \cite{pohl}. Erlykin \ea \cite{erl96} reviewed the various AGN contributions and quoted that the fraction of the DGRB produced by the observed AGNs is $\sim 65 \%$. The DGRB fractions previously reported may be subject to a revision ($\sim 25 \%$ increase) if the recent reanalysis \cite{swz} of the EGRET data is adopted. On account of the large theoretical uncertainties and of the present observational precision of the EGRET detectors, it is still hard to discriminate among the different proposed possibilities, even though a fluctuation analysis of the EGRET data should give more precise indications on the nature and origin of the DGRB. Beside the discrete, unresolved source case pictured for the origin of the DGRB, there have been some pioneering works \cite{hw} \cite{ds} \cite{bbp} \cite{volk} suggesting that a relevant fraction of the DGRB could be produced by {\it extended} sources through hadronic collisions of cosmic ray (hereafter CR) protons interacting with the protons of the Inter Galactic Medium (hereafter IGM) which is abundantly present within galaxy clusters (see \cite{sa88} for a review). In this alternative picture, the CR's are assumed to be produced within clusters (we will discuss in Sect.4 some of the possible sources) where also a population of protons and electrons is residing in the form of a hot (with temperatures $T \sim 10^7 \div 10^8$ K), tenuous (with electron number densities $n_e \sim 10^{-3}~cm^{-3}$), chemically enriched and massive (with mass fractions $M_{IGM}/M \sim 0.05 \div 0.3$) plasma: the IGM. The proposed mechanism has an essential ingredient in the confinement of the CR's within clusters where they are produced; this point, already realized by some authors \cite{bbp} \cite{volk}, is responsible for the net increase in the probability of interaction per proton with respect to the case of a straight line propagation. The increase factor can be estimated to be $\sim c t_{cl} /R_{cl}\simgt 600$, where $t_{cl} \simlt H_0^{-1}$ is the age of the cluster and $R_{cl}$ is its size. Cosmic rays produced within a cluster during all its lifetime can thus produce gamma rays through the production and the subsequent decay of neutral pions: \begin{equation} p+p\to \pi^0+X~,~~~~~~~~~~~~~~\pi^0\to \gamma+\gamma. \end{equation} Note that in the same interactions, charged pions are also produced, which determine a neutrino emission through the following channels: \begin{equation} p+p\to\pi^{\pm}+X, ~~~~~\pi^{\pm}\to \mu^{\pm} \nu_{\mu}(\bar{\nu}_{\mu}), ~~~~~\mu^{\pm}\to e^{\pm} + \bar{\nu}_{\mu}(\nu_{\mu}) + \nu_e (\bar{\nu}_e) ~. \end{equation} We will also discuss the relevance of these last processes in Section 7 below. Using the gamma ray production from clusters of galaxies according to eq. (1), Houston \ea \cite{hw} suggested that the total extragalactic gamma ray intensity detected above $35$ MeV \cite{ft82}, $I_{\gamma} \approx 5.5 \cdot 10^{-5}~ cm^{-2} s^{-1} sr^{-1}$, could be ascribed, for a large fraction, to galaxy clusters. They predicted a level $I_{\gamma} \approx 5 \cdot 10^{-5} cm^{-2} s^{-1} sr^{-1}$ at energies above $35$ MeV, assuming an observed local cluster space density, $n_{cl} \approx 7.3 \cdot 10^{-5} Mpc^{-3}$, integrated out to the Hubble radius, $R_H=6 \cdot 10^3$ Mpc, and neglecting any cosmological effect. More recently, Dar \& Shaviv (hereafter DS \cite{ds}) reanalyzed the problem in the light of the EGRET data \cite{owz94} and calculated the contribution to the DGRB from CR interactions in the intracluster gas, under the assumption that the energy density of CR's in clusters is the same as in our own galaxy (universality). With this assumption, Dar \& Shaviv \cite{ds} predicted a level $I_{\gamma}(> 100~ MeV) \approx 1.2 \cdot 10^{-5}$ photons cm$^{-2}$ s$^{-1}$ sr$^{-1}$, which could explain the whole amount of the DGRB of extragalactic origin. In a following paper, Berezinsky, Blasi \& Ptuskin (hereafter BBP \cite{bbp}) relaxed the {\it ad hoc} assumption of universality, and estimated the CR energy density in clusters due to various possible sources of CR, using the condition of diffusive confinement of CRs. In their approach, BBP \cite{bbp} showed that it is impossible to fulfill the universality condition with the usual CR sources in clusters, emphasizing that the DGRB due to the CR interactions in clusters should be a small fraction of the total diffuse flux observed by EGRET. This conclusion was reached by the previous authors under the hypothesis that a large fraction of the baryons in the universe is contained inside clusters of galaxies (BBP considered that clusters are a fair sample of the baryons in the universe \cite{wf91}, \cite{wetal93}, \cite{wf95}) assumed to have a homogeneous inner distribution of gas, $n_e =const$ (here $n_e$ is the IGM electron number density). Because of these assumptions, their results depend only on overall cosmological parameters like the baryon fraction in the universe $\Omega_b$, and on the cluster size. Dar \& Shaviv \cite{ds} also predicted the gamma ray fluxes from a few nearby clusters (Coma, Perseus and Virgo): for these three clusters they found $\gamma$-ray fluxes in the range $F_{\gamma}(>100 ~MeV) \approx 5 \div 20 \cdot 10^{-8} cm^{-2} s^{-1}$. In particular, the value which they predicted for A1656 (Coma), $F_{\gamma}(>100 ~MeV) \approx 5 \cdot 10^{-8} cm^{-2} s^{-1}$, is close to - or slightly higher than - the upper limits given by EGRET for this source. Similar results were obtained for these clusters by Ensslin \ea \cite{ensslin96} assuming a population of CRs from radio sources located within galaxy clusters in almost equipartition with the IGM thermal energy. We stress here that in all the previous works a uniform IGM density profile was assumed. Moreover, the cluster population was not assumed to evolve with cosmic time, and the same working hypothesis of no-evolution was assumed for the IGM content of each cluster. However, X-ray studies of galaxy clusters, have shown that these cosmic structures are indeed well structured, having a gas density profile $n(r) \propto [1+(r/r_c)^2]^{- 3 \beta /2}$, with core radii $r_c \approx 0.1 \div 0.3 \hmpc$ and $\beta \approx 0.6 \div 0.8$ (see \eg \cite{jf92}; see also \cite{sa88} and references therein). Beside this, the IGM is indeed evolving as indicated by its sensitive metal enrichment, $Fe/H \sim 0.2 \div 0.5$ (in solar units, see \eg \cite{e90}, \cite{ar}), shown even for the brighter clusters observable at redshifts $z \sim 0.5$ \cite{lm97}. Nonetheless, there is also an increasing debate on the possible evolution of the X-ray luminosity function observed out to $z \simlt 0.5$ with the EINSTEIN \cite{gio90} \cite{h92} and ROSAT satellites \cite{ebe97} \cite{nich} and on the possible evolution of the cluster temperature function \cite{cmv} \cite{eke98} \cite {vl98}. If an evolution is present in the cluster population this can be, in fact, understood as a result of two competing effects: \newline {\it i)} a luminosity evolution, where the cluster X-ray luminosity, $L \propto n^2 T^{1/2} R^3$ (mainly due to thermal bremsstrahlung), changes with redshift due to variations in the gas mass density, $n \propto f_g \rho_{cl}$ (where the cluster gas mass is taken to be a fraction $f_g \equiv M_{gas}/M$ of the total cluster mass), and/or changes in the IGM temperature $T$ at fixed mass, $M \propto \rho_{cl} R^3$ (here $\rho_{cl}$ is the cluster total mass density); \newline {\it ii)} a change in redshift of the number density, $N(M,z)$, (usually referred to as mass function, hereafter MF) of clusters that are found to be collapsed (or virialized) in the mass range $M, M+dM$ at redshift $z$. Detailed studies of cluster evolution in X-rays (see \eg \cite{cv}, \cite{ob}, \cite{cmv}) considered in fact that a combination of the previous mechanisms is responsible for the actual cluster evolution when they fit the available data (see \cite{cv} for a detailed discussion). In this paper we predict the amount of high energy, non-thermal, gamma-ray emission from galaxy clusters using detailed modelling of the realistic cluster structure, as well as viable modelling for the evolution of the IGM and of the cluster MF. Based on these phenomenological cluster models, we predict the amount of DGRB that can be produced in the viable cosmological models: here we consider flat and low-density (open or vacuum-dominated) CDM models as well as mixed Dark Matter models with a fraction $\Omega_{\nu} \approx 0.3$ of the total density of the universe in form of massive neutrinos. We use $h=H_0/100$ km s$^{-1}$ Mpc$^{-1}$ throughout the paper unless otherwise specified. The plan of the paper is the following. In Sect.2 we briefly summarize the cluster formation hystory in hierarchical scenarios for structure formation. In Sect.3 we describe a model for the production of diffuse gamma-ray emission due to the interaction of CR's with the target protons present in the extended, diffuse IGM. We consider in Sect. 4 different CR sources that can be found in connection with galaxy clusters. We discuss in Sect.5 the correlation between the extended gamma-ray emission and the much better known thermal X-ray emission coming from the IGM. Based on these properties, we construct a list of predicted $\gamma$-ray fluxes for a compilation of X-ray clusters with detailed informations on their IGM structure, IGM temperatures and X-ray fluxes. In Sect. 6 we present predictions for the amount of DGRB produced by galaxy clusters in different cosmological scenarios. We briefly discuss in Sect.7 the extended neutrino fluxes emerging from these objects and their contribution to a possibly detectable diffuse neutrino background (hereafter DNB). Finally, in Sect.8 we discuss our results in the light of the current limits obtained from EGRET and in the light of the future experiments for gamma-ray and neutrino astronomy. | In this paper we presented a detailed study of the diffuse emission of $\gamma$-rays and neutrinos from clusters of galaxies. Using realistic modelling of the cluster structure, of their formation history and of their evolution with cosmic time, we found that galaxy clusters can provide $\simlt 1\%$ of the DGRB measured by EGRET (in the first release by OWZ \cite{owz94}). Our estimate of $I_{\gamma}$ is quite independent on the geometry of the universe, on the assumed cosmological model and on the amount of IGM evolution, because most of their contribution to the DGRB comes from nearby, $z \simlt 0.2$, clusters. In fact, at these redshifts the effects of curvature do not take place strongly in changing the perturbation growth factor, ${\cal D}(z,\Omega_0)$, (normalized at the present epoch), the difference in cluster evolution are small when the different models are normalized to the local abundance of clusters observed in X-rays and the available amount of IGM evolution - even if considered to be quite strong, $f_g \propto (1+z)^{-1 \div -2}$ - can provide only small variations to the cluster $\gamma$-ray luminosities, as $L_{\gamma} \propto f_g$ (see eq.18). On account of all these aspects, we consider that our results for the contribution of galaxy clusters to the DGRB are quite robust. Our approach differs substantially from the previous ones in several (among others) aspects: \par\noindent {\it i)} we considered - differently from all the previous approaches - a self-consistent approach to the formation of clusters following the spherical collapse model \cite{peeb80} complemented with a realistic IGM density profile, consistent with the most recent determinations from X-ray observations. This fact has important effects on the CR confinement within cluster cores and hence on the relative $\gamma$-ray and neutrino emission rates; \par\noindent {\it ii)} we considered (as BBP did) here an energy dependent diffusion coefficient which results in a very general picture of the CR confinement within cluster cores; \par\noindent {\it iii)} we also considered here - at variance with the previous approaches - the effects of a possible evolution in the cluster IGM content. This is consistent with the present indications of a variation in the IGM content from groups to rich clusters in the local frame and with the X-ray, shock (or entropy) induced, luminosity evolution observed from numerical simulations \cite{tm97} and predicted in analytical models (both shock and entropy models) for the evolution of X-ray clusters \cite{cola97} \cite{b97}; \par\noindent {\it iv)} we use the PS cluster MF that was found to be consistent with N-body simulations over a large dynamical range and up to $z \simgt 2 $ \cite{eke96}. We normalized it to the local abundance of clusters detected in X-rays. In the previous approaches average values for the overall cluster abundance, $n_{cl} \approx 4\div 7 \cdot 10^{-5}~ Mpc^{-3}$, were used \cite{hw} \cite{bbp} without considering the effect of an evolving cluster mass function; \par\noindent {\it v)} using a self-consistent modelling of the IGM we found an analytical correlation between $\gamma$-ray and $X$-ray emission for clusters. The predicted ratio $F_{\gamma} / F_X \propto f_g^{-1} r_c^{-1} T^{-1/2}$ [see eq.(32)] provides a behaviour of the $F_{\gamma}-F_X$ relation different from that obtained by Ensslin \ea \cite{ensslin96}, $F_{\gamma} / F_X \propto T^{1/2}$, because we did not assume any (partial) equipartition between IGM thermal energy and relativistic jet particles. Our correlation results only from the basic electromagnetic and hadronic emission mechanisms in which the IGM protons are the targets for both the X-ray thermal bremsstrahlung emission and for the $p p$ collisions responsible for $\gamma$-rays. \par\noindent {\it vi)} using such a correlation we derived a sample of nearby clusters with predicted $\gamma$-ray fluxes observable with the next generation $\gamma$-ray telescopes. Incidentally, we found a $\gamma$-ray flux for Coma, $F_{\gamma}(>100 MeV)\approx 8.5 \cdot 10^{-9}~ photons$ $s^{-1} ~cm^{-2}$ which is consistent with the EGRET upper limits for this cluster (previous specific predictions \cite{ds} \cite{ensslin96} seem to exceed the EGRET upper limit). Our numerical results for $I_{\gamma,\nu}$ are in reasonable agreement with those obtained by BBP, even though based on a quite different description of the cluster structure and evolution. This agreement is due to the fact that BBP considered a constant comoving density of clusters, $n_{cl} \sim 5 \cdot 10^{-5}$ Mpc$^{-3}$, assumed to be a fair sample of the baryons in the universe, and containing a fraction $\Omega_b\approx 0.5 \Omega_{BBN}$ (where $\Omega_{BBN}$ is the value of the baryon density predicted by Big Bang Nucleosynthesis). Under these assumptions, BBP obtained a value for $I_{\gamma}$ higher by a factor $\sim 3\div 4$ with respect to our result, based on values $f_g \sim 0.1$ (see Sect.2). Our refined calculations show why their assumption of considering $n_{cl} \sim const$ was reasonable: the clusters effectively contributing to the DGRB are located at $z \simlt 0.2$ (see Fig.5), where the effects of evolution do not have room to take definitely place (see Fig.7). Because of the inherent uncertainties in the predictions of quantities whose calculation involve to set the values of parameters which are not known precisely, we also estimated the range spanned by $I_{\gamma, \nu}$ for the combination of parameters allowed by the present observational ranges. In fact, the description of the cluster structure and evolution that we used in our analytic approach consider only ensemble averaged quantities. But we observe a whole distribution of the real cluster properties with respect to the average cluster moulding. Some amount of variance is needed to be considered in cluster modelling in order to ensure the predictive power of the viable models for structure formation. To explore the role of the uncertainties in the relevant quantities we considered several sources of uncertainties (see Sect.7). From an inspection of Figs. 9 and 10 we note that the effects of the possible theoretical uncertainties in the description of the cluster and IGM evolution could change the predicted contributions for the DGRB and for the DNB by a factor $\simlt 3$, setting the maximal level of $I_{\gamma}$ to a few $\%$ of the EGRET value. The DGRB seems to be mostly produced by AGNs (FSRQ and/or BLLacs) and/or blazars (we take here an estimate of $\sim 60 \div 65 \%$ \cite{erl96} of the EGRET diffuse flux \cite{owz94}). Diffuse $\gamma$-ray emission could be observed also from normal galaxies yielding a contribution $\sim 5 \%$ \cite{erl96}. When added to the $\sim 10 \div 15 \%$ of the DGRB contributed by their high-$E$ photons interacting with other existing backgrounds (\eg IR, CMB \cite{erl96}) one gets only $\sim 15 \div 20 \%$ of the DGRB left for truly diffuse or extended sources. Of this amount, a fraction of the diffuse $\gamma$-ray flux $\sim 3 \div 5 \%$ is predicted \cite{ww} to originate from decaying topological defects (see \cite{sigl}) and interactions of UHE particles with the CMB. Note, however, that the amount and the spectral distribution of this possible diffuse background depend sensitively on the amplitude of the primordial magnetic field on scales larger than supercluster sizes. So, according to the previous estimates, the presence of all these sources of diffuse $\gamma$-ray emission (even though partially model dependent) determines an upper limit to the contribution of extended extragalactic sources to the DGRB, that is $\sim 10 \div 22 \%$ of the OWZ EGRET level \cite{owz94}. This sets rather weak constraints on the level of CR production in clusters and hence on the presence and activity of AGNs in clusters or on the formation and efficiency of accretion shocks around clusters. However, if we consider the revised level of the DGRB as derived by SWZ \cite{swz}, then the previous upper limit reduces to $\simlt 2 \%$. Our predictions of the DGRB contributed by galaxy clusters $I_{\gamma,cl} / I_{EGRET} \sim 0.005 \div 0.02$ is perfectly compatible with the presence of both a population of evolving FSRQ and AGNs dominating the $\gamma$-ray sky and with the presence of truly diffuse backgrounds like those previously discussed. Note, however, that the major source of uncertainty in the level of the extragalactic DGRB comes from the contribution of the AGNs. A fluctuation analysis of the EGRET data is needed to have more definite indications on the level of the DGRB contributed by discrete sources. If, on the other hand, CR acceleration will be found to be relevant in clusters (yielding $L_{CR}$ substantially larger than $ 10^{44}$ erg/s), then the predicted level of DGRB produced by galaxy clusters can set interesting limits to the space density and evolution of $\gamma$-ray AGNs. The sensitivities and angular resolutions achievable by the next generation gamma-ray (INTEGRAL, GLAST, AMS) and neutrino (see \cite{halz} for a review) detectors will be able to shed a new light on the high energy phenomena occurring in large scale structures. \vskip 1.5truecm \newline {\bf Acknowledgements} We thank the Referee for useful comments and suggestions which improved substantially the presentation of the paper. We also aknowledge interesting and stimulating discussions with V.S. Berezinsky during a recent visit of S.C. at the LNGS. S.C. aknowledges also interesting discussions with G. Kanbach, A. Dar and F. Halzen, among others, at the 1997 Moriond Meeting {\it High Energy Phenomena in the Universe}. The research of P.B. is funded by a INFN PostDoctoral Fellowship at the University of Chicago. \newpage | 98 | 4 | astro-ph9804262_arXiv.txt |
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9804 | astro-ph9804054_arXiv.txt | We present the results of an {\asca} observation of the LINER NGC~4579. A point-like X-ray source is detected at the nucleus with a 2--10~keV luminosity of $1.5\times10^{41}$~{\eps}assuming a distance of 16.8~Mpc. The X-ray spectrum is represented by a combination of a power-law with a photon index of $\sim$1.7 and soft thermal component with $kT\sim$0.9~keV. An iron K emission line is detected at $6.73\pm0.13$~keV (rest frame) with an equivalent width of $490^{+180}_{-190}$~eV and is statistically significant at more than 99.9\% confidence. The line center energy is consistent with Helium-like iron and is significantly higher than 6.4~keV which is expected from fluorescence by "cold" (or a lower ionization state of) iron. The iron line profile shows no significant red tail in contrast to Seyfert 1 galaxies although the statistics are limited. The line center energy, equivalent width, and profile are consistent with an origin in an ionized accretion disk. However the large mass accretion rate necessary to ionize the accretion disk is not consistent with the observed luminosity and normal accretion models. | Recent optical spectroscopic surveys have shown that there are many active galactic nuclei (AGNs) in nearby galaxies and about 40\% of bright galaxies are classified as Seyfert galaxies or LINERs (Low Ionization Nuclear Emission-line Regions; \cite{hec80}) (Ho, Filippenko, \& Sargent 1997a). The luminosity of these objects are rather low compared to previously known AGN with a median value of the H$\alpha$ luminosity being only $2\times10^{39}$ {\eps} in the sample of Ho et al (1997a). Such objects (low luminosity AGNs; hereafter LLAGNs) are important for investigating the physics of AGN under an extreme condition, i.e. very low luminosity. X-ray observations probe the innermost regions of AGNs and specifically the iron K line provides information on the ionization state, density, and motion of matter very close to the central energy source. {\asca} observations of Seyfert 1 galaxies revealed that as a class these objects have a broad iron K line with a profile skewed to lower energies, thought to be caused by the reprocessing of the continuum by a relativistic accretion disk (e.g., \cite{tan95}, Nandra et al. 1997a). The center energy of the iron line from Seyfert 1 galaxies is consistent with 6.4~keV, which is expected from fluorescence by neutral or lower ionization states of ($<$\ion{Fe}{16}) iron in a disk with an inclination of $<$ 30 degrees. In the Seyfert 1.9 galaxy IRAS ~18325--5926, a higher peak energy of iron emission is seen, which is compatible with a highly-inclined disk ($i=40-50^{\circ}$) origin (\cite{iwa96}). Highly-ionized iron emission lines are detected from several radio-quiet quasars e.g. E1821+643 (\cite{kii91}, \cite{yam97}) and PG~1116+215 (\cite{nan96}). Nandra et al. (1997b) studied the luminosity dependence of the iron line profile in a large sample of AGN and found that the center energy increases and the red-tail becomes weaker with increasing luminosity. They attributed such behavior to an increasing ionization of the accretion disk with increasing luminosity. Thus X-ray measurements of iron emission lines are powerful diagnostic tools of matter in the vicinity of the nucleus. There are only a few observations of iron emission lines from low luminosity AGNs ( {\LX} (2--10~keV) $\sim 10^{40}-10^{41}$~{\eps}). M81 (NGC~3031) with an X-ray luminosity of {\LX} (2--10~keV) $\sim 2\times10^{40}$~{\eps} shows a broad iron line centered at $\sim6.7$~keV with an equivalent width of $\sim 200$~eV. This line center energy is significantly higher than Seyfert 1 galaxies and similar to luminous quasars. An iron line at 6.4~keV with an equivalent width of $\sim300$~eV is detected from the low luminosity Seyfert 1 galaxy NGC~5033 ({\LX} (2--10~keV) = $2\times10^{41}$~{\eps}, \cite{te98b}) but only an upper limit on the equivalent width of {\simlt} 300~eV is obtained for NGC~1097 ({\LX} (2--10~keV) = $1\times10^{41}$~{\eps}, Iyomoto et al. 1996). Although strong iron emission lines are also detected from M51 (= NGC~5194, Terashima et al. 1998a), NGC~1365 and NGC~1386 (Iyomoto et al. 1997), the iron lines in these objects are interpreted as being caused by reprocessed emission from an obscuring tori and/or extended ionized scatterer outside of our line of sight, i.e., these nuclei are heavily obscured. Thus, at present, the number of LLAGNs with small intrinsic absorption from which iron lines are detected is rather limited. NGC~4579 (M58) is a Sab galaxy in the Virgo cluster of galaxies and classified as a LINER or Seyfert 1.9 galaxy based on the optical emission lines (\cite{ho97a}, \cite{kee83}, \cite{sta82}) and the broad H$\alpha$ component, detected with a FWHM $\sim 2300$~km s$^{-1}$ (\cite{ho97b}). There exists a flat-spectrum radio core (\cite{hum87}). An {\Einstein} HRI observation showed the presence of an unresolved X-ray nucleus and the X-ray flux was measured to be {\FX} = $7.9\times10^{-12}$~{\eps} cm$^{-1}$ in the 0.2--4.0~keV band with the {\Einstein} IPC (\cite{fab92}, Halpern \& Steiner 1983) which corresponds to the X-ray luminosity of $2.7\times10^{41}$~{\eps} (we assume a distance of 16.8~Mpc (\cite{tul88}) throughout this paper). These facts indicate the presence of a LLAGN in this galaxy. A recent ultraviolet imaging observation by {\it Hubble Space Telescope} ({\HST}) Faint Object Camera (FOC) detected a point source at the nucleus (\cite{mao95}). Its UV spectra were taken by {\HST} Faint Object Spectrograph (FOS) and a featureless UV continuum is detected as well as various emission lines. Comparison of the FOC and FOS data also indicate a factor of 3.3 decrease of UV flux in 19 months. The narrow UV emission lines are incompatible with shock excitation model and a photoionization model is preferred (Barth et al. 1996). Several broad UV emission lines are also detected. These UV results provide further support for the presence of a LLAGN in NGC~4579. On the other hand, \cite{mao98} estimated the ionizing photon number by extrapolating the UV luminosity at 1300 A towards higher energies and argued that the observed UV continuum is not sufficient to explain the H$\alpha$ luminosity. They also suggest that emission from AGNs is most prominent at energies higher than the UV. Measurements of an X-ray flux and continuum slope provide information on the ionization source in this LINER. In this paper we report the detection of an Iron K emission line centered at 6.7~keV and discuss X-ray properties of the LLAGN in NGC~4579 and origin of the iron emission line. | \subsection{X-ray emission from a low luminosity AGN} We obtained X-ray images and spectra in the 0.5--10~keV band and a point-like X-ray source with a photon index of $\Gamma = 1.72\pm0.05$ is detected. An iron line is also detected at 6.7~keV. In the soft energy band, a broad line like feature identified with iron-L line complex indicates the presence of thin-thermal plasmas of temperature $kT\sim 0.9$~keV. The X-ray luminosity ($1.5\times 10^{41}~${\eps} in 2--10~keV ) is 1--3 orders of magnitude smaller than typical Seyfert galaxies and falls in the classes of LINERs and "low luminosity" Seyfert galaxies (Serlemitsos et al. 1996, Iyomoto et al. 1996, Ishisaki et al. 1996, Terashima et al. 1998b). In normal spiral galaxies, the X-ray emission is dominated by discrete sources, specifically low mass X-ray binaries (LMXBs) (e.g. \cite{fa89}, Makishima et al. 1989). The X-ray luminosity from LMXBs are roughly proportional to B-band luminosity {\LB} and their X-ray spectra can be approximated by a thermal bremsstrahlung of a temperature of several keV. The {\asca} X-ray spectrum of NGC~4579 is also fitted by $kT\sim8$~keV thermal plasma model. However the strong iron line at 6.7~keV is not compatible with the X-ray spectra of LMXBs, since the equivalent width of iron emission lines from LMXBs are small (several tens of eV, \cite{hir87}). Additionally, the {\LX}/{\LB} value $1.3\times10^{-3}$ is more than an order of magnitude higher than normal spiral galaxies (e.g. \LX /\LB =$3.5\times10^{-5}$ for M31; \cite{mak89}). Additionally an upper limit on the size of an archival {\rosat} PSPC image is 14" (Gaussian $\sigma$), which corresponds to 1.1 kpc at 16.8 Mpc. This upper limit is significantly lower than the size of the galaxy. Therefore we conclude that contribution from LMXBs to X-ray emission of NGC~4579 is negligible. Hot plasmas with temperatures on the order of $\sim10$ keV are present in the Galactic center region and their X-ray spectra show prominent, ionized iron K emission (e.g. Koyama et al. 1996). The X-ray spectral shape of NGC 4579 in the hard X-ray band is similar to such hot gas. However, the X-ray luminosity of NGC 4579 is three orders of magnitude higher than the Galactic ridge emission ({\LX}$\sim2\times10^{38}$ {\eps}; Kaneda et al. 1997, Yamasaki 1996, Warwick et al. 1985). Starburst galaxies also show a hard spectral component with a temperature of $\sim$10 keV and their X-ray luminosities are around $10^{40}$ {\eps} (e.g. $3.4\times10^{40}$ \eps in 2--10 keV for M82; Ptak et al. 1997). However the starburst activity in NGC 4579 is weaker than M82, since the far-infrared luminosity of NGC 4579 is about an order of magnitude lower than that of M82. Furthermore, starburst galaxies show weak or no iron emission contrary to NGC 4579. Therefore hot plasma is unlikely as the origin of the hard component and iron emission line in NGC 4579 and we conclude that the AGN emission dominates the {\it ASCA} spectra and that other components such as a hot gas contribution is small, if any. We note that errors in background subtraction of the Virgo cluster hot gas do not affect the detection of the iron emission line at 6.7 keV, since the cluster gas is very dim in this region and temperature is low ($kT\sim2$ keV; Matsumoto 1998, \cite{boh94}). Actually no significant iron emission is detected from the GIS field around NGC 4579. If the primary ionizing mechanism of LINER optical emission lines in this galaxy is photoionization by a LLAGN, {\LX}/{\LHa} might be expected to be similar to Seyfert 1 galaxies, for which there is a good positive correlation between {\LX} and {\LHa} (e.g. Ward et al. 1988, \cite{kor95}, Serlemitsos et al. 1996). Using the H$\alpha$ luminosity of broad plus narrow component {\LHa} = $5.9\times10^{39}$~{\eps} (\cite{ho97b}) and the observed X-ray luminosity in the 2--10~keV band, we obtain {\LX}/{\LHa} $\approx$ 26 for NGC~4579. This value is in excellent agreement with those of Seyfert 1 galaxies (\cite{war88}) and strongly supports a low luminosity AGN as the ionizing source of the LINER in NGC~4579. Less luminous Seyfert 1 galaxies tend to show rapid and large amplitude variability (Mushotzky, Done, \& Pounds 1993 and references therein). However NGC 4579 shows no significant short term variability. Lack of variability on short time scales seems to be a common property of LLAGNs (\cite{mus92}, \cite{pet93}), for example the LLAGN in NGC 1097 (Iyomoto et al. 1996) and NGC 3998 (\cite{awa92}) also show no significant variability on timescales less than a day. Direct comparison of {\rosat} PSPC and {\asca} flux in the 0.5--2 keV band show a factor of two increase in $\sim$3.5 years. The X-ray spectral slope $\Gamma = 1.72\pm0.05$ is identical to the average value found for hard X-ray selected Seyfert 1 galaxies (Mushotzky et al. 1993) but the luminosity is lower than that of any Seyfert 1 galaxy but NGC 4051. Based on the FW0I (full width at 0 intensity) of a broad emission line and an estimate of the size of the broad line region, mass of the central black hole is roughly estimated to be $M_{\bullet} \sim 4\times10^6M_{\odot}$. Then the Eddington ratio $L/L_{\rm Edd}$ is $\sim 10^{-3}$ for the observed luminosity of $\sim5\times10^{41}$~{\eps} (Barth et al. 1996), although their blackhole mass estimation is crude. Therefore the X-ray spectral slope does not seem to be drastically changed even at a very low Eddington ratio. This is also true for M81, for which $L/L_{\rm Edd}$ is estimated to be $\sim (2-10)\times 10^{-4}$ (\cite{ho96}) and the photon index is $1.85\pm0.04$ (\cite{ish96}). Soft thermal emission of $kT\sim0.5-1$~keV is often observed from low luminosity AGNs (Terashima 1997, \cite{pta97}, Serlemitsos et al. 1996). In some cases, such emission is associated with starburst activity (e.g. \cite{iyo96}, \cite{te98a}). Since the far-infrared luminosity of NGC~4579 is $1.5\times10^{43}$~{\eps} some star formation activity may be present which may explain the thermal emission. The soft thermal X-ray to far infrared luminosity ratio {\LX}/{\LFIR} = $6\times10^{-4} - 1.1\times10^{-3}$ is consistent with starburst galaxies (e.g. \cite{dav92}) within the scatter. \subsection{Iron-K line} A marginally broad ($\sigma \approx 0.17$~keV) iron emission line is clearly detected at $6.73^{+0.13}_{-0.12}$~keV and the equivalent width is $490^{+180}_{-190}$~eV for the broad Gaussian model fit. The line center energy is significantly higher than 6.4~keV, which is typically observed from Seyfert 1 galaxies, and consistent with He-like iron. A similar broad iron line centered at $\sim 6.7$~keV is detected from the low luminosity Seyfert galaxy M81 (Ishisaki et al. 1996, Serlemitsos et al. 1996). The line can also be represented by line blending of neutral, He-like, and H-like iron and dominated by He-like iron. The disk-line profile (Fabian et al. 1989) is probably inconsistent with the data for 6.4~keV or 6.7~keV intrinsic line energy because of following reasons. The $\chi^2$ value is worse than a single broad Gaussian fit and systematic residuals remain in the disk-line fit, since a significant red tail is not clearly seen in the data. Furthermore, the disk-line model provides the very large equivalent width $\sim900$~eV, which is about 4 times larger than the results of the disk-line fit to Seyfert 1 galaxies ($<$EW$>=(230\pm60)$ ~eV, \cite{na97a}). Therefore our data prefer a symmetric Gaussian-shape profile with intrinsic line center energy of 6.7~keV (He-like) rather than 6.4~keV ($<$\ion{Fe}{16}). Thus the ionization state of the iron line emitter may be different from that of higher luminosity Seyfert 1 galaxies in at least some LLAGNs (NGC 4579 and M81) Strong ionized iron emission lines are observed in heavily obscured Seyfert 2 galaxies (NGC 1068, \cite{uen94}, \cite{iwa97}; NGC 1365, \cite{iyo97}; see also \cite{tu97a}, \cite{tu97b}). In these objects continuum emission from the nucleus is completely blocked and only scattered radiation is observed. Ionized iron lines are interpreted as originating from a photoionized scattering medium. If the continuum of NGC~4579 is scattered radiation, then the observed X-ray luminosity is only a fraction of its intrinsic luminosity. Since the scattering fraction is typically less than 10 \% for Seyfert 2 galaxies (\cite{uen95}), the {\LX}/{\LOIII} should be less than 10 \% of those of Seyfert 1 galaxies as is the case for NGC 1068 (\cite{mul94}). However the observed X-ray to [OIII]$\lambda$5007 luminosity ratio {\LX}/{\LOIII} is very similar to Seyfert 1 galaxies. Therefore the observed X-ray continuum is not likely to be due to a scattered component. Then the observed iron line should be emitted from the matter close to the nucleus in order to be ionized and/or broadened due to the Doppler effect. If the iron line is emitted by an accretion disk, a line profile with significant red tail is expected (\cite{fab89}). On the other hand, the observed profile seems to be symmetric in shape although the statistics are limited. Broad lines with weaker red tails than Seyfert 1s are observed in AGNs with much higher luminosity; {\LX}$>10^{44}$~{\eps} (\cite{na97b}). If the inner-most part of the accretion disk is almost fully ionized, the red component is expected to be very weak or absent. Thus the observed profile is consistent with the interpretation that the observed iron K emission is from an ionized disk. The obtained equivalent width ($\sim 500$~eV for the Gaussian model) is rather large compared to that seen in most Seyfert 1 galaxies. If the disk is highly ionized, the fluorescence yield of iron increases and absorption by lighter elements decreases as light elements are almost completely ionized. In such a situation the equivalent width of an iron line can increase by a factor of two (Matt et al. 1993, \cite{zyc94}). Therefore the large equivalent width is also naturally explained by an ionization effect. The ionization state of photoionized matter is determined by an ionization parameter $\xi = L/nR^2$ (\cite{kal82}), where $L$, $n$, and $R$ is the luminosity of ionizing photons, the number density of photoionized matter, and the distance from light source to photoionized matter, respectively. The X-ray luminosity of NGC 4579 is only $1.5\times10^{41}$ {\eps}, which is 1--3 orders of magnitude smaller than for Seyfert 1 galaxies, and the X-ray luminosity of M81, from which an iron line centered at $\sim6.7$~keV is detected, is even lower ($\sim 2\times10^{40}$ {\eps}). In order to photoionize iron atoms to be He-like, $\xi$ should be at least $\sim$ 500, while $\xi<100$ is required for less ionized species ($<$ \ion{Fe}{16}) which is probably appropriate for usual Seyfert 1 galaxies. Therefore $nR^2$ in the iron line emitting region should be more than two orders of magnitude smaller than that of luminous Seyfert 1 galaxies. An expected ionization parameter under an assumption of standard $\alpha$ disk is calculated by \cite{mat93}. According to their results, the ionization parameter has a strong dependence on the mass accretion rate $\xi \propto \dot{m}^3$ (equations (5) and (6) in Matt et al. 1993), where $\dot{m}$ is denoted in units of the critical accretion rate $\dot{m} = L/L_{\rm Edd}$. In order to ionize iron to He-like, $\dot{m}$ should be at least 0.2 (Figs. 2 and 5 in Matt et al. 1993). However the order-of-magnitude estimate of the central black hole mass by \cite{bar96} combined with the observed luminosity gives a significantly smaller value of $\dot{m}$ $\sim1\times10^{-3}$. Then we cannot explain the very low luminosity and the ionized iron line at the same time in the standard disk model. This may suggest that the accretion processes in AGN is different in a very low luminosity situations with very small $\dot{m}$. An advection dominated accretion flow (ADAF) model is proposed for AGNs specifically for objects radiating at very low Eddington ratio (e.g. $\dot{m} \sim 10^{-4}$ for NGC 4258, Lasota et al. 1996). In the model by Lasota et al. (1996), a standard disk is assumed outside of $r_{\rm in}$ and an ADAF inside of $r_{\rm in}$. In an ADAF, accreting matter is heated up to very high temperatures ($T_i\sim10^{12}$K, $T_e\sim10^{9}$K). However our detection of an iron line indicates the presence of highly ionized (but not fully ionized) matter surrounding a large solid angle viewed from the light source. This means that $r_{\rm in}$ should be small and a geometrically thin disk is appropriate. Therefore the iron line in NGC~4579 cannot be explained solely by an ADAF model and the real situation in NGC~4579 may correspond to a condition near the transition from the $\alpha$ disk to an ADAF. Future sophisticated modeling of accretion in LLAGNs and calculation of expected iron emission as well as precise measurements of an iron K line and mass determination by {\HST} Space Telescope Imaging Spectrograph will be important to understand physical processes in extremely low luminosity AGNs.\\ | 98 | 4 | astro-ph9804054_arXiv.txt |
9804 | gr-qc9804048_arXiv.txt | We conduct a direct comparison of three different representative numerical codes for constructing models of rapidly rotating neutron stars in general relativity. Our aim is to evaluate the accuracy of the codes and to investigate how the accuracy is affected by the choice of interpolation, domain of integration and equation of state. In all three codes, the same physical parameters, equations of state and interpolation method are used. We construct 25 selected models for polytropic equations of state and 22 models with realistic neutron star matter equations of state. The three codes agree well with each other (typical agreement is better than 0.1 \% to 0.01 \%) for most models, except for the extreme assumption of uniform density stars. We conclude that the codes can be used for the construction of highly accurate initial data configurations for polytropes of index $N>0.5$ (which typically correspond to realistic neutron stars), when the domain of integration includes all space and for realistic equations with no phase transitions. With the exception of the uniform density case, the obtained values of physical parameters for the models considered in this paper can be regarded as ``standard'' and we display them in detail for all models. | The physical state of the neutron star matter has not been fully understood yet because it is very difficult to investigate particle interactions beyond nuclear matter density ($\varepsilon_{\rm N}/c^2 \sim 2 \times 10^{14}$ g cm$^{-3}$) either from nuclear experiments or from nuclear theories, (here $\varepsilon_{\rm N}$ is the energy density of the nuclear matter and $c$ is the velocity of light). Given this situation, one promising approach to explore the behavior of very high density matter is to make use of the macroscopic quantities of neutron stars. In particular, the mass and the rotational period of neutron stars depend crucially on the softness of the equation of state (EOS) at very high densities (see e.g. Friedman et al. 1984, 1986, 1989), thus, observational constrains, matched with theoretical models, may help in reconstructing the equation of state of very high density matter. Given a particular equation of state, the mass of neutron stars varies with central energy density and always reaches a maximum. This implies that if the maximum mass of neutron star models constructed with a certain equation of state is smaller than the mass of observed neutron stars, that equation of state must be discarded. Currently, the largest accurately measured mass of slowly rotating neutron stars is $M_{\rm BP} = 1.44 M_{\odot}$, where $M_{\rm BP}$ is the mass of one of the components of the binary pulsar PSR1913+16 (Taylor \& Weisberg \cite{tayl89}) and $M_{\odot}$ is one solar mass. Individual masses of neutron stars have also been estimated in six other binary pulsars (Thorsett et al. \cite{thor93}, Wolszczan \cite{wol97}), as well as in six X-ray binaries (van Kerkwijk et al. \cite{vankerk95}) but the accuracy is not as good as in PSR1913+16. Thus, equations of state which give larger masses than $M_{\rm BP}$ for slowly rotating stars, can be valid as candidates for the real equation of state at very high densities. Since the maximum mass of neutron stars is smaller for more compressible (soft) equations of state than for less compressible (stiff) equations of state, the true equation of state at high densities cannot be extremely soft. On the other hand, stiff equations of state can be limited by considering the neutron star with the shortest rotational period, i.e. the most rapidly rotating pulsar. There exists a lower limit on the rotational period for each equation of state, because if the centrifugal force exceeds the self-gravity at the equatorial surface, no equilibrium states are allowed. The lower limit of the rotational period depends on the softness of the equation of state - the radius of neutron stars with softer equations of state is smaller, which allows for higher rotation rates. Thus, if very rapidly rotating neutron stars should be found, we could exclude most stiff equations of state. At the moment, the shortest period of observed pulsars is 1.56 ms, of PSR1937+21. Consequently, equations of state for which the shortest rotational period is larger than this value, must be excluded as candidates for the real equation of state for neutron star matter. The discussions above require us to make use of highly accurate schemes for constructing rotating neutron star models, in order to compute precise theoretical values of masses and rotational periods. Highly accurate relativistic equilibrium models are also needed as initial data for relativistic time-evolution codes (modeling of nonlinear pulsations, collapse and generation of gravitational waves). Recently, a number of groups have succeeded in constructing models of rapidly rotating neutron stars (Friedman et al. 1984, 1986, 1988, 1989, Eriguchi et al. \cite{erig94}, Salgado et al. \cite{salg94}, \cite{salg94b}, Cook et al. \cite{cook94b}, Stergioulas \& Friedman \cite{stag95} -- for a recent review see Stergioulas \cite{S98}). However, the obtained models by those authors do not always agree with each other (see e.g. Friedman et al. \cite{frie89}, Eriguchi et al. \cite{erig94}, Salgado et al. \cite{salg94}, Cook et al. \cite{cook94b}, Stergioulas \& Friedman \cite{stag95}). Although Stergioulas \& Friedman~(1995) have determined the cause of the discrepancy between models in Friedman et al.~(1989) and Eriguchi et al.~(1994), (which was due to the use of a slightly different equation of state table), the reasons of smaller differences which remain, even after using exactly the same equation of state, have not been clarified yet. This is because numerical techniques used in the different codes, such as the choice of parameters defining the model, the interpolation method, the method of integrating the field equations, a.s.o. are not the same. In this paper, three groups using their own codes (Komatsu et al. \cite{koma89a}, Eriguchi et al. \cite{erig94}, Salgado et al. \cite{salg94}, Stergioulas \& Friedman \cite{stag95}) will decrease the differences between their results to a minimum possible, by tuning each code and using the same parameters, the same schemes of interpolation, the same equations of state, and so on. Since the basic schemes used by the three groups are different, it will be impossible to have exactly the same results and the relative differences between results are a measure of the accuracy of the codes. Models obtained with small relative differences between the three codes can be considered as ``standard" models for each equation of state. Furthermore, this direct comparison allows us to investigate the effect that the choice of interpolation method, equation of state and domain of integration has on the accuracy of the codes. | \subsection{Discussion} \subsubsection{Metric Potentials} As redshift factors differ by about 10 \% for the constant density, relativistic model N00$rr$ between the three codes (while the agreement global quantities is within a few \%) we compare directly the local values of metric potentials for several models. Figures 1 to 16 show the four metric potentials (upper panel) and the relative differences in them between BGSM and KEH(OR) (middle panel) and KEH(SF) and BGSM (lower panel) for the models N05$mr$, N15$mr$, L(L)$mr$ and WFF(FPS)$br$. The metric potentials are graphed against the coordinate $r$ in the equatorial plane ($\theta= \pi/2$, solid line) and along the axis of rotation ($\theta=0$, dashed line). The range of the coordinate $r$ is the twice the equatorial radius of the star. In general, the agreement in the local values of the metric potentials reflects the agreement in the computed physical parameters of models. In these graphs, several significant behaviors can be pointed out: First, there are high frequency and small amplitude oscillations at the inner part of the stars for all models. Second, the differences between the results of KEH(OR) and those of the other two codes are growing outside the stars as $r$ increases. Third, although the differences between the KEH(SF) and BGSM codes are very small for models N05$mr$, N15$mr$ and WFF(FPS)$br$, there appear larger differences for the stiff model L(L)$mr$. Fourth, there appears a larger amplitude oscillation in the metric potential $\omega$ on the axis of rotation, close to the surface. The first behavior is due to the integration scheme of the KEH code, i.e. the Simpson scheme. In general, the Simpson scheme gives results with higher precision, compared with those obtained by the trapezoidal scheme. However, in the KEH scheme, the integrands contain nonsmooth functions with respect to the radial coordinate, because of the nature of the Green's functions. The non-uniform distribution of the weight factor in Simpson's scheme for nonsmooth functions results in oscillating behaviors with very small amplitudes, which cannot be noticed in the behavior the integrated quantities. The second behavior in the original KEH code is caused by the "truncation" of the domain of integration at a finite distance from the star, instead of integrating over the whole space. The large differences in the metric potentials between KEH(SF) and BGSM for EOS L, could be accounted to the stiffness of the equation of state, but the differences between KEH(OR) and BGSM for the same model are not as large, and we have not an explanation for that. The oscillations in $\omega$ on the axis of rotation near the surface are present also for the soft $N=1.5$ polytropes, while for $N=0.5$ they are larger. This indicates that terms in the field equations for $\omega$ are very sensitive to the presence of the surface and the accompanying Gibbs phenomenon. Even for $N=1.5$ polytropes, where the density goes to zero smoothly at the surface, there is a small scale Gibbs phenomenon, due to the finite number of grid points used to represent the region of integration. \subsubsection{Method of Interpolation} An important factor for the local accuracy of models constructed with realistic equations of state is the method of interpolation of the energy vs. pressure data given in an EOS table. While global quantities are not affected significantly, the virial identities for realistic EOSs, are sensitive to the interpolation scheme This can be considered to reflect the nature of the interpolation scheme as mentioned before. If we define the enthalpy ($H$) by \begin{equation} H \equiv \ln \left( { \varepsilon + p \over \rho c^2} \right), \end{equation} the Gibbs-Duhem relation, which follows directly from the first law of thermodynamics, implies \begin{equation} {dp\over dH} = \varepsilon + p \ . \end{equation} In the cubic Hermite interpolation, the Gibbs-Duhem relation is used to replace by $\nabla H$ the term $\nabla p /(\varepsilon+p)$ which appears in the hydrostationary equilibrium equation. If the tabulated function $p(H)$ fails to satisfy the above relation, then the hydrostationary equilibrium equation, which is derived from the Bianchi identity, is only approximately verified by the numerical solution, which results in increased error in the GRV2 and GRV3 virial identities. The four point Lagrange interpolation does not satisfy the Gibbs-Duhem relation because it only reproduces the values of the discrete points, but there is no guarantee for the reproduction of the derivatives. This explains why the GRV2 and GRV3 errors are bad, even in the nonrotating case (GRV2 = 3E-03, GRV3 = 1E-02 for model L$sr$) as compared to ${\rm GRV2}\sim 10^{-14}$ for polytropic models (see e.g. Bonazzola et al. \cite{bona93}). The GRV2 and GRV3 error indicators thus do not reflect the precision of the code but the bad thermodynamical behavior of the tabulated EOS. The advantage of the cubic Hermite interpolation is that the Gibbs-Duhem relation is automatically satisfied because this interpolation reproduces not only the values themselves but also the derivatives (Swesty 1996). Moreover, in our case, the energy density and the baryon number density are obtained by \begin{eqnarray} \varepsilon & = & {p \over H} {d\log p\over d\log H} - p, \\ n & = & {\varepsilon + p \over m_{\rm B} c^2} \exp(-H) \ . \end{eqnarray} Because of these equations, the Gibbs-Duhem relation is satisfied in every point. Note also that the value of $\varepsilon$ obtained in this way coincides exactly with $\varepsilon_i$ at the points in the tabulated equation of state. \subsection{Conclusion} The comparison of three different codes for constructing rapidly rotating relativistic neutron star models demonstrates that the BGSM and KEH schemes used are highly accurate for typical polytropic models - when the field equations are solved to infinity - and for models constructed with realistic equations of state, that do not have phase transitions. If one approximates neutron stars as constant density stars, then Gibbs phenomena at the discontinuous surface reduce the accuracy of the computed models. If high accuracy in such models and in models with phase transitions is desired, then modified numerical schemes - free of Gibbs phenomena - need to be used. Such numerical schemes could employ, for example, surface fitted coordinates. Such a scheme has been presented recently by Bonazzola et al. (1998a) in the framework of spectral methods and looks promising for rotating stellar models. Further, we demonstrated that the metric potentials are subject to various local oscillatory behaviors, even if integrated quantities have very good accuracy. This observation is important for the effort of constructing initial data for the numerical evolution of rotating relativistic neutron star models. \begin{figure} \resizebox{\hsize}{10.5cm}{\includegraphics{zeta_19.eps}} \caption{Same as Figure 1 but for the metric potential $\zeta$ of model N15$mr$.} \end{figure} \begin{figure} \resizebox{\hsize}{10.5cm}{\includegraphics{zeta_27.eps}} \caption{Same as Figure 1 but for the metric potential $\zeta$ of model WFF(FPS)$br$.} \end{figure} \begin{figure} \resizebox{\hsize}{10.5cm}{\includegraphics{zeta_4.eps}} \caption{Same as Figure 1 but for the metric potential $\zeta$ of model N05$mr$.} \end{figure} \begin{figure} \resizebox{\hsize}{10.5cm}{\includegraphics{zeta_26.eps}} \caption{Same as Figure 1 but for the metric potential $\zeta$ of model L(L)$mr$.} \end{figure} | 98 | 4 | gr-qc9804048_arXiv.txt |
9804 | gr-qc9804081_arXiv.txt | In this paper we investigate extended inflation with an exponential potential $V(\sigma)= V_0\,e^{-\kappa\sigma}$, which provides a simple cosmological scenario where the distribution of the constants of Nature is mostly determined by $\kappa$. In particular, we show that this theory predicts a uniform distribution for the Planck mass at the end of inflation, for the entire ensemble of universes that undergo stochastic inflation. Eternal inflation takes place in this scenario for a broad family of initial conditions, all of which lead up to the same value of the Planck mass at the end of inflation. The predicted value of the Planck mass is consistent with the observed value within a comfortable range of values of the parameters involved. | The extended inflation action is given by \cite{extended} \begin{equation} \label{a} S =\int d^{4}x \,\sqrt{-g}\left[\Phi R - {\omega\over\Phi}(\partial \Phi)^{2} - \frac{1}{2}(\partial \sigma)^{2} \\ - V(\sigma)\right] , \end{equation} where $R$ is the curvature scalar and the potential is $V(\sigma) = V_0 e^{-\kappa\sigma}$. The coupling $\omega$ plays a similar r\^{o}le as that of the coupling functions $B_i(\Psi)$ of the dilaton field $\Psi$ in string theory. Based on this analogy, several authors have investigated the so-called {\it hyperextended inflation} models \cite{extended2,extended3,mikel2}, where $\omega$ becomes dependent on the BD field. In this paper however, we will merely examine the $\omega={\rm const}$ model. The BoI boundary is given by $V(\sigma)= M^{4}_{\rm p}(\Phi)$ or equivalently, \begin{equation} \label{c1} \Phi = \frac{V_0^{1/2}}{16\pi}\,e^{-\kappa\sigma/2}. \end{equation} The BoI is the quantum limit where the metric fluctuations become significant and the inflaton field cannot take the values for which the potential is above this boundary. The EoI boundary is on the other hand \begin{equation} \label{d} \frac{1}{2}\dot \sigma^{2} + \omega\frac{\dot \Phi^{2}}{\Phi} \approx V(\sigma). \end{equation} The equations of motion in a flat FRW background are \begin{equation} \label{e1} \left(D^2+\frac{1}{2\omega}R\right)\,\Phi=0\,, \end{equation} \begin{equation} \label{e2} D^{2}\sigma = -V^{\prime}(\sigma)\,, \end{equation} \begin{equation} \label{e3} H^{2}+H\frac{\dot \Phi}{\Phi} = \frac{\omega}{6} \left({\dot\Phi\over\Phi}\right)^2 + \frac{1}{6\Phi} \left[\frac{1}{2}\dot\sigma^{2} + V(\sigma)\right] , \end{equation} and the differential operator $D$ is defined \begin{equation} \label{f} D^{2} \equiv \partial^{2}_{t} + 3H\partial_{t} . \end{equation} In the slow-roll approximation, $\ddot \Phi \ll H\dot\Phi \ll H^{2}\Phi$ and $\dot \sigma^{2}+2\omega\, \dot\Phi^{2}/\Phi \ll 2V(\sigma)$, (\ref{e1})-(\ref{e3}) read \begin{eqnarray} \label{g1} \frac{\dot\Phi}{\Phi} &=&2\frac{H}{\omega} \,, \\ \dot \sigma &=& -\frac{1}{3H}V^{\prime}(\sigma) \,,\\ H^{2} &=&\frac{1}{6\Phi}V , \end{eqnarray} and the curvature scalar is given by $R=-12H^{2}$. The slow-roll equations (\ref{g1}) enable us to rewrite (\ref{d}): \begin{equation} \label{h} \Phi_* =\left(3-\frac{2}{\omega}\right)\frac{1}{\kappa^{2}}, \end{equation} where the $*$ subindex denotes the value at the end of inflation. Hence, $\Phi_*$ is independent of $\sigma$ and, for reasonably large $\omega$, it is solely determined by the slope of the potential. The condition $\Phi>0$ also imposes the constraint $\omega$, as can be seen in Fig.~1, such that the range $0<\omega<2/3$ is excluded to prevent imaginary values of the Planck mass. The classical trajectories of the fields are given by the following conservation law \cite{mikel2} \begin{equation} \label{h1} \frac{d}{dt}\left[\omega\Phi + \int d\sigma \frac{V(\sigma)}{V^{\prime}(\sigma)} \right]=0, \end{equation} which in the case of the exponential potential yields \begin{equation} \label{h2} \Phi= \frac{\sigma}{\kappa \omega} + \left(\Phi_0- \frac{\sigma_0}{\kappa \omega}\right). \end{equation} In Fig.~2 we have put together the BoI and EoI curves and the classical trajectories of the fields on the ($\sigma$,$\Phi$) plane, i.e. (\ref{c1}),(\ref{h}) and (\ref{h2}) respectively. It can be seen in the figure that inflation takes place in the region enclosed by BoI and EoI to the right of the intersection point $A$. The trajectories given by (\ref{h2}) are straight lines parallel to the segment $BC$, and the fields drift along these curves in the direction $B\to C$ during the course of inflation. The region enclosed by BoI and EoI to the left of $A$ does not undergo inflation, because the orientation of the classical trajectories is such that the fields would move from EoI towards BoI, which is not an acceptable solution. In addition to the classical trajectories quantum diffusion is responsible for the jumps of the fields between neighbouring classical trajectories. It can be seen that, unlike with powerlaw potentials, for which $\sigma$ decreases as $\Phi$ increases during the course of inflation, in the case of the exponential potential both fields increase during the slow roll. \begin{figure}[t] \centering \leavevmode\epsfysize=5.5cm \epsfbox{omega.ps}\\ \vskip 0.2cm \caption[fig1]{BD field at EoI, $\Phi_*$, vs. $\omega$ for an arbitrary value of $\kappa$. $\Phi_*$ is given in units of $\kappa^{-2}$. Inflation takes place in the range $\omega<0$ or $\omega\geq 2/3$.} \end{figure} \begin{figure}[t] \centering \leavevmode\epsfysize=5.5cm \epsfbox{fields.ps}\\ \vskip 0.2cm \caption[fig2]{Predicted BoI and EoI curves, dashed and thick solid curves respectively. Classical trajectories are straight lines parallel to $BC$. At the intersection point $A$ the onset and end of inflation coincide. $\sigma_{\rm max}$ determines the scale of validity of the slow-roll approximation. Inflation takes place within the region enclosed by BoI, EoI and $\sigma\approx\sigma_{\rm max}$.} \end{figure} The EoI boundary (\ref{h}) gives a definite and unique prediction for $\Phi_*$, and also it implies that if $\Phi_0>\Phi_*$ inflation will not occur. In the case of $0<\Phi_0\lsim\Phi_*$, inflation takes place for values of the inflaton \begin{equation} \label{sigma0} \sigma_0 \gsim -\frac{2}{\kappa}\, {\rm log}\left({16\pi\Phi_0 \over V_0^{1/2}}\right). \end{equation} Naturally if $\Phi_0=\Phi_*$, then the RHS of (\ref{sigma0}) is $\sigma_A$, the value of the field at the intersection point $A$ of BoI and EoI in Fig.~2. It must be noted that the slow-roll approximation does not hold in the case of the exponential potential for arbitrarily large values of $\sigma$. For a given value of $\kappa$ it is straightforward to compute $\sigma_{\rm max}$ for which the potential and kinetic energy of the fields are comparable and thus the slow-roll conditions break down. It is easy to show from (\ref{d}) that this scale is \begin{equation} \label{smax} \sigma_{\rm max}\approx \left({3\omega-2\over\kappa}\right) . \end{equation} Therefore it follows that the EoI boundary does not span from $\sigma_A$ to infinity, but inflation will occur within a finite region $\sigma_A\lsim\sigma\lsim\sigma_{\rm max}$. | In this paper we have examined extended inflation with an exponential potential. The remarkable feature of this model is the prediction of a constant distribution of the Planck mass at the end of inflation, given by (\ref{h}). The parameter $\kappa$ of the theory is therefore estimated via the observed Planck mass in this region of the universe, which in turn fixes the parameter $\sigma_{\rm max}$ that determines the range of values of $\sigma$ for which inflation takes place. The amplitude of the potential $V_0$ is left unconstrained by astrophysical bounds on the spectrum of fluctuations, as described by the argument given in \S 4. The dynamics as is given in \S 2 and the likelihood distributions in \S 3 are shown to be insensitive to the numerical value of this parameter. As is shown in Fig.~2, the BoI and EoI curves in this model cross at $\sigma=\sigma_A$ and the area enclosed between them is thus infinite. However the breakdown of the slow-roll approximation for the exponential potential over the range $\sigma\gsim \sigma_{\rm max}$ (where $\sigma_{\rm max}$ is given by (\ref{smax})) implies that in practical terms only a finite region of the ($\sigma$,$\Phi$) plane undergoes inflation. In the classification of \cite{V} this means that the exponential potential is {\it class I}, i.e. the values of the fields at EoI remain finite. \subsection* | 98 | 4 | gr-qc9804081_arXiv.txt |
9804 | astro-ph9804239_arXiv.txt | This study is the first known attempt to search for gamma-ray burst repeaters combining data from gamma-ray experiments flying on board different satellites and making use of information derived from the bursts detected simultaneously by all the experiments. The proposed method is suitable to correlate GRB data provided by experiments that overlap partially or totally in time. As an application of this method we have correlated the positions of 57 gamma-ray bursts observed by WATCH/GRA\-NAT and WATCH/EURECA with 1905 bursts detected by BATSE. Comparing the so-called~~~``added correlation''~~~between~~~the WATCH and BATSE bursts with that obtained with simulated WATCH catalogues, we conclude that there is no indication of recurrent activity of WATCH bursts in the BATSE sample. We derive an upper limit of $15.8\%$, with a confidence level of $94\%$, for the number of WATCH gamma-ray bursts that could represent a population of repeaters in the BATSE sample. | Despite the advances carried out so far, the origin of the gamma-ray bursts (hereafter GRBs) remains unknown. The identification of absorption lines in the optical spectrum of GRB 970508 strongly supports models arising from sources at cosmological distances (Metzger et al. 1997), but there is still a lack of knowledge on the mechanisms originating these enigmatic phenomena. One of the most important clues that could clarify the nature of the GRBs would be the detection of a repeater behaviour. Initial studies showed an apparent evidence of repetition for the BATSE 1B catalogue (Quashnock and Lamb 1993), suggesting that it would be possible to have an excess of pairs of GRBs clustered in both time and space (Wang and Lingenfelter 1995). This fact was not confirmed by the work carried out using the BATSE 2B catalogue (Brainerd et al. 1995), although other studies provided marginal evidence for both temporal and angular clustering (Petrosian and Efron 1995). Analyses based on autocorrelations with data from the BATSE 3B catalogue did not find any evidence of repetition (Bennett and Rhie 1996) and have imposed several constraints to the number of repeaters (Tegmark et al. 1996). Finally, recent studies confirm the lack of repetition in the 4B catalogue and lead to an upper limit to the repetition rate of $ 0.04$ burst source$^{-1}$ yr$^{-1}$ (Hakkila et al. 1997). The BATSE~4B catalogue was obtained by the BATSE experiment on board the {\it CGRO} satellite and contains 1637 GRBs detected from April 1991 to August 1996 (Paciesas et al. 1998). The BATSE experiment consists of eight identical detector modules, placed at the corners of the {\it CGRO} spacecraft and covering energy channels from $\sim 25$ keV to $\sim$ 2 MeV. It provides error boxes with a minimum radius of $1.6^{\circ}$ (1$\sigma$ confidence level, Fishman et al. 1994). BATSE is detecting bursts at a rate of 0.8 bursts per day. The bursts are daily added to the so-called Current GRB Catalogue, which contains the BATSE~4B catalogue plus all bursts detected after August 1996. When this study was started, the catalogue contained 1905 sources; this sample constitutes the basis of the present work. The WATCH X-ray all-sky monitor is based on the rotation modulation principle (Lund 1986). The instrument has a circular field of view of 4 steradians and an effective area of $\sim$ 30 cm$^2$ (averaged over the field of view). Position sensitivity is achieved using the rotation collimator principle, with the collimator grids rotating with a frequency $\omega$=1 Hz. The phoswich detectors consist of interleaved scintillator-strips of NaI and CsI crystals. The geometric area of the scintillator is 95~cm$^2$. Four units were mounted on board the Soviet {\it GRANAT} satellite in a tetrahedral configuration covering the whole sky, and one unit on board the European Space Agency {\it EURECA} spacecraft. The total energy range is 8-80~keV, therefore overlapping with the lower BATSE energy band. WATCH\-/GRANAT detected bursts in 1990-94 and WATCH/EURECA in 1992-93, thus both experiments also overlapped in time with BATSE. One of the main advantages of WATCH was the capability of locating bursts with relatively small error boxes ($3\sigma$ error radii with $\sim$ 1$^{\circ}$) (Brandt et al. 1990). WATCH\-/GRANAT detected 47 GRBs in this period and WATCH\-/EURECA 12 (Castro-Tirado et al. 1994, Brandt et al. 1994, Sazonov et al. 1998). Two GRBs (GRB 920814 and GRB 921022) were detected by both the WATCH\-/GRANAT and WATCH\-/EURECA experiments. Therefore, the sample of WATCH GRBs used in this study comprises 57 GRBs: 45 WATCH\-/GRANAT bursts, 10 WATCH\-/EURECA bursts and the above-mentioned two GRBs. BATSE also detected 27 of them. Fig.~\ref{figure1} shows the sample of 57 WATCH GRBs used in this study. \begin{figure*} \centering \resizebox{!}{!}{\includegraphics[width=\hsize,totalheight=10.7cm]{7502.f1}} \caption{Error boxes for the 57 GRBs detected by WATCH, represented in galactic coordinates. The sample contains 45 GRBs detected by WATCH/GRANAT, 10 by WATCH/EURECA and two localized by both experiments at the same time. The typical radii of the error boxes are $\sim 1 ^{\circ}$, with a $3 \sigma$ confidence level.} \label{figure1} \end{figure*} The distribution of time amplitudes for GRBs shows two classes of bursts: a) durations shorter than $\sim$ 2 s and b) longer than $\sim$ 2 s (Kouveliotou et al. 1993). It was noticed that the energy spectra of the short bursts were generally harder than those of the long ones (Kouveliotou et al. 1993, Lestrade et al. 1993). The fraction of short events in the WATCH sample is smaller than that in the 4B catalogue. This fact can be justified by at least three selection effects: i) The availability of WATCH for localizing sources is governed by the rotation speed of the collimator grids (1~Hz). So, a source needs to be bright enough for at least one rotation of the modulation collimator in order to be localized, implying a burst duration longer than 1~s. In contrast, the BATSE experiment is able to detect bursts with durations as short as 64 ms. ii) The low energy band of the WATCH experiment ($\sim$8-20~KeV) is sensitive to the soft GRBs, below the BATSE lower limit ($\sim$25 KeV), which generally belong to the class of bursts with durations longer than 2 s. iii) On the other hand, since WATCH is about an order of magnitude less sensitive than the large-area detectors of BATSE, the WATCH catalogue contains bursts which are brighter than those in the BATSE sample. The above three reasons~~explain~~why~~the GRBs~~in the WATCH sample are longer, softer and brighter than the average BATSE 4B bursts. This study is the first known attempt to search for repeaters combining data of $\gamma$-ray experiments flying on board different satellites. The method proposed makes use of the so-called ``simultaneous bursts'' and is suitable to correlate GRB data provided by experiments that overlap partially or totally in time. In the future, this work could also be used to detect systematic pointing errors between different $\gamma$-ray experiments, allowing to improve their capability for locating GRBs. | In this study we have developed a method that allows us to search for GRBs common to two catalogues of sources, each one based on a different instrument. The method makes use of the GRBs detected simultaneously by both experiments, so it is necessary that the experiments overlap in time. We have applied~~the~~method~~to the WATCH (WATCH/GRANAT + WATCH\-/EU\-RE\-CA) and BATSE (BATSE 4B + bursts detected after August 1996) catalogues. We conclude that there is no evidence of recurrent activity of WATCH bursts in the BATSE sample. We claim (with a $94\%$ confidence level) that no more than a $15.8\%$ of the 57 GRBs detected by WATCH are present in the sample of 1905 BATSE bursts (excluding the simultaneous bursts). However, the possibility of finding repeaters in each single catalogue cannot be ruled out. Our results support models which do not predict repetitions of GRBs (for instance the merging of neutron stars at cosmological distances). | 98 | 4 | astro-ph9804239_arXiv.txt |
9804 | astro-ph9804180_arXiv.txt | s{ We address the problem of map-making with data from the \ps\ High Frequency Instrument, with an emphasis on the understanding and modelling of instrumental effects, and in particular that of sidelobe straylight. } | We have shown that the scan strategy of \ps\ along rings on the sky allows to decompose the problem of converting data streams into CMB anisotropy maps in two independent steps. This makes the problem tractable numerically, and helps in analysing and monitoring the impact of systematic effects. In particular, a promising method for the identification and removal of sidelobe signals has been developped. | 98 | 4 | astro-ph9804180_arXiv.txt |
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9804 | astro-ph9804149_arXiv.txt | We review recent results of Sunyaev-Zel'dovich effect (SZE) observations toward galaxy clusters. Using cm-wave receivers mounted on the OVRO and BIMA mm-wave arrays we have obtained high signal to noise images of the effect for more than 20 clusters. We present current estimates of the Hubble constant and cosmological parameters and discuss the potential of conducting statistical studies with large SZE cluster samples. | Over the last few years there has been a tremendous increase in the study of galaxy clusters as cosmological probes, initially through the use of X-ray emission observations, and in recent years, through the use of Sunyaev-Zel'dovich effect (SZE). Briefly, the SZE is a distortion of the cosmic microwave background (CMB) radiation by inverse-Compton scattering of thermal electrons within the hot intercluster medium (Sunyaev \& Zel'dovich 1980, see Birkinshaw 1998 for a recent review). The change in the CMB brightness temperature observed is: \begin{equation} \frac{\Delta T}{T_{\rm CMB}} = \left[ \frac{x (e^{x}+1)}{e^{x}-1} -4 \right] \int \left(\frac{k_B T_e}{m_e c^2}\right) n_e \sigma_T dl, \end{equation} where $x = h \nu/k_B T_{\rm CMB}$, and $n_e$, $T_e$ and $\sigma_T$ are the electron density, electron temperature and the cross section for Thomson scattering. The integral is performed along the line of sight through the cluster. The other important observable of the hot intercluster gas is the thermal Bremsstrahlung X-ray emission, whose surface brightness $S_X$ can be written as: \begin{equation} S_X = \frac{1}{4 \pi (1+z)^3} \int n ^{2}_{e} \Lambda_e dl, \end{equation} where $z$ is the redshift and $\Lambda_e(\Delta E,T_e)$ is the X-ray spectral emissivity of the cluster gas due to thermal Bremsstrahlung within a certain energy band $\Delta E$. By combining the intensity of the SZE and the X-ray emission observations, and knowing the cluster gas temperature $T_e$, the angular diameter distance to the cluster can be derived due to the different dependence of the X-ray emission and SZE on the electron density, $n_e$. Combining such distance measurements with redshift allows a determination of the Hubble constant, H$_0$, as a function of certain cosmological parameters (e.g., Hughes \& Birkinshaw 1998a). If distance measurements for a sample of clusters exist, then the angular diameter distance with redshift relation can be used to put constraints on the cosmological models, similar to current supernovae constraints at high redshift. | 98 | 4 | astro-ph9804149_arXiv.txt |
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9804 | astro-ph9804133_arXiv.txt | A new method of shower-image analysis is presented which appears very powerful as applied to those Cherenkov Imaging Telescopes with very high definition imaging capability. It provides hadron rejection on the basis of a single cut on the image shape, and simultaneously determines the energy of the electromagnetic shower and the position of the shower axis with respect to the detector. The source location is also reconstructed for each individual $\gamma$-ray shower, even with one single telescope, so for a point source the hadron rejection can be further improved. As an example, this new method is applied to data from the C{\small AT} (Cherenkov Array at Th\'emis) imaging telescope, which has been operational since Autumn, 1996. | Many of the sources in the E{\small GRET} catalogue \cite{egret} have well-identified radio, optical, or X-ray counterparts, which give the source position to an accuracy much better than can be achieved by {\small ACT} telescopes. Such sources with known position are usually placed at the centre of the field of an imaging telescope. However, many unidentified sources in the E{\small GRET} catalogue are located in error boxes with typical size of 1$^\circ$. Ground-based Cherenkov detectors are, in principle, able to localize such sources with a higher precision. In order to observe these sources, different methods have been developed by imaging Cherenkov telescope groups. Stereoscopy is the most direct way to find the direction of a source, but requires at least two telescopes and reduces the collection area \cite{hegratel}. For single-telescope experiments with sufficient background rejection on an image-shape criterion, it is possible to perform de-localized analyses assuming the source at the different points of a grid covering the field of view or by examining the distribution of the intersections of the image axes \cite{grid}. However, an event-by-event analysis method such as that described here is preferable since in the former methods the signal is more easily drowned-out by the background. The method has been tested on simulated $\gamma$-images provided by the Monte-Carlo simulation program described above and a realistic simulation of the detector response, including the measured variation in collection efficiency and gain between the phototubes, and measured wavelength response of the mirrors and Winston cones. For optimization of the cut values, these simulated $\gamma$-images have been used together with the real background events from data from off-source runs with the C{\small AT} imaging telescope. Gammas from a point-like source with the Crab nebula spectrum \cite{catcrab} (see equation (1) above) were simulated at various elevations. The capability of the method both for source detection and for source spectrum measurement have been examined. \subsection{Source detection} As applied to the data, the method consists of minimizing the $\chi^2$ with respect to $E_\gamma$, $D$, $\vec\xi$, and $\phi$. Fig~\ref{fig:chi2} shows the $\chi^2$ probability distributions obtained with this fit for the simulated $\gamma$-ray events and real background events. \begin{figure} \epsfxsize=14.5 cm \leavevmode \centering \epsffile[0 25 590 520]{fig6.eps} \caption{$\chi^2$ probability distribution for a fit with the source coordinates considered as free parameters (constraint is from shape alone): a) vertical $\gamma$-ray showers; b) real background (off-source data) showers. The upper line in each figure is for all events above threshold, the shaded distributions for events with a fitted energy, $E_{\mathrm{f}}$, greater than $350\:{\mathrm {GeV}}$ and a fitted impact parameter, $D_{\mathrm{f}}$, between 30 and $125\:{\mathrm{m}}$.} \label{fig:chi2} \end{figure} A cut on the $\chi^2$ probability value, $P(\chi^2)$, provides a selection of $\gamma$-like events on the basis of the image shape alone. The reconstructed angular origins obtained for simulated $\gamma$-events accumulate around the actual source position which in this case is at the centre of the field. The dispersion around the actual source position has a typical {\small RMS} spread of $0.14^\circ$. For each event, the accuracy of the angular origin determination is better by a factor two in the direction perpendicular to the image axis than in the direction of the image axis (Fig.~\ref{fig:pointerr}). \begin{figure} \epsfxsize=14.5 cm \leavevmode \centering \epsffile[20 25 540 230]{fig9.eps} \caption{Distributions of longitudinal and transverse errors (with respect to the image major axis) in the reconstruction of the source position for vertical showers with $P(\chi^2)>0.2$. The small bump at negative values of longitudinal error results from wrong direction reconstruction (mainly for events close to the energy threshold). Shaded histograms correspond to events with a fitted energy greater than $350\:{\mathrm GeV}$.} \label{fig:pointerr} \end{figure} The {\small RMS} longitudinal error typically varies from $0.2^\circ$ to $0.1^\circ$ as the energy varies from the threshold to $2\:{\mathrm {TeV}}$. The angular origins obtained for background events are spread over the whole field with an approximately Gaussian distribution with a $1.8^\circ$ {\small FWHM}. Since this distribution is fairly flat, 2-dimensional skymaps of the angular origins of the showers could be used for source detection, as can be seen from the reconstructed positions in data taken on Markarian 501 in Fig.~\ref{fig:m501}. \begin{figure} \epsfxsize=14.5 cm \leavevmode \centering \epsffile[0 25 580 520]{fig10.eps} \caption{The distribution of reconstructed angular origins for the data from a 30-minute run on Markarian 501 from April 16, 1997, for events with $P(\chi^2)>0.2$. The concentric circles represent the trigger region and the small-pixel region, respectively. During the run, the optic axis described the small arc indicated due to the mechanical flexibility of the structure, which is monitored as described in [9]. The number of events reconstructed in each bin of $(0.05{^0})^2$ is shown. No background subtraction has been performed.} \label{fig:m501} \end{figure} The errors in angular reconstruction given above are for a single shower; a point source with poorly defined position could thus be localized to $\sim 1-2'$ with the combination of $\sim 100$ such events. For the present, a conservative procedure of monitoring the background is used, based on the pointing angle $\alpha$, similar to the ``orientation'' of Whipple \cite{scuts}: $\alpha$ is the angle at the image barycentre between the actual source position and the reconstructed source position. The pointing angle does not use the full information contained in the results of the fit, but has a fairly flat distribution from $0^\circ$ to about $120^\circ$ for background events, which allows the background level to be easily monitored. The cut on $\alpha$ is more efficient than a cut on the angular distance between the source position and the reconstructed $\gamma$ origin since, as seen in Fig.~\ref{fig:pointerr}, the position reconstructed is not symmetric about the source position. The distribution of $\alpha$ for $\gamma$-events from a Crab-like source exhibits a peak at $0^\circ$ (Fig.~\ref{fig:alpha}) and a small accumulation at $180^\circ$ corresponding to events which are wrongly found to point away from the source. \begin{figure} \epsfxsize=14.5 cm \leavevmode \centering \epsffile[5 25 580 520]{fig7.eps} \caption{Distribution of the pointing angle $\alpha$ for events with $P(\chi^2)>0.2$ (constraint from shape alone): a) vertical $\gamma$-ray showers; b) Real background (off-source data) showers. The upper line is for all events above threshold, the shaded histograms for events with a fitted energy greater than $350\:{\mathrm {GeV}}$.} \label{fig:alpha} \end{figure} Around 17\% of the $\gamma$-events from a Crab-like source are in this situation. The proportion of events with a wrongly reconstructed direction decreases with increasing energy, from 22\% at $200\:{\mathrm {GeV}}$ to 9\% at $600\:{\mathrm {GeV}}$ and 4\% at $1\:{\mathrm {TeV}}$. The significance of a signal is calculated using the usual formula: $(ON-OFF)/\sqrt{ON+OFF}$ \cite{lima}, assuming equal time on and off-source. The significance per hour on a simulated Crab-like source at zenith has been calculated for various cut values on $\alpha$ and $P(\chi^2)$ (Fig.~\ref{fig:signif1}.a). \begin{figure}[t] $$ \epsfxsize=6.5 cm (a)\epsffile[40 30 540 490]{fig8.eps} \; \epsfxsize=6.5 cm (b)\epsffile[40 30 540 490]{fig11.eps} $$ \vspace{-1cm} \caption{Detection significance for a Crab-like source at zenith in one hour of observation, shown as a function of the two cuts on $P(\chi^2)$ and $\alpha$: (a) for a source at the centre of the camera; (b) for a source at $1^\circ$ from the centre. The full lines indicate contours of equal significance; dotted lines show fixed $\gamma$-ray selection efficiency.} \label{fig:signif1} \end{figure} The best result in terms of both significance and efficiency for $\gamma$-events is obtained for a $P(\chi^2)>0.2$ and $\alpha<6^\circ$, which gives $5.1\sqrt{t}\ \sigma$ where $t$ is the on-source observation time in hours, retaining 34\% of the $\gamma$-events while giving a rejection factor of 120 on background events. This rejection factor is smaller than for some comparable experiments as there is a large rejection factor at the trigger level, allowing a moderate background rate of $15\:{\mathrm {Hz}}$ at the zenith. At $45^\circ$ from zenith the best significance for the same cuts falls to $2.7\sqrt{t}\ \sigma$. This is essentially due to the higher energy threshold, leading to a lower event rate; on the other hand, for the same $\gamma$-ray selection efficiency the background rejection factor is comparable to that at zenith. In order to estimate the efficiency of the $\chi^2$-method in the case of a source with a poorly-defined position, a simulated Crab-like source has been set on the edge of the trigger area ($1^\circ$ from the centre). In this case, the equivalent detection area is divided by a factor of the order of two. Even for a source with known position, when the source is not at the centre of the field it is possible to use one side of the camera as the off region for the other side \cite{offc}. For a Crab-like source at zenith, the same cuts in $P(\chi^2)$ and $\alpha$ as for a source in the centre of the camera give a $2.3\sqrt{t}\ \sigma$ significance, with a selection efficiency for $\gamma$'s of 36\% and a rejection factor of 116 above threshold (Fig.~\ref{fig:signif1}.b). This means that the on-source run time has to be five times larger for a source on the edge of the trigger area than for a source at the centre of the field for the same significance. The corresponding significance obtained at $45^\circ$ from zenith for an off-centre source is $1.5\sqrt{t}\ \sigma$. \subsection{$\gamma$-ray energy measurement} For a source detected with a strong enough significance, the energy spectrum can be studied by a detector with good energy resolution. The fit described above in which the source position is a free parameter gives a first estimate of the energy of each event to within about 30\%. However, more precise spectral studies can be carried out on point sources of $\gamma$-rays. The use of the source position as a constraint in the fit provides a higher accuracy for impact parameter measurement and, as a consequence, for energy measurement. If trigger selection effects are ignored in the Monte-Carlo program (thus accepting all events above $100\:{\mathrm {GeV}}$), the $\chi^2$ minimization with respect to $E_\gamma$, $D$, and $\phi$ provides an unbiased energy measurement within about 25\% (statistical error only). Trigger selection effects are small for events well above the threshold, as can be seen for simulated $400\:{\mathrm {GeV}}$ $\gamma$-rays in Fig.~\ref{fig:ener500}. \begin{figure} \epsfxsize=14.5 cm \leavevmode \centering \epsffile[15 25 550 300]{fig12.eps} \caption{Distribution of $\ln({E_{\mathrm {f}}}/{E_\gamma})$ for vertical $400\:{\mathrm {GeV}}$ $\gamma$-ray showers satisfying the selection cuts (the source location, $\vec{\xi}$, being fixed in the fit). The shaded histogram is further restricted to events with a fitted impact parameter $30\:{\mathrm {m}} < D_{\mathrm {f}} < 125\:{\mathrm {m}}$.} \label{fig:ener500} \end{figure} This figure also shows that the distribution of the fitted event energies about the true energy is Gaussian on a logarithmic scale. Consequently, the slope of a power-law spectrum can be directly estimated with this technique. Close to the threshold, however, the fitted energy $E_{\mathrm {f}}$ is overestimated as a consequence of the trigger selection. Similarly, the small remaining bias in $ \log (E_{\mathrm {f}}/E_\gamma)$ at $400\:{\mathrm {GeV}}$ is due to showers with large impact parameters for which the trigger selection is critical at this energy since the telescope is then located at the border of the light pool (Fig.~\ref{fig:densite}). This effect is largely removed if only showers with a fitted impact parameter $D_{\mathrm {f}}$ lower than $125\:{\mathrm {m}}$ are included (shaded histogram in Fig.~\ref{fig:ener500}). The bias induced by the trigger selection at different energies is best illustrated by plotting 68\% confidence intervals for $E_{\mathrm {f}}$ as a function of the true value $E_\gamma$ used in the simulation (Fig.~\ref{fig:enerfit}). \begin{figure} \epsfxsize=14.5 cm \leavevmode \centering \epsffile[5 25 580 520]{fig14.eps} \caption{Fitted energy, $E_{\mathrm {f}}$, versus true energy, $E_\gamma$, for $\gamma$-ray showers satisfying the selection cuts (the source location, $\vec{\xi}$, being fixed in the fit) and for which $30\:{\mathrm {m}}< D_{\mathrm {f}}\cos Z < 125\:{\mathrm {m}}$ ($Z=$~zenith angle). The shaded interval shows the 68\% confidence intervals for a source at the zenith. The effect of trigger selection on the energy estimation can be seen.} \label{fig:enerfit} \end{figure} It can be seen that for $E_f$ below $350\:{\mathrm {GeV}}$ for vertical showers, only an upper limit can be safely derived for $E_\gamma$. Therefore, spectrum measurement is reliable only above a spectrometric threshold, that is, in the region in which $E_{\mathrm {f}}$ depends linearly on $E$. It is somewhat higher than the nominal threshold which is relevant for source discovery. Fig.~\ref{fig:enerfit} also shows the variation of the average values of $\log(E_{\mathrm {f}})$ on $\log(E_\gamma)$ for zenith angles $0^{\circ}$, $30^{\circ}$, $45^{\circ}$, and $60^{\circ}$, showing the increase in the spectrometric thresholds with increasing zenith angle. Restricting to events with $E_{\mathrm {f}}$ above the spectrometric threshold and $30\:{\mathrm {m}}< D_{\mathrm {f}} \cos Z < 125\:{\mathrm {m}}$, the accuracy of the preceding method is about 20\% (statistical error only), independent of the zenith angle $Z$ up to $45^{\circ}$. | The method described above, based on a realistic analytic description of electromagnetic air showers, is best suited for those Cherenkov Imaging Telescopes with a high-resolution camera. The light distribution in the focal plane is fully exploited, yielding the shower direction from the asymmetry of the longitudinal profile as well as the source position in the focal plane. Selection of $\gamma$-rays on the basis of the image shape is performed by using a single $\chi^2$-variable instead of a series of cuts on various image parameters. By combining the $\chi^2$ probability cut and a direction ($\alpha$) cut, a significance of 5$\sigma$ per hour can be achieved for a Crab-like source at the zenith. A future development of the method would be to use the distribution of selected $\gamma$-ray origins on the celestial sphere together with the known energy-dependent point spread function of the method to estimate the significance by a maximum-likelihood method. With the C{\small AT} telescope ($250\:{\mathrm {GeV}}$ threshold), sources with the intensity of the Crab nebula can be detected in one hour. This has been confirmed with the results obtained in the 96/97 observing campaign. Moreover, sources with poorly defined position can be localized with an accuracy of the order of an arc minute on the basis of about 100 showers. The accuracy on $\gamma$-ray energy is of the order of 20--25\%. Biases induced by the trigger selection have been investigated; in particular, care should be taken in spectrum measurement, which is accurate only above a specific threshold, higher than that used for source detection. | 98 | 4 | astro-ph9804133_arXiv.txt |
9804 | astro-ph9804305_arXiv.txt | We analyze the stability of g-modes in variable white dwarfs with hydrogen envelopes. All the relevant physical processes take place in the outer layer of hydrogen rich material which consists of a radiative layer overlain by a convective envelope. The radiative layer contributes to mode damping because its opacity decreases upon compression and the amplitude of the Lagrangian pressure perturbation increases outward. The convective envelope is the seat of mode excitation because it acts as an insulating blanket with respect to the perturbed flux that enters it from below. A crucial point is that the convective motions respond to the instantaneous pulsational state. Driving exceeds damping by as much as a factor of two provided $\omega\tau_c\geq 1$, where $\omega$ is the radian frequency of the mode and $\tau_c\approx 4\tau_{\rm th}$ with $\tau_{\rm th}$ being the thermal time constant evaluated at the base of the convective envelope. As a white dwarf cools, its convection zone deepens, and modes of lower frequency become overstable. However, the deeper convection zone impedes the passage of flux perturbations from the base of the convection zone to the photosphere. Thus the photometric variation of a mode with constant velocity amplitude decreases. These factors account for the observed trend that longer period modes are found in cooler DAVs. Overstable modes have growth rates of order $\gamma\sim 1/(n\tau_\omega)$, where $n$ is the mode's radial order and $\tau_\omega$ is the thermal time-scale evaluated at the top of the mode's cavity. The growth time, $\gamma^{-1}$, ranges from hours for the longest period observed modes ($P\approx 20$ minutes) to thousands of years for those of shortest period ($P\approx 2 $ minutes). The linear growth time probably sets the time-scale for variations of mode amplitude and phase. This is consistent with observations showing that longer period modes are more variable than shorter period ones. Our investigation confirms many results obtained by Brickhill in his pioneering studies of ZZ Cetis. However, it suffers from at least two serious shortcomings. It is based on the quasiadiabatic approximation that strictly applies only in the limit $\omega\tau_c\gg 1$, and it ignores damping associated with turbulent viscosity in the convection zone. We will remove these shortcomings in future papers. | } ZZ Cetis, also called DAVs, are variable white dwarfs with hydrogen atmospheres. Their photometric variations are associated with nonradial gravity-modes (g-modes); for the first conclusive proof, see Robinson \etal (\cite{scaling-robinson82}). These stars have shallow surface convection zones overlying stably stratified interiors. As the result of gravitational settling, different elements are well separated . With increasing depth, the composition changes from hydrogen to helium, and then in most cases to a mixture of carbon and oxygen. From center to surface the luminosity is carried first by electron conduction, then by radiative diffusion, and finally by convection. Our aim is to describe the mechanism responsible for the overstability of g-modes in ZZ Ceti stars. This topic has received attention in the past. Initial calculations of overstable modes were presented in Dziembowski \& Koester (\cite{adia-dziem81}), Dolez \& Vauclair (\cite{adia-dolez81}), and Winget \etal (\cite{adia-winget82}). These were based on the assumption that the convective flux does not respond to pulsation; this is often referred to as the frozen convection hypothesis. Because hydrogen is partially ionized in the surface layers of ZZ Ceti stars, these workers attributed mode excitation to the $\kappa$-mechanism. In so doing, they ignored the fact that the thermal time-scale in the layer of partial ionization is many orders of magnitude smaller than the periods of the overstable modes. Pesnell (\cite{adia-pesnell87}) pointed out that in calculations such as those just referred to, mode excitation results from the outward decay of the perturbed radiative flux at the bottom of the convective envelope. He coined the term `convective blocking' for this excitation mechanism.\footnote{This mechanism was described in a general way by Cox \& Guili (\cite{adia-cox68}), and explained in more detail by Goldreich \& Keeley (\cite{adia-goldreich77}).} Although convective blocking is responsible for mode excitation in the above cited references, it does not occur in the convective envelopes of ZZ Ceti stars. This is because the dynamic time-scale for convective readjustment (i.e., convective turn-over time) in these stars is much shorter than the g-mode periods. Noting this, Brickhill (\cite{adia-brick83}, \cite{adia-brick90}, \cite{adia-brick91a}, \cite{adia-brick91b}) assumed that convection responds instantaneously to the pulsational state. He demonstrated that this leads to a new type of mode excitation, which he referred to as convective driving. Brickhill went on to presents the first physically consistent calculations of mode overstability, mode visibility, and instability strip width. Our investigation supports most of his conclusions. Additional support for convective driving is provided by Gautschy \etal (\cite{adia-gautschy96}) who found overstable modes in calculations in which convection is modeled by hydrodynamic simulation. In this paper we elucidate the manner in which instantaneous convective adjustment promotes mode overstability. We adopt the quasiadiabatic approximation in the radiative interior. We also ignore the effects of turbulent viscosity in the convection zone. These simplifications enable us to keep our investigation analytical, although we appeal to numerically computed stellar models and eigenfunctions for guidance. The DA white dwarf models we use are those produced by Bradley (\cite{scaling-bradley96}) for asteroseismology. Fully nonadiabatic results, which require numerical computations, will be reported in a subsequent paper. These modify the details, but not the principal conclusions arrived at in the present paper. The plan of our paper is as follows. The linearized wave equation is derived in \S \ref{sec:adia-prepare}. In \S \ref{sec:adia-perturb}, we evaluate the perturbations associated with a g-mode in different parts of the star. We devote \S \ref{sec:adia-driving} to the derivation of a simple overstability criterion. Relevant time-scales and the validity of the quasiadiabatic approximation are discussed in \S \ref{sec:adia-discussion}. The appendix contains derivations of convenient scaling relations for the dispersion relation, the WKB eigenfunction, and the amplitude normalization. | } \subsection{Time-Scales \label{sec:adia-relevant}} Three time-scales are relevant for convective driving in DAVs. The first is the period of an overstable g-mode, $P = 2\pi/\omega$, which is typically of order a few hundred seconds. The second is the dynamical time constant, $\tcv \sim H_p/\vcv$, on which convective motions respond to perturbations; $\tcv \leq 1 \s$ throughout the convection zones of even the coolest ZZ Cetis. This is why the convective motions adjust to the instantaneous pulsational state. The third is the thermal time constant, $\tau_c$, during which the convection zone can bottle up flux perturbations that enter it from below. Given the central role of $\tau_c$, we elaborate on its relation both to $\tcv$ and to the more conventional definition of thermal time constant, $\tau_{\rm th}$, at depth $z$. The latter is the heat capacity of the material above that depth divided by the luminosity. In a plane parallel, fully ionized atmosphere this is equivalent to \begin{equation} \tau_{\rm th}\equiv {1\over F} \int_0^z dz \, c_p\, {\rho k_B\over m_p T} \approx {5pz\over 7F}. \label{eq:adia-tauth}\end{equation} Appeal to equation \refnew{eq:adia-Fvcv} establishes that inside the convection zone \begin{equation} {\tcv\over \tau_{\rm th}}\sim \left({\vcv\over c_s}\right)^2\ll 1. \end{equation} Now $\tau_c\equiv (B+C)\tau_b$, where $\tau_b$ is defined by equation \refnew{eq:adia-taub}. To the extent that $c_p\approx 5$ is constant in the convection zone, $\tau_b\approx \tau_{\rm th}/5$, where the latter is evaluated at $z_b$. Next we address the relation between $\tau_c$ and $\tau_b$. Here we are concerned with the relatively large value of $B+C$, typically about 20 for DAVs.\footnote{In our models, $B$ and $C$ have comparable value.} Recall from equations \refnew{eq:adia-Dsb} and \refnew{eq:adia-DFph} that \begin{equation} {\delta F_{\ph} \over F}\approx {\delta s_b\over B+C}. \label{eq:adia-dFds}\end{equation} So the photosphere and superadiabatic layer add an insulating blanket on top of the convection zone. The large value of $B$ follows because the photospheres of DAVs are composed of lightly ionized hydrogen. In this state, the values of $\kappa_T$ and $s_T$ are both large and positive; typical values in the middle of the instability strip are $\kappa_T\approx 6$ and $s_T\approx 24$. The large and positive $\kappa_T$ arises because the population of hydrogen atoms in excited states which the ambient radiation field can photoionize increases exponentially with increasing $T$. The large and positive $s_T$ occurs as the ionization fraction increases exponentially with increasing $T$, and the much larger entropy contributed by a free as compared to a bound electron. The large and positive value of $C$ reflects the increase in entropy gradient that accompanies an increase in convective flux. It is obtained from mixing length theory with an unperturbed mixing length. \subsection{Validity of the Quasiadiabatic Approximation} The validity of the quasiadiabatic approximation requires that the nonadiabatic parts of the expressions for $\delta \rho/\rho$ and $\delta T/T$, as given by equations \refnew{eq:adia-eqdelrho} and \refnew{eq:adia-eqdelT}, be small in comparison to the adiabatic parts. Thus the ratio \begin{equation} {\cal R_{\rm na}}\equiv {\delta s\over\delta p/p} \label{eq:adia-rnonad}\end{equation} is a quantitative measure of nonadiabaticity. We estimate ${\cal R}_{\rm na}$ for the radiative interior and the convection zone. We calculate $\delta s$ in the radiative interior from \begin{equation} \delta s \approx {iF\over \omega}{m_p\over\rho k_B T}{d\over dz}\left({\delta F\over F}\right). \label{eq:adia-delsrad}\end{equation} In the upper evanescent layer, $z_b<z<z_\omega$, this leads to \begin{equation} |{\cal R_{\rm na}}|\sim {1\over \omega\tau_{\th}}, \label{eq:adia-ratioevu}\end{equation} whereas in the propagating cavity, $z>z_\omega$, we find \begin{equation} |{\cal R_{\rm na}}|\sim {1\over \omega\tau_{\th}}\left({z\over z_\omega}\right). \label{eq:adia-ratioevl}\end{equation} Nonadiabatic effects in the radiative interior are maximal at $z=z_b$, where \begin{equation} |{\cal R_{\rm na}}|\sim {1\over \omega\tau_b}, \label{eq:adia-ratioevmax}\end{equation} since $\tau_{\rm th}/z$ increases with depth. The measure of nonadiabaticity in the convection zone is given by equation \refnew{eq:adia-Dsb}. Since $\omega\tau_b>1$ is required for the validity of the quasiadiabatic approximation in the radiative zone, we restrict consideration to the limiting case $\omega\tau_c\gg 1$. In this limit, equation \refnew{eq:adia-Dsb} yields \begin{equation} |{\cal R_{\rm na}}|\sim {1\over \omega\tau_b}, \label{eq:adia-ratiocvz}\end{equation} which is identical to the value arrived at for the radiative zone in equation \refnew{eq:adia-ratioevmax}. The requirement $\omega\tau_b\gtrsim 1$ for the validity of the quasiadiabatic approximation severely limits the applicability of the current investigation. The perturbed flux at the photosphere is related to that at the bottom of the convection zone by \begin{equation} {\Delta F_{\ph}\over F}\approx {1\over 1-i\omega\tau_c}{\Delta F_b\over F}. \label{eq:adia-relDelF}\end{equation} Since $\tau_c$ is at least an order of magnitude larger than $\tau_b$, modes with $\omega\tau_b\gtrsim 1$ are likely to exhibit small photometric variations. However, this may not render them undetectable because their horizontal velocity perturbations pass undiminished through the convection zone. \subsection{Brickhill's Papers\label{subsec:adia-brickhill}} Our investigation is closely related to studies of ZZ Cetis by Brickhill (\cite{adia-brick83},\cite{adia-brick90},\cite{adia-brick91a}). Brickhill recognized that the convective flux must respond to the instantaneous pulsational state. To determine the manner in which the convection zone changes during a pulsational cycle, he compared equilibrium stellar models covering a narrow range of effective temperature. Brickhill provided a physical description of convective driving and obtained an overstability criterion equivalent to ours.\footnote{Our time constant $\tau_c$ is equivalent to the quantity $D$ which Brickhill defined in equation (9) of his 1983 paper.} Moreover, he recognized that the convection zone reduces the perturbed flux and delays its phase. Our excuses for revisiting this topic are that Brickhill's papers are not widely appreciated, that our approach is different from his, and that our paper provides the foundation for future papers which will examine issues beyond those he treated. \begin{appendix} | 98 | 4 | astro-ph9804305_arXiv.txt |
9804 | astro-ph9804075_arXiv.txt | The decay of massive neutrinos to final states containing only invisible particles is poorly constrained experimentally. In this letter we describe the constraints that can be put on neutrino mass and lifetime using CMBR measurements. We find that very tight lifetime limits on neutrinos in the mass range 10 eV - 100 keV can be derived using CMBR data from upcoming satellite measurements. \\ Keywords: Neutrino decay, Cosmology: Theory, Cosmic Microwave Background \\ PACS: 13.35.Hb, 14.60.St, 98.70.Vc | 98 | 4 | astro-ph9804075_arXiv.txt |
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9804 | astro-ph9804243_arXiv.txt | Although the unified formula for $\gamma$-ray absorption process involving both the magnetic field and a perpendicular electric field derived by Daugherty \& Lerche (1975) is correct, we argued in this paper that their conclusion that the induced electric fields are important in the pair formation process in the pulsar magnetospheres is wrong and misleading. The key point is that usually the direction of a $\gamma$ photon at the emission point observed in the laboratory frame should be $(v/c, 0, [1-(v/c)^2]^{1/2})$ rather than $(0, 0, 1)$, where $v$ is the co-rotating velocity. This emission direction is just the one which results in zero attenuation coefficient of the $\gamma$ photon. Calculation shows that after the photon has moved a distance, its direction lead to the result that the induced electric field is also of minor importance. Thus only $\gamma-B$ process is the important mechanism for the pair production in the pulsar magnetospheres. The implications of the modification by ejecting the induced electric field are also discussed. | Pair production process plays an important role in pulsar physics. It is not only a necessary process for the multiplication of the particles to account for emissions of different bands from pulsars (e.g. Sturrock 1971; Ruderman \& Sutherland 1975; Arons \& Scharleman 1979; Arons 1983; Cheng, Ho, \& Ruderman 1986), but also an important mechanism to absorb $\gamma$-rays produced in the pulsar magnetospheres, especially near the polar cap region (e.g. Hardee 1977; Harding, Tademaru, \& Esposito 1978; Harding 1981; Daugherty \& Harding 1982, 1996; Zhao {\it et al.} 1989; Lu \& Shi 1990; Lu, Wei, \& Song 1994; Dermer \& Sturner 1994; Sturner, Dermer, \& Michel 1995; Wei, Song, \& Lu 1997). Furthermore, the way by which the $\gamma$-rays are absorbed is also the key factor to limit the parameters of the inner magnetospheric accelerators of pulsars (e.g. Ruderman \& Sutherland 1975, hereafter RS75; Zhang \& Qiao 1996; Qiao \& Zhang 1996; Zhang {\it et al.} 1997a; Zhang, Qiao, \& Han 1997b, hereafter ZQH97b). Pair formation in intense magnetic fields ($\gamma-B$ process) has been studied explicitly by different authors (e.g. Erber 1966; Tsai \& Erber 1974; Daugherty \& Harding 1983, Rifert, M\'{e}sz\'{a}ros \& Bagoly 1989), and its importance in pulsar physics was first pointed out by Sturrock (1971). Daugherty \& Lerche (1975, hereafter DL75) first dealt with the case involving a relatively weaker electric field perpendicular to the magnetic field (${\bf E}^2-{\bf B}^2\le0$, ${\bf E\cdot B}=0$), and came to a unified formula of the attenuation coefficient of the $\gamma$ photons. The more general case involving both the perpendicular and the parallel components of the electric field with respect to the magnetic field ($E_\perp$ and $E_\parallel$) was presented by Daugherty \& Lerche (1976) and Urrutia (1978). In the specific case of pulsars, although $E_\parallel$ is usually sufficiently small so that its effect is negligible, $E_\perp$ induced by the fast spin of the neutron stars was demonstrated to be very important in pair formation process by DL75. This leads many authors to take this effect seriously into account in their studies (e.g. Hardee 1977; Lu \& Shi 1990; Lu, Wei, \& Song 1994; Qiao \& Zhang 1996). In this paper, we'll argue that although the unified formula of DL75 is correct, their conclusion that the induced electric fields are important in the pair formation process in the pulsar magnetospheres is wrong and misleading. The detailed argument is presented in Section 2 and Section 3. Finally, we discuss the possible implications of this modification. | Although the unified formula for the pair production process involving both the magnetic field and the perpendicular electric field derived by DL75 is correct, we have argued in this paper that their conclusion that the induced electric fields are important in the pair formation process in the pulsar magnetospheres is wrong and misleading at least for the ``aligned rotator'' case. At the emission point, the photon emitted by a certain mechanism (e.g. the curvature radiation or the inverse Compton scattering) just moves along the very direction in which the attenuation coefficient is zero. Considering the propagation of the photon, we found that this rotation-induced electric field still plays a minor role in the $\gamma$-ray absorption process in the polar cap region of a pulsar. For the general case of an ``oblique rotator'' in which the magnetic and the rotational axes are misaligned, the co-rotating velocity is no more perpendicular to the magnetic field so that $\theta_B^{\prime}$ (also $\theta_u^{\prime}$ and $\theta_\gamma^{\prime}$) is not $\pi/2$. The photon direction consequently deviate from $(v/c, 0, [1-(v/c)^2]^{1/2})$ slightly. From Fig.2, we see that the attenuation coefficient is also very small around the direction $(v/c, 0, [1-(v/c)^2]^{1/2})$, so that the conclusion that the induced electric field plays a minor role in $\gamma$-ray absorption still holds for the oblique rotator case. Actually, DL75's result can only be applied to the aligned case strictly, since generally the co-rotating velocity ${\bf v}_{r}$ is not equal to ${\bf v}_{drift}=c({\bf E\times B})/B^2$, with which one can define a frame where the electric field vanishes completely. In the co-rotating frame of an oblique rotator, an electric field component parallel to the magnetic field will still remain so that DL75's application condition fails. Our results in this paper may have some implications for some previous studies which regard the electric field as the important effect of $\gamma$-ray absorption (e.g. Hardee 1977; Zhao {\it et al.} 1989; Lu \& Shi 1990; Lu, Wei, \& Song 1994; Qiao \& Zhang 1996). Although the polar cap models of the $\gamma$-ray pulsars are by all means sound in principle, their concrete details will alter much by ejecting the induced electric field. Specifically, based on DL75's result, Hardee (1977) got an absolute upper limit to the photon energies $$E_\gamma <9.6\times 10^9 B_{12}^{-1}r_6^2 P {\rm eV}, \eqno(6)$$ (his Eq.(38)) at which escape of the $\gamma$-rays from the magnetosphere is possible. It should be replaced by a threshold in the pure magnetic field absorption scheme $$E_\gamma <2.6\times 10^9 B_{12}^{-1} P^{1/2} {\rm eV} \eqno(7)$$ (Wei, Song, \& Lu 1997, their Eq.(10), $r_6=1$ and the last open field line is assumed). Thus the generation order parameters proposed by Lu, Wei, \& Song (1994) should take the form in Wei, Song, \& Lu (1997, their Eq.(17)). The three boundary lines (birth line, death line and appearance line) in the $\dot P-P$ diagram of pulsars derived by Qiao \& Zhang (1996) are also based on the electric field absorption. The details will also be changed by ejecting the electric field, but the picture still remains and may give a hint to us about the magnetic field configuration in the neutron star vicinity (Qiao \& Zhang, discussions). | 98 | 4 | astro-ph9804243_arXiv.txt |
9804 | astro-ph9804257_arXiv.txt | Molecules dominate the cooling function of neutral metal-poor gas at high density. Observation of molecules at high redshift is thus an important tool toward understanding the physical conditions prevailing in collapsing gas. Up to now, detections are sparse because of small filling factor and/or sensitivity limitations. However, we are at an exciting time where new capabilities offer the propect of a systematic search either in absorption using the UV Lyman-Werner H$_2$ bands or in emission using the CO emission lines redshifted in the sub-millimeter. | \subsection{Introduction} QSO absorption line systems probe the baryonic matter over most of the history of the Universe (0~$<$~$z$~$<$~5). The so-called damped Ly$\alpha$ (hereafter DLA) systems are characterized by a very large H~{\sc i} column density ($N$(H~{\sc i})~$>$~2$\times$10$^{20}$ ~cm$^{-2}$), similar to what is usually observed through local spiral disks. The case for these systems to be produced by proto-galactic disks is supported by the fact that the cosmological density of gas associated with these systems is of the same order of magnitude as the cosmological density of stars at present epochs (Wolfe 1996). The presence of heavy elements ($Z \sim 1/10 ~ Z_\odot$) and the redshift evolution of metallicity suggest ongoing star formation activities in these systems (Lu et al. 1996, Pettini et al. 1996, 1997). Moreover, strong metal line systems have been demonstrated to be associated with galaxies at low and intermediate $z$ (e.g. Bergeron \& Boiss\'e 1991). It has also been shown that the profiles of the lines arising in the neutral gas show evidence for rotation (Wolfe 1996, Prochaska \& Wolfe 1997). Whether these arguments are enough to demonstrate that DLA systems arise in large disks is a matter of debate however. Indeed simulations have shown that the progenitors of present day disks of galaxies could look like an aggregate of well separated dense clumps at high redshift. The kinematics could be explained by relative motions of the clumps with very little rotation (Haehnelt et al. 1997, Ledoux et al. 1998). Moreover, using {\sl HST} high spatial resolution images of the field of seven quasars whose spectra contain DLA lines at intermediate redshifts (0.4~$<$~$z$~$<$~1), Le~Brun et al. (1997) show that, in all cases, at least one galaxy candidate is present within 4~arcsec from the quasar. There is no dominant morphological type in their sample: three candidates are spiral galaxies, three are compact objects and two are amorphous low surface brightness galaxies. Therefore, although the nature of the DLA systems is unclear they trace the densest regions of the Universe where star formation occurs.\par \subsection{Molecular hydrogen} \begin{figure} \epsfysize=15cm % \epsfxsize=11cm % \hspace{1.5cm}\epsfbox{q0528mol.ps} % \caption[h]{Fit result for a few rotational transitions of the H$_2$ Lyman absorption bands in the $z_{\rm abs}$~=~2.8112 system toward PKS~0528--250. The spectrum has been obtained with the echelle spectrograph CASPEC attached on the ESO 3.6~m at La Silla. The resolution is $R$~=~36000 and the integration time 5~hours.} \end{figure} It is thus surprising that despite intensive searches, the amount of H$_2$ molecules seems quite low in damped Ly$\alpha$ systems in contrast to what is observed in our own galaxy. Two detections of H$_2$ molecules in high redshift DLA systems have been reported. Recently Ge \& Bechtold (1997) have found strong absorptions in the $z_{\rm abs}$~=~1.9731 DLA system toward Q~0013--004. They derive $N$(H$_2$)~=~6.9$\times$ 10$^{19}$~cm$^{-2}$, $b$~=~15~km~s$^{-1}$, $T_{\rm ex}$~$\sim$~70~K and $n$(H)~$\sim$~300~cm$^{-3}$ for a total hydrogen column density $N$(H)~=~6.4$\times$10$^{20}$~cm$^{-2}$. This system has by far the largest H$_2$ abundance $f$~=~2$N$(H$_2$)/[2$N$(H$_2$)~+~$N$(H~{\sc i})] $\sim$~0.22$\pm$0.05 observed in DLA systems. However the exact number should be confirmed using higher resolution data. Other searches have led to much smaller values or upper limits ($f$~$<$~10$^{-6}$, Black et al. 1987, Chaffee et al. 1988, Levshakov et al. 1992). Table~1 summarizes the caracteristics of damped Ly$\alpha$ systems that have been searched for molecules. Levshakov \& Varshalovich (1985) suggested that molecules could be present toward PKS~0528--250 at a redshift ($z_{\rm abs}$~=~2.8112), slighly larger than the emission redshift of the quasar. This claim has been confirmed by Foltz et al. (1988) using a 1~\AA~ resolution spectrum. The latter authors derive $N$(H$_2$)~=~10$^{18}$~cm$^{-2}$, $b$~=~5~km~s$^{-1}$, $T_{\rm ex}$~=~100~K and log~$N$(H~{\sc i})~=~21.1$\pm$0.3. By fitting the damped absorption together with the Ly$\alpha$ emission from the quasar, M\o ller \& Warren (1993) find log~$N$(H~{\sc i})~=~21.35. New high resolution data has been recently obtained by Srianand \& Petitjean (1998). They estimate the column density of H$_2$ molecules $N$(H$_2$)~$\sim$~6$\times$10$^{16}$~cm$^{-2}$ and the fractional abundance of H$_2$, $f$~=~5.4$\times$10$^{-5}$ (see Fig.~1). The excitation temperature derived for different transitions suggests that the kinetic temperature of the cloud is $\sim$200~K and the density $n$~$\sim$~1000~cm$^{-3}$. The cloud has therefore a dimension along the line of sight smaller than 1~pc. Since it obscurs the broad-line emission region, its transverse dimension should be larger than 10~pc. Upper limits are obtained on the column densities of C~{\sc i} ($<$~10$^{12.7}$~cm$^{-2}$) and CO ($<$~10$^{13.2}$~cm$^{-2}$; $N$(CO)/$N$(H~{\sc i})~$<$~7$\times$10$^{-9}$). It is suggested that the ratio $N$(H$_2$)/$N$(C~{\sc i}) is a useful indicator of the physical conditions in the absorber. Photo-ionization models show that radiation fields with spectra similar to typical AGNs or starbursts are unable to reproduce all the constraints and in particular the surprizingly small $N$(C~{\sc i})/$N$(H$_2$) and $N$(Mg~{\sc i})/$N$(H$_2$) ratios. In view of the models explored, the most likely ionizing spectrum is a composite of a UV-"big bump" possibly produced by a local starburst and a power-law spectrum from the QSO that provides the X-rays. This suggests that the gas is not predominantly ionized by the quasar and that star-formation may occur in the clouds, a conclusion reached as well by Warren \& M\o ller (1996) and Ge et al. (1997). Dust is needed to explain the production of molecules in the cloud. The amount of dust is broadly consistent with the [Cr/Zn] abundance determination. \par\noindent \vspace{0.5cm} % \par\noindent \centerline{\bf Tab. 1 - H$_2$ molecules in DLA systems} \begin{table}[h] \hspace{1.5cm} % \begin{tabular}{|l|c|c|c|c|c|c|} \hline Name & 000-263 & 0013-004 & 0100+130 & 0528-250 & 1331+170 & 1337+113 \\ \hline $z_{\rm em}$ & 4.110 & 2.084 & 2.681 & 2.770 & 2.081 & 2.919\\ $z_{\rm abs}$ & 3.391 & 1.9731 & 2.309 & 2.811 & 1.776 & 2.796\\ $N$(HI) (10$^{21}$~cm$^{-2}$ & 2.0 & 0.64 & 2.5 & 2.2 & 1.5 & 0.80\\ $N$(H$_2$) (10$^{16}$~cm$^{-2}$& $<$~0.3 & 6900 & $<$~0.5 & 6 & ... & $<$~5\\ $f_{{\rm H}2}$(10$^{-4}$) & $<$~0.03 & 2200 & $<$~0.04 & 0.5 & ... & $<$~1.3\\ \hline \multicolumn{7}{l}{Levshakov et al. (1992); Ge \& Bechtold (1997); Srianand \& Petitjean (1998)} \\ \end{tabular} \label{tsys} \end{table} \begin{figure} \centerline{ \psfig{figure=boisse.ps,width=9.cm,height=8.cm,angle=270} } \caption[h]{ [Zn/H] versus log~$N$(H~{\sc i}) for 37 damped Ly$\alpha$ systems; small symbols correspond to $z_{\rm abs}$~$<$~2.15 and large symbols to $z_{\rm abs}$~$>$~2.15. The line in the upper right corresponds to $N$(Zn~{\sc ii}) = 1.4$\times$10$^{13}$~cm$^{-2}$ or to Galactic material inducing $A_{\rm V}$~$\sim$~0.27. The figure is taken from Boiss\'e et al. (1998).} \label{boisse} \end{figure} \subsection{Is there a bias against detection of H$_2$ molecules ?} The small number of H$_2$ detections in damped systems is intriguing. Indeed in the interstellar medium of our Galaxy, all the clouds with log~$N$(H~{\sc i})~$>$~21 have log~$N$(H$_2$)~$>$~19 (Jenkins \& Shaya 1979). Formation of H$_2$ is expected on the surface of dust grains if the gas is cool, dense and mostly neutral, and from the formation of negative hydrogen if the gas is warm and dust free (see e.g. Jenkins \& Peimbert 1997). Destruction is mainly due to UV photons. The effective photodissociation of H$_2$ takes place in the energy range 11.1--13.6 eV, through Lyman-Werner band line absorption. In the DLA system toward PKS~0528-250, (i) abundances are of the order of 0.1~$Z_{\odot}$; (ii) the ratio [Cr/Zn] indicates a depletion factor into dust-grains of the order of half of that in the Galactic ISM; (iii) although it has been shown that the cloud is located at a distance larger than 10~kpc from the quasar, it is still close to it and exposed to its UV flux. Nonetheless, molecular hydrogen is detected. This suggests that indeed, molecular hydrogen should be seen in most of the damped systems. The small number of detections may be explained if observations are biased against the presence of molecules. Indeed it can be speculated that molecules should be found predominantly in gas with a non negligible amount of dust. However the corresponding extinction of the background quasar due to the dust in the damped system could be large enough to drop the quasar out of the sample of quasars that are usually observed for such studies. Boiss\'e et al. (1998) notice that for the damped systems studied up to now, the larger the H~{\sc i} column density, the smaller the abundances (see Fig.~\ref{boisse}). This suggests that the high column density DLA systems detected up to now are those with the smallest metallicities and consequently those with the smallest amount of dust. One way to clear up this problem is to observe a complete sample of quasars (if possible constructed without color selection) and to search for high column density damped systems. | 98 | 4 | astro-ph9804257_arXiv.txt |
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9804 | astro-ph9804061_arXiv.txt | We have computed stellar evolutionary models for stars in a mass range characteristic of Cepheid variables ($3<m/\Msol<12$) for different metallicities representative of the Galaxy and the Magellanic Clouds populations. The stellar evolution calculations are coupled to a linear non adiabatic stability analysis to get self-consistent mass-period-luminosity relations. The period - luminosity relation as a function of metallicity is analysed and compared to the recent EROS observations in the Magellanic Clouds. The models reproduce the observed width of the instability strips for the SMC and LMC. We determine a statistical P-L relationship, taking into account the evolutionary timescales and a mass distribution given by a Salpeter mass function. Excellent agreement is found with the SMC PL relationship determined by Sasselov et al. (1997). The models reproduce the change of slope in the P-L relationship near $P\sim 2.5$ days discovered recently by the EROS collaboration (Bauer 1997; Bauer et al. 1998) and thus explain this feature in term of stellar evolution. Some discrepancy, however, remains for the LMC Cepheids. The models are also in good agreement with Beat Cepheids observed by the MACHO and EROS collaborations. We show that most of the 1H/2H Beat Cepheids have not yet ignited central helium burning; they are just evolving off the Main Sequence toward the red giant branch. | In Fig. 2a-b we compare our calculations with observed fundamental mode pulsators in the LMC and SMC in a P - $M_V$ diagram. The EROS ($B_E$, $R_E$) magnitudes were transformed into the Johnson-V magnitude according to Beaulieu et al. (1995). We adopt the extinction E(B-V) = 0.10 for the LMC and E(B-V) = 0.125 for the SMC with R$_V$ = 3.3 and distance moduli $(m-M)_0$ = 18.5 for the LMC and $(m-M)_0$ = 19.13 for the SMC (cf. Beaulieu et al. 1997a). Comparison is also made with the Laney \& Stobie (1994, LS94) data, which extend to longer periods than the EROS data. The transformation of the theoretical quantities (L, $\te$) into M$_V$ are based on the Allard and Hauschildt (1998) most recent atmosphere models. These models originally developed for M-dwarfs do not extend below log g $<$ 3.5. We therefore use the bolometric corrections at constant log g = 3.5 but take metallicity effects into account. We verified however that gravity effects in the V-band are small. We determine a theoretical statistical P - $\mv$ relationship under the form $\mv = A \, {\rm log}P + B$ with a weighted least square fit by minimizing the quantity : $$ Q = \sum_i \alpha _i \big( M_{Vi} - A \, {\rm log}P_i - B \big) ^2 $$ where the summation extends over all the stellar masses. The coefficients $\alpha_i(m_i, t)$ depend on the mass distribution and on the evolutionary times. We adopt a Salpeter mass function (MF) $dN/dm \propto m^{-2.35}$ and the time dependence derives directly from the coupled evolution and pulsation calculations. Although the present day MF in the MC may differ from an initial Salpeter MF, we verified that variations of the slope of the MF between -2 and -4 barely affect the slope of the PL relationship. Such a steep MF favors the lowest mass stars because (i) of the number of stars itself and (ii) of the longer time spent in the instability strip as the mass decreases. The slope of the PL relationship is thus hardly affected by stars with m $\ge 7 \msol$. In Fig. 2a, comparison is made between the SMC observations and the Z=0.004 models with masses m= 3 - 12 $\msol$. The first crossing unstable models are indicated by open circles but only for m $\le 4 \msol$. Although included in the calculations, the first crossing instability phase for m $\ge 5 \msol$ is statistically insignificant, since the time spent during this phase is more than 300 times smaller than the time spent by a 4 $\msol$ in the instability strip (first crossing and blue loop). We thus predict the existence of Cepheids with periods $P \simle \,1$ day, as the signature of the first instability strip, but the observation of these objects requires larger statistics. The models agree reasonably well with the observed width of the instability strip, although they do not reach the observed blue edge of the EROS data. The faintest objects ($M_V \ge -1.5, {\rm log} P \le 0.4$) seem to indicate a change of slope in the PL relationship, becoming steeper for log P $\simle 0.4$, as discovered in the EROS-2 Cepheid sample and analysed carefully by Bauer (1997) and Bauer et al. (1998). We note that such a trend is observed in the theoretical relation near $\sim$3 $\msol$. This change of slope is thus real and stems from stellar evolution, illustrating the reduction of the He blue loop as mass decreases (cf. \S 2). Note also that the period of the minimum mass undergoing a blue loop ($\sim 3 \msol$) is consistent with the faintest observed SMC Cepheids. Finally the average slope of our P - $\mv$ relationship is $A$ = -2.92, in excellent agreement with the slope derived by S97. Fig. 2b shows the results for the LMC with Z=0.01 and masses from 3.75 to 12 $\msol$. As in Fig. 2a, the models reproduce correctly the observed width of the instability strip and the overall agreement with observations is good. We note that the minimum theoretical ``unstable'' mass undergoing a blue loop $m_{min} \sim 3.75 \msol$ for Z=0.01 does not correspond to the faintest objects observed, which correspond to $\sim 4.25 \msol$. We predict that fundamental pulsators with ${\rm log} P \simle 0.3$ days should be in the first crossing instability. The slope of the PL relationship with a minimum mass of $4.25 \msol$ is $A$ = -2.50, shallower than the one observed in the LMC by S97. In order to test the influence of chemical composition, we have recomputed the whole grid of models with (Z=0.008, Y=0.25) and (Z=0.01, Y=0.28), without any substantial modification of the theoretical slope. % For solar metallicity models, we find also a slope $A$ = -2.55 shallower than the one observed in the Milky Way. A detailed analysis of the possible sources for such discrepancies is under progress. | 98 | 4 | astro-ph9804061_arXiv.txt |
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9804 | astro-ph9804127_arXiv.txt | A revision of Stod\'o\l kiewicz's Monte--Carlo code is used to simulate evolution of star clusters. The new method treats each {\it superstar} as a single star and follows the evolution and motion of all individual stellar objects. The first calculations for isolated, equal--mass $N$--body systems with three--body energy generation according to Spitzer's formulae show good agreement with direct $N$--body calculations for $N=2000$, $4096$ and $10000$ particles. The density, velocity, mass distributions, energy generation, number of binaries etc. follow the $N$--body results. Only the number of escapers is slightly too high compared to $N$--body results and there is no level off anisotropy for advanced post--collapse evolution of Monte--Carlo models as is seen in $N$-- body simulations for $N \leq 2000$. For simulations with $N > 10000$ gravothermal oscillations are clearly visible. The calculations of $N=2000$, $4096$, $10000$, $32000$ and $100000$ models take about $2$, $6$ $20$, $130$ and $2500$ hours, respectively. The Monte--Carlo code is at least $10^5$ times faster than the $N$--body one for $N=32768$ with special--purpose hardware (Makino 1996ab). Thus it becomes possible to run several different models to improve statistical quality of the data and run individual models with $N$ as large as $100000$. The Monte--Carlo scheme can be regarded as a method which lies in the middle between direct $N$--body and Fokker--Planck models and combines most advantages of both methods. | Our knowledge about the stellar content, kinematics, and the influence of the environment on observational features of globular clusters and even richer stellar systems are increasing dramatically (Janes 1991, Djorgovski \& Meylan 1993, Smith \& Brodie 1993, Hut \& Makino 1996, Meylan \& Heggie 1997). First, observations are reaching the point where segregation of mass within globular clusters can be observed directly and quantitatively. Second, observations have revealed that clusters with dense (collapsed) cores are relatively more concentrated to the galactic center than uncollapsed ones. Thus the influences of the environment and mass spectrum are crucial for cluster evolution. Third, observations give clear evidence that post--collapse globular clusters have bluer cores. This suggests strong influence of dynamical interactions between stars on observational properties of globular clusters. Fourth, recent observations show that many different and fascinating types of binaries and binary remnants are present in abundance in globular clusters. Binaries, in addition to being a diagnostic of the evolutionary status of clusters, are directly involved in the physical processes of energy generation, providing the energy source necessary to stop the core collapse and then drive the core expansion. So, to model the evolution of real stellar systems and make meaningful comparison with observation one has to take into account the complex interactions between stellar evolution, stellar dynamics and the environment. Of course all these demands can be fulfilled by direct $N$--body codes (but even the $N$--body method will have trouble with stellar evolution of binary stars). But they are very time--consuming and they need a special--purpose hardware to be run efficiently (Makino 1996ab). Another possibility is to use a code which is very fast and properly reproduces the standard relaxation process and at the same time provides a clear and unambiguous way of introducing all the physical processes which are important during globular cluster evolution. This task might seem unachievable, but actually this kind of code was in use in the past. Monte--Carlo codes, which use a statistical method of solving the Fokker--Planck equation provide all the necessary flexibility. They were developed by Spitzer (1975, and references therein) and H\'enon (1975, and references therein) in the early seventies, and substantially improved by Marchant \& Shapiro (1980, and references therein) and Stod\'o\l kiewicz (1986a, and references therein). Unfortunately, lack of fast computers with sufficient memory at that time and development of the direct Fokker--Planck and gaseous models contribute to the abandonment of this method. But recent developments in computer hardware, speed and memory now make it possible to run a Monte--Carlo code efficiently, even on general--purpose workstations. The great advantages of this method, beside of its simplicity and speed, are connected with the inclusion of anisotropy and with the fact that added realism does not slow it down. The Monte--Carlo method can practically cope as easily as the $N$--body method with internal freedom of single and binary stars and external environment, with one exception, a stellar system must be spherically symmetric. The Monte--Carlo code can have another possible use. Despite the simplified nature of continuum models (Fokker--Planck and gaseous models) they will continue for a while to be the most commonly used codes for stellar dynamical evolution. The Monte--Carlo models can be used to optimise physical free parameters and approximations of continuum models to check their validity as it was done in comparison between small $N$--body simulations and continuum ones (Giersz \& Heggie 1994ab, Giersz \& Spurzem 1994). This procedure should further increase our confidence in results obtained by Fokker--Planck or gaseous simulations. On the other hand the Monte--Carlo techniques can be incorporated in continuum models to describe the stochastic processes of binary formation, energy generation and movement (Spurzem \& Giersz 1996, Giersz \& Spurzem 1997). This, for example, will enable a very detailed investigation of evolution of primordial binaries in evolving background given by an anisotropic gaseous model. The plan of the paper is as follows. In Section 2 a short review of the `old' and `new' Monte--Carlo methods will be presented. In Section 3 the first results of the `new' Monte--Carlo simulation will be presented. And finally in Section 4 the conclusions and future development of the code will be discussed. | \medskip A successful revision of Stod{\'o}{\l}kiewicz's Monte--Carlo code was presented. The updated method treats each {\it superstar} as a single star and follows the evolution and motion of all individual stellar objects. This improvement was possible thanks to the recent developments in computer hardware and computer speed. Two essential changes was added to the original Monte--Carlo code. Firstly, the procedure which deal with problems of radial velocity determination after the system rearrangement (changes of mechanical energy of the stars due to changes of mass distribution) was slightly changed. This assures better energy conservation. Secondly, the new procedure which deals with star escapers was added. This practically resolves the problem with too high escape rate observed in Monte--Carlo simulations. The Monte--Carlo scheme presented here (as previous Monte--Carlo schemes) takes full advantage of the undisputed physical knowledge on the secular evolution of (spherical) star clusters as inferred from continuum model simulations. Additionally it describes in a proper way the graininess of the gravitational field and the stochasticity of the real $N$--body systems. This does not include any additional physical approximations or assumptions which are common in Fokker--Planck and gas models (e.g. conductivity or isotropic distribution function for field stars). From that respect Monte--Carlo scheme can be regarded as a method which lies in the middle between direct $N$--body and Fokker--Planck models and combines most advantages of the both methods. The first calculations for equal--mass $N$--body systems with three--body energy generation according to Spitzer's formulae show good agreement with direct $N$--body calculations for $N=2000$, $4096$ and $10000$ particles. The density, velocity, mass distributions, energy generation, number of binaries etc. follow the $N$--body results. Only the number of escapers is slightly too high compared to $N$--body results (but this can be resolved by the time--dependent shift of the escape rate) and there is no level off anisotropy for advanced post--collapse evolution of Monte--Carlo models as is seen in $N$--body simulations for $N \leq 2000$. For simulations with $N > 10000$ gravothermal oscillations are clearly visible. This is the first unambiguous detection of gravothermal oscillations in Monte--Carlo simulations. Moreover, this is a first unambiguous detection of gravothermal oscillations for stochastic $N$--body system with $N$ as large as $100000$. The speed of the new code makes it possible to run individual models with $N$ as large as $100000$ and also enables, in an unambiguous way, the inclusion of several different physical processes which operate during different stages of evolution of real globular clusters. The new Monte--Carlo code described in this paper is seen as a first step towards realistic models of globular clusters. Several important physical processes have to be included to make the simulations of the stellar systems more realistic. The final code will contain the following physical processes: {\bf (1)} formation of binaries due to dynamical and tidal interactions, {\bf (2)} primordial binaries, {\bf (3)} stellar evolution, {\bf (4)} tidal field of Galaxy and tidal shocks connected with crossing the galactic plane and with large molecular clouds, {\bf (5)} collisions between stars, {\bf (6)} interactions between binaries and stars and between binaries themselves, improving the presently used scattering cross-sections for binary hardening. In the first stage all processes connected with interactions between objects were modelled using analytical cross sections available in the literature. This allowed the code to be tested, and made possible comparison with continuum models. In the next stage interactions between groups of three and four stars will be modelled by numerical integrations of their orbits (the first attempts are tested now). If during the integration the distance between two or more stars becomes smaller than the sum of their radii then a physical collision takes place. This more realistic approach ensures that processes of energy generation (the most important factor in the dynamical evolution of globular clusters) will be modelled more closely. The final stage will be the inclusion of detailed 3--D hydrodynamical modelling of collisions between stars. This will be done by use of Smooth Particle Hydrodynamics (SPH) for a limited number of particles per star (a few hundred). This will allow close comparison between numerical models and observations of real globular clusters. I refer here to observations of various, peculiar objects like blue stragglers and milliseconds pulsars, which can be formed during collisions and encounters between stars. \bigskip \bigskip {\parindent=0pt {\bf Acknowledgments} I would like to thank Douglas C. Heggie and Rainer Spurzem for stimulating discussions, comments and suggestions to a draft version of this paper. I also thank Douglas C. Heggie, who made the $N$--body results for $N = 4096$ particles available. This work was supported in part by the Polish National Committee for Scientific Research under grant 2--P304--009-06.} | 98 | 4 | astro-ph9804127_arXiv.txt |
9804 | astro-ph9804194_arXiv.txt | We present a Hubble Space Telescope (HST) study of the nuclear region of the E4 radio galaxy NGC 7052, which has a nuclear disk of dust and gas. The Second Wide Field and Planetary Camera (WFPC2) was used to obtain {\BB}, {\VV} and~{\II} broad-band images and an {\halnii} narrow-band image. The images yield the stellar surface brightness profile, the optical depth of the dust, and the flux distribution of the ionized gas. The Faint Object Spectrograph (FOS) was used to obtain {\halnii} spectra at six different positions along the major axis, using a $0.26''$ diameter circular aperture. The emission lines yield the rotation curve of the ionized gas and the radial profile of its velocity dispersion. The observed rotation velocity at $r=0.2''$ from the nucleus is $V = 155 \pm 17 \kms$. The Gaussian dispersion of the emission lines increases from $\sigma \approx 70\kms$ at $r=1''$, to $\sigma \approx 400\kms$ on the nucleus. To interpret the gas kinematics we construct axisymmetric models in which the gas and dust reside in a disk in the equatorial plane of the stellar body, and are viewed at an inclination of $70^{\circ}$. It is assumed that the gas moves on circular orbits, with an intrinsic velocity dispersion due to turbulence (or otherwise non-gravitational motion). The latter is required to fit the observed increase in the line widths towards the nucleus, and must reach a value in excess of $500 \kms$ in the central $0.1''$. The circular velocity is calculated from the combined gravitational potential of the stars and a possible nuclear black hole. Models without a black hole predict a rotation curve that is shallower than observed ($V_{\rm pred} = 92\kms$ at $r=0.2''$), and are ruled out at $>99$\% confidence. Models with a black hole of mass $\Mbh = 3.3^{+2.3}_{-1.3} \times 10^8 \Msun$ provide an acceptable fit. The best-fitting model with a black hole adequately reproduces the observed emission line shapes on the nucleus, which have a narrower peak and broader wings than a Gaussian. NGC 7052 can be added to the list of active galaxies for which HST spectra of a nuclear gas disk provide evidence for the presence of a central black hole. The black hole masses inferred for M87, M84, NGC 6251, NGC 4261 and NGC 7052 span a range of a factor $10$, with NGC 7052 falling on the low end. By contrast, the luminosities of these galaxies are identical to within $\sim\!25$\%. Any relation between black hole mass and luminosity, as suggested by independent arguments, must therefore have a scatter of at least a factor $10$. | Astronomers have been searching for direct evidence for the presence of black holes (BHs) in galactic nuclei for more than two decades. Initially, the only constraints on the central mass distributions of galaxies were obtained from ground-based stellar kinematical observations. More recently, the launch of the Hubble Space Telescope (HST) and the subsequent refurbishment in 1993 have provided an important increase in spatial resolution. Combined with new techniques for data analysis and dynamical modeling this has strengthened the stellar kinematical evidence for BHs in several quiescent galaxies (e.g., Kormendy \etal 1996a,b; van der Marel \etal 1997a; Cretton \& van den Bosch 1998; Gebhardt \etal 1998). New tools for the detection of BHs were also developed. HST observations of the rotation velocities of nuclear disks of ionized gas provided accurate BH mass determinations for several active galaxies (e.g., Harms \etal 1994; Ferrarese, Ford \& Jaffe 1996, 1998; Macchetto \etal 1997; Bower \etal 1998), while for other galaxies BHs were detected through VLBI observations of nuclear water maser sources (e.g., Miyoshi \etal 1995). The case for a BH in our own galaxy improved drastically through measurements of stellar proper motions exceeding $1000 \kms$ in the central $0.1\pc$ (Genzel \etal 1997). There are now a total of 10---20 galaxies for which a nuclear dark mass, most likely a BH, has been convincingly detected. The combined results for these galaxies are summarized and reviewed in, e.g., Kormendy \& Richstone (1995), Ford \etal (1998), Ho (1998), Richstone (1998), and van der Marel (1998). This sample is now large enough to study the BH mass distribution in galaxies, which is further constrained by ground-based stellar kinematical observations (Magorrian \etal 1998), HST photometry (van der Marel 1998) and quasar evolution (e.g., Haehnelt \etal 1998). Our understanding remains sketchy, but is consistent with a picture in which a majority of galaxies has BHs, and in which the BH mass $\Mbh$ correlates with the luminosity or mass of the host spheroid. In this paper we present and analyze HST data for the E4 galaxy NGC 7052. This galaxy is a radio source with a core and jet, but no lobes (Morganti \etal 1987). Ground-based optical images show a nuclear dust disk aligned with the major axis of the galaxy (Nieto \etal 1990). The physical properties of this disk were discussed by de Juan, Colina \& Golombek (1996). In a previous paper (van den Bosch \& van der Marel 1995; hereafter Paper~I) we presented ground-based narrow-band imaging and long-slit spectroscopy obtained with the 4.2m William Herschel Telescope (WHT). These observations showed that there is also a rotating nuclear disk of ionized gas in NGC 7052. The gas has a steep central rotation curve, rising to nearly $300 \kms$ at $1''$ from the center. However, the spatial resolution of the spectra was insufficient to convincingly detect a BH, due in part to the relatively large distance of NGC 7052 ($58.7 \Mpc$; i.e., $1'' = 284.6 \pc$). The velocity dispersion of the gas was found to increase from $\sigma \approx 70\kms$ at $1''$ from the center to $\sigma \approx 200\kms$ on the nucleus. We showed that this cannot be the sole result of rotational broadening, which would have predicted double-peaked line profile shapes that are not observed. Instead, the observed central increase in the line width must be at least partly intrinsic. The ground-based kinematics yield an upper limit of $\sim 10^9 \Msun$ on the mass of any possible BH. To improve the constraints on the presence of a central BH we obtained broad- and narrow-band images of NGC 7052 with the Second Wide Field and Planetary Camera (WFPC2) and spectroscopy with the Faint Object Spectrograph (FOS), both in the context of HST project GO-5848. We discuss the imaging and photometric analysis in Section~\ref{s:WF}, and the spectroscopy and kinematical analysis in Section~\ref{s:FOS}. In Section~\ref{s:dyn} we construct dynamical models to interpret the results. With the high spatial resolution of these data we are able to better constrain the nuclear mass distribution, and we find that NGC 7052 has a BH with mass $\Mbh = 3.3^{+2.3}_{-1.3} \times 10^8 \Msun$. We summarize and discuss our findings in Section~\ref{s:disc}. Some observational details are presented in an appendix. We adopt $H_0 = 80 \kms \Mpc^{-1}$ throughout this paper. This does not directly influence the data-model comparison for any of our models, but does set the length, mass and luminosity scales of the models in physical units. Specifically, distances, lengths and masses scale as $H_0^{-1}$, while mass-to-light ratios scale as $H_0$. | \label{s:disc} We have presented HST observations of the nuclear gas and dust disk in the E4 radio galaxy NGC 7052. WFPC2 broad- and narrow-band images were used to constrain the stellar surface brightness profile, the optical depth of the dust, and the flux distribution of the ionized gas. We have built axisymmetric models in which the gas and dust reside in the equatorial plane, and in which the gas moves on circular orbits with an additional velocity dispersion due to turbulence (or otherwise non-gravitational motion). These models were used to interpret the ionized gas kinematics inferred from our new FOS spectra and from existing ground-based spectra. The models fit the observed central rotation gradient only if there is a central BH with mass $\Mbh = 3.3^{+2.3}_{-1.3} \times 10^8 \Msun$. Models without a black hole are ruled out at $>99$\% confidence. The models provide an adequate fit to the available observations with a minimum number of free parameters. The assumptions that we make are similar to those that have been made in HST studies of other galaxies with nuclear gas disks. In several areas our models are in fact more sophisticated than some of the previous work. In particular: we use our multi-colour photometry in order to constrain the central cusp steepness of the stellar mass distribution; we explicitly take into account the contribution of the axisymmetric stellar mass distribution to the circular velocity of the gas, and we do not assume the rotation field to be purely Keplerian; we explicitly model the convolution with the HST/FOS PSF and the binning over the size of the aperture; we model the full line profile shapes, and fit the widths of the emission lines as well as their mean; and we fit Gaussians to the models as we do the data, to properly take into account the fact that Gaussian fits to lines that may be skewed or have broad wings yield biased estimate of the true moments. Still, our models remain only an approximation to the true structure of NGC 7052. In particular: the thickness of the gas disk may not be negligible; the mean motion of the gas may not be circular; and the observed rotation curve may not perfectly reflect the intrinsic rotation curve, because of partial absorption of the emission line flux by dust. The limited sky coverage of the FOS spectra prevents a direct check on whether the gas motions in NGC 7052 are indeed circular. However, several consistency checks are available that may have signaled errors in our assumptions; none did. The stellar mass-to-light ratio and systemic velocity inferred with our models from the nuclear gas kinematics agree with those inferred from stellar kinematical measurements outside the region influenced by dust absorption. The best-fitting model for the gas kinematics reproduces the shapes of the emission lines on the nucleus, despite the fact that these shapes were not included as constraints in the fit. These agreements do not rule out a conspiracy of some sort, but they do make it less likely that the observed gas kinematics are the result of vastly non-circular motion, or have been strongly modified by dust absorption. Models of adiabatic BH growth for the stellar surface brightness cusp provide another successful check: the BH mass implied by these models is fully consistent with that inferred from the gas kinematics. Figure~\ref{f:allBHs} shows a scatter plot of $\Mbh$ versus {\BB}-band spheroid luminosity $L_{B,{\rm sph}}$ for all galaxies with reasonably secure BH mass determinations (adapted from van der Marel 1998, with the addition of NGC 7052; all for $H_0 = 80 \kms \Mpc^{-1}$). There is a trend of increasing $\Mbh$ with increasing $L_{B,{\rm sph}}$, although it remains difficult to rule out that systematic biases play some role in this relation (van der Marel 1998). Besides NGC 7052, the other galaxies for which the BH detections are based on kinematical studies of nuclear gas disks with the HST are M87 (Harms \etal 1994; Macchetto \etal 1997), M84 (Bower \etal 1998), NGC 6251 and NGC 4261 (Ferrarese, Ford \& Jaffe, 1998, 1996). The $\Mbh$ in these galaxies are $3.2 \times 10^9$, $1.4 \times 10^9$, $6.6 \times 10^8$ and $\Mbh = 4.9 \times 10^8 \Msun$, respectively. NGC 7052 falls at the low end of this range. The five galaxies with BH evidence from nuclear gas disks form a very homogeneous set. Each of these galaxies is a radio source and is morphologically classified as an elliptical. The luminosities are identical to within $\sim\!25$\% ($\log L_B$ in the range $10.6$---$10.8$ for all five galaxies). By contrast, the black hole masses span a range of a factor $10$. The results for these galaxies therefore show that any relation between $\Mbh$ and $L_{B,{\rm sph}}$ must have a scatter of at least a factor $10$, even if the comparison is restricted to galaxies of similar type. \placefigure{f:allBHs} | 98 | 4 | astro-ph9804194_arXiv.txt |
9804 | astro-ph9804037_arXiv.txt | Two dimensional realizations of self-consistent models for the ``perfect elliptic disks'' were tested for global stability by gravitational N-body integration. The family of perfect elliptic disk potentials have two isolating integrals; time independent distribution functions $f(E,I_2)$ which self-consistently reproduce the density distribution can be found numerically, using a modified marching scheme to compute the relative contributions of each member in a library of orbits. The possible solutions are not unique: for a given ellipticity, the models can have a range of angular momenta. Here results are presented for cases with minimal angular momentum, hence maximal random motion. As in previous work, N-body realizations were constructed using a modified quiet start technique to place particles on these orbits uniformly in action-angle space, making the initial conditions as smooth as possible. The most elliptical models initially showed bending instabilities; by the end of the run they had become slightly rounder. The most nearly axisymmetric models tended to become more elongated, reminiscent of the radial orbit instability in spherical systems. Between these extremes, there is a range of axial ratios $0.305 \lesssim b/a \lesssim 0.570$ for which the minimum streaming models appear to be stable. | Recent studies of elliptical galaxies and of bulges of spiral galaxies indicate that their figures are likely to be at least slightly triaxial (for reviews see \cite{bin82}; \cite{dzf91}; \cite{bs93}). Most elliptical galaxies appear to be supported at least in part by anisotropies in the velocity distributions rather than by rapid rotation: see, for example, the work on the dwarf elliptical galaxies NGC 147, 185 and 205 by \cite{bpn91} and \cite{hdmp92}. A class of non-rotating potentials, known as the perfect ellipsoids, has been advanced as a possible model for elliptical galaxies (e.g. \cite{dez85}). In these potentials, the mass density is stratified on concentric, similar ellipsoids, and is non-singular in the center. Many of the properties of these potentials can be derived analytically; the orbits all have three isolating integrals, and hence properties such as the time-averaged density distribution can be computed exactly. This simplifies the task of finding {\it self-consistent models\/}: time-steady phase-space distribution functions $f({\bf x}, {\bf v})$ such that the resulting mass density generates the desired gravitational potential. \cite{stat87} and \cite{teub87} have demonstrated that distribution functions for the perfect ellipsoids, and the analogous two-dimensional elliptic disks, can be constructed. Various sub-families of the axisymmetric perfect ellipsoids have been tested for stability (\cite{dzs89}; \cite{ms90}; \cite{mh91}; \cite{rdz91}). Flattened perfect ellipsoids could also be viewed as models for galactic bars. The only analytical bar models are Freeman's (\cite{free66a}, \cite{free66b}, \cite{free66c}) bars, which are based upon a rotating two dimensional harmonic oscillator potential, and the perfect elliptic disk models, which have no figure rotation. \cite{tdz87} showed that in the limit of the needle ($b \rightarrow 0$) the two dimensional perfect elliptic disk is neutrally stable. Prompted by the large streaming velocities seen in barred spiral galaxies, the stability of perfect elliptic disks with maximum angular momentum has already been studied (\cite{slls94}). The roundest disks were unstable to spiral mode formation, the most elongated elliptical models were unstable to bending modes, while the models with axial ratio $b/a$ in the range $0.250 \lesssim b/a \lesssim 0.570$ appeared stable. This paper extends that previous work to the study of a set of low angular momentum perfect elliptical disks. The minimal angular momentum cases allow us to study the ability of internal velocity dispersion to support an elliptic figure, and forms a natural complement to the earlier work as the other bound of the whole class of perfect elliptic disks. As before, we tested for global stability by constructing a discrete, self-consistent model, loading it into an $N$-body integrator and allowing it to evolve. | In this paper and in LS, we have constructed discrete self-consistent representations of the distribution functions of a range of perfect elliptic disks with minimal and maximal angular momentum. These models were then integrated forward in time using an $N$-body integrator to see if they were stable. The nearly axisymmetric and the most elongated models were unstable. The perfect elliptic disks with moderate axial ratios appear to be stable in both the maximum and minimum streaming cases. In the maximum streaming case, the nearly axisymmetric models developed spiral and bar instabilities as expected, since their limiting case, a cold axisymmetric disk, is known to be violently unstable to spiral instabilities. In the minimal angular momentum case, nearly-round disks became more elliptical, in a manner very similar to the radial orbit instability of spherical systems. This is not too surprising, since the velocity distribution is anisotropic, with the radial velocity dispersion being substantially higher than the tangential dispersion, even in the very nearly axisymmetric models. This comes about because of the substantial presence of box and marginal loop orbits in the models. In both angular momentum extremes, the most elongated models developed a bending instability. The similarity in behavior is not very surprising given the decreasing importance of rotational support with increasing ellipticity in these models. \cite{tdz87} have shown that the limiting case of the needle ($b\rightarrow0$) is neutrally stable to bending, while Merritt \& Hernquist (1991)\nocite{mh91} have demonstrated a bending instability in a very prolate (E9) system. It is thus not surprising that the most elliptic models should develop this instability. For the minimum angular momentum family, the instabilities change the shape of the disk towards a more moderate ellipticity. In the nearly axisymmetric disks, as the angular momentum is decreased from a maximum, we expect that the increasing velocity dispersion should help to stabilize against spiral instabilities. It appears likely that there is a stable region for nearly axisymmetric disks with values of Toomre's (\cite{t64}) stability parameter $Q$ which lies between the points $Q \sim 2$ and $Q \sim 3$ where the velocity dispersion has increased to the point of being able to support the disk against the spiral instability (fig \ref{fig-angmom}). As the rotational support becomes negligible and the radial velocity dispersion increases, a radial--orbit instability develops; the disks with lower angular momentum become unstable to elliptical distortions when $T_{\rm radial} / T_{\rm tangential} \gtrsim 1.2$ (after the discussion of \cite{fp84} and \cite{p87} for stability of spherical systems). We expect that there is a range of angular momentum between the two extremes for which the nearly round disks are stable. The moderately elliptical disks with maximum and minimum angular momentum appear to be stable, so we would anticipate that disks of similar ellipticity and intermediate angular momentum will also be stable. The stability of the two-dimensional models with moderate ellipticity gives us hope that the three dimensional perfect ellipsoids of intermediate triaxiality (which is probably the appropriate range for elliptical galaxies \cite{mb81}; \cite{ddc76}; \cite{dzf91}), will also be stable. It is known that some very flattened systems, such as the extreme oblate spheroids constructed from thin short-axis tube orbits (\cite{ms90}), are unstable, but the simple fact that two longer axes are unequal is not likely to be the cause of further trouble. The three dimensional extension of this work will be interesting to see in light of the work of \cite{app92} showing that three dimensional systems with a small amount of rotational streaming are unstable to a tumbling bar instability, both when the models have largely radial orbits and when the orbits are mostly circular. The techniques developed in this work have laid the foundation for investigating the stability of three dimensional perfect ellipsoids, and indeed of any integrable potential. The methods for choosing orbits, whether simulated annealing (as in LS) or the marching scheme of ZHS, can be easily expanded to take account of a variety of possible cost terms related to angular momentum or line of sight velocities (e.g. \cite{rix97}). The procedure of LS for generating a quiet start which minimizes random noise due to particle discreteness, can be carried over to any integrable system. This is potentially most useful in $N$-body studies which attempt to measure the growth rate of instabilities, because the detection of instabilities which are still in the linear regime is limited by particle noise. For example, \cite{s91} found that his linear stability theory was consistent with the results of $N$-body simulations for highly unstable spherical systems, but predicted slow growing instabilities which could not be seen in the simulations because of particle noise. \cite{app90} have constructed an analytic potential--smoothing integration technique which decreases the $\sqrt{N}$ noise associated with binning and softening in $N$-body codes, and permits better examination of the linear growth regime. Their method would also benefit from a quiet start, because the particle discreteness then makes a larger relative contribution to the noise. | 98 | 4 | astro-ph9804037_arXiv.txt |
9804 | astro-ph9804171_arXiv.txt | In this letter we present an idea which reconciles a homogeneous and isotropic Friedmann universe with a fractal distribution of galaxies. We use two observational facts: The flat rotation curves of galaxies and the (still debated) fractal distribution of galaxies with fractal dimension $D=2$. Our idea can also be interpreted as a redefinition of the notion of bias.\\ {\bf Key Words:} Large scale structure of universe-cosmology: theory-cosmology:dark matter galaxies:general | 98 | 4 | astro-ph9804171_arXiv.txt |
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9804 | hep-ph9804444_arXiv.txt | The process of photon splitting $\gamma \to \gamma \gamma$ in a strong magnetic field is investigated both below and above the pair creation threshold. Contrary to the statement by Baier et al., the ``allowed'' channel $\alw$ is shown not to be a comprehensive description of splitting in the strong field because the ``forbidden'' channel $\frb$ is also essential. The partial amplitudes and the splitting probabilities are calculated taking account of the photon dispersion and large radiative corrections near the resonance. \\\\ PACS numbers: 12.20.Ds, 95.30.Cq, 98.70.Rz | 98 | 4 | hep-ph9804444_arXiv.txt |
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9804 | astro-ph9804165_arXiv.txt | Images and longslit, echelle spectra of the \Ha emission from 14 dwarf galaxies and M82 have been used to identify expanding shells of ionized gas. Supershells (radius $>~300$~pc) are found in 12 of the dwarfs. The measured shell sizes and expansion speeds constrain the ages and power requirements of the bubbles. The dynamical age of the larger bubbles is typically about 10~Myr, and ionized shells older than 20~Myr are rare. An energy equivalent to 100 to 10,000 supernova explosions over this period is needed to drive the shock front that sweeps out the cavity. The current star formation rates are high enough to meet these power requirements. Many of the shells will breakthrough the surrounding layer of HI supersonically, but the projected expansion speeds are typically less than the lower limits on the escape velocity. Some of the shell material may permanently escape from a few galaxies such as \n1569. Whether bound to the galaxy or not, these outflows probably play an important role in regulating the star formation rate and are expected to significantly influence the chemical evolution of the galaxies. The shells lift gas out of the disk at rates comparable to, or even greater than, the current galactic star formation rates. They will only displace a substantial fraction of the interstellar gas if their duty cycle is much longer than the rotational period of the disk. | The interplay between massive stars and the interstellar medium (ISM) plays a fundamental role in the formation and evolution of galaxies. In addition to ionizing radiation and newly synthesized elements, massive stars deliver kinetic energy and momentum to the surrounding gas through stellar winds and supernova explosions. Shock waves driven by an ensemble of massive stars may trigger additional star formation and/or sweep the interstellar gas out of the region actively forming stars (Tenorio-Tagle \& Bodenheimer 1988). The gas flows create a turbulent pressure which helps support the weight of the ISM (e.g. McKee 1990) and cavities which apparently enhance the distance ionizing radiation propagates (Hunter \& Gallagher 1997; Martin 1997). This feedback from star formation may have a particularly strong influence on the evolution of low mass galaxies. Owing to their low escape velocity, Larson (1974) suggested that the loss of supernova-heated gas would begin earlier and carry away a larger fraction of their initial mass. This idea was further developed by Dekel \& Silk (1988) who used the supernova feedback to regulate the star formation history of the evolving dwarfs. Their starburst-driven wind models were consistent with the observed mass-metallicity and mass-radius scaling relations of dwarfs when a halo similar to those produced in cold dark matter cosmological simulations was included. Mass loss has subsequently been proposed to explain a number of peculiarities about dwarf galaxies such as their abundance patterns (Marconi,Matteucci, \& Tosi 1994) and rapid evolution at moderate redshifts (Phillipps \& Driver 1995; Babul \& Rees 1992). The ejection of the ISM may not be as easy as previously thought, however. In particular, the rupture of a supershell perpendicular to a galactic disk may vent much of the energy leaving most of the disk intact (DeYoung \& Heckman 1994). Observations of dwarf galaxies reveal an environment conducive to the growth of large bubbles. Their rotation is typically nearly solid body, so shells are not sheared apart; and metallicities are generally sub-solar so cooling times are longer. Indeed, small bubbles permeate the star forming regions of the Magellanic Clouds, and a hierarchy of giant shells ($R < 300$~pc) and supergiant shells ($R \ge 300$~pc) is plainly visible (Davies, Elliot, \& Meaburn 1976; Meaburn 1980; Kennicutt \et 1995). The formation of regions like 30~Doradus, which will evolve into a supergiant shell (Chu \& Kennicutt 1994), may be thought of as the first step in the formation of a galactic outflow. Deep imaging of the ionized gas in other dwarfs yields a plethora of candidate structures for supergiant shells. Indeed, roughly one out of every four high-surface brightness dwarfs exhibit at least one shell and/or filaments (Hunter \et 1993). It is not always obvious, however, which arcs and filaments will show the kinematic signature of an expanding shell (Hunter \& Gallagher 1990). The kinematic evidence is mounting that some shells do breakthrough the ambient neutral gas. In the LMC, for example, the kinematics of many supershells are surprisingly quiescent compared to the giant shells (Hunter 1994). Some of these supershells are believed to be the inner ionized surface of cylindrical HI holes (Meaburn 1979; Meaburn 1980; Hunter 1994), which may have formed as a superbubble blew out perpendicular to the galactic plane. In another Magellanic irregular galaxy, \n4449, the very large HI hole may be associated with a shell that expanded out of the galactic plane (Hunter \& Gallagher 1997). In less luminous galaxies like the blue compact dwarf \n1705, the expansion of the shell around the central starburst is decidedly non-spherical (Meurer \et 1992). The kiloparsec scale, expanding shells in amorphous dwarfs (Marlowe \et 1995) and the faint galaxy IZw18 (Martin 1996) also seem to be elongated in the general direction of the HI minor axis. At issue, however, is whether any of these disk outflows develop into freely flowing winds in which the gas actually escapes from the gravitational potential of the galaxy. Only one member of Marlowe's sample, \n3955, was a strong wind candidate. The most convincing arguments for actual mass ejection are based on the detection of \x emitting gas well above the galactic plane of \n1569 (Heckman \et 1995). The association of a hot bubble with the cavity formed by the expanding network of extended \Ha filaments is reminiscent of the minor axis outflow from M82 (Bland \& Tully 1988; Strickland \et 1996; Shopbell \et 1997), although it is not yet clear whether the dynamics of these two classes of galactic outflows are completely analogous. A more extensive kinematic census is desired to assess the frequency of blowout and the amount of mass loss. This paper presents a catalog of large-scale expanding structures in 14 nearby dwarfs. Although M82 does not strictly meet the sample selection criteria, it was added to the sample to provide a common galaxy between this study and studies of superwinds from more luminous starbursts (Heckman, Armus, \& Miley 1990, hereafter HAM). Galaxies were selected from a volume of radius $d \le 10$~Mpc, right ascension $4 h \le \alpha \le 14~h $, and declination $\delta \ge -35$\deg. An effort was made to pick the galaxies with the most intense star formation over a range in absolute luminosity from $M_B = -13.5$ to $M_B = -18.5$. Each radial velocity field was sampled with deep, high-resolution spectra of the \Ha emission. Additional properties of the galaxies are summarized in Table~\ref{tab:sam}. While similar scale shells are found throughout the sample, the net impact on the host galaxy's evolution may be quite varied. Two factors which largely determine the bubble's fate -- i.e. the distribution of the HI and the gravitational potential -- are not at all uniform across the sample. Hence, the prospects for mass ejection are discussed on a galaxy by galaxy basis. The results have interesting applications for the chemical evolution of dwarf galaxies and the regulation of their star formation rate. This paper is organized as follows. The observations and data reduction are described in \S~\ref{sec:obs}. Section~\ref{sec:results} describes the kinematics of the ionized gas, and \S~\ref{sec:dynamics} discusses the dynamics of the expanding shells. Rotation curves are sketched in \S~\ref{sec:mdot}, and the shell expansion speeds are compared to the escape velocity. Section~\ref{sec:sum} summarizes the main results. | \label{sec:sum} An extensive set of \Ha echellograms and images were used to reconstruct the large-scale kinematics of the ionized gas in 14 dwarf galaxies and M82. Details of the results for individual galaxies are included in their respective subsections of the paper, and \fig~\ref{fig:weaver} provides a concise summary of the shell expansion speeds and sizes. The main results regarding the formation of winds in dwarf galaxies are summarized here. \begin{itemize} \item The formation of supershells must be a common byproduct of massive star formation in dwarf galaxies. Expanding, supergiant ($R > 300$~pc) shells were found in all but two of the galaxies. This sample was drawn from the population of nearby dwarf galaxies with prominent arcs and/or extended filaments in their \Ha emission, and roughly one out of four catalogued dwarf galaxies fits this description (Hunter, Hawley, \& Gallagher 1993). Indeed, the hierarchical growth of these structures probably began in star forming regions akin to 30~Doradus in the Large Magellanic Cloud (Chu \& Kennicut 1994). The most powerful outflows, i.e. \n1569 and M82, were found to be composed of multiple {\em cells} whose walls are probably the interface between polar shells. Star formation in the lowest luminosity galaxies, e.g. IZw18, also generates kiloparsec-scale shells. \item Although many of the expanding complexes survive for 10~Myr, none older than $\sim 20$~Myr were identified. The lack of shells older than this likely reflects their disruption timescale and provides an indirect measure of the scale height of the ISM. Alternatively, the ionization rate of the shells might drop abruptly on this timescale due to changes in the birthrate of massive stars and/or the illumination geometry. Although bright, extraplanar HI shells have not been detected in any of the galaxies in this sample, some HI holes in the LMC (Meaburn 1980) and \n4449 (Hunter \& Gallagher 1997) are probably relics of expanding supershells. The power requirements of the ionized supershells typically exceed the critical power for supersonic disk breakthrough, so a disruption scenario must be favored for them. The sequence of echellograms stepped across the southern lobe of \n1569 constrains their deceleration and shows multiple velocity components at least up to 640~pc above the galactic plane. Future observational work must aim to detect the remains of the hot gas and ruptured shell following blowout. One might speculate that the quiescent filaments in \n5253 could be fragments of a ruptured shell or clumps of infalling material ejected in a previous wind epoch. A better census of the local dwarf population would also be helpful for constraining the duty cycle of the winds. \item Presuming the shells do rupture, the escape of hot, X-ray emitting gas from their interiors seems certain. A diffuse, thermal component of the \x emission has been resolved in three of the galaxies in the sample, but it is only a significant fraction ($> 10\%$) of the interstellar HI mass in M82. In contrast, much of the interstellar gas swept into the warm ionized shells probably remains bound to the galaxy. The structure of the dark matter halo has been measured in several low surface brightness dwarf galaxies with large HI disks (e.g. Carignan \& Beaulieu 1989; Meurer \et 1994) and appears to have a universal structure (Burkert 1995; Navarro, Eke, \& Frenk 1996). Hence, a conservative approach to mass loss is to assume that the bursting dwarfs are embedded in similar halos. These dark halo models often do not provide enough mass to explain the HI rotation speed in the inner galaxy, however. Stars and atomic gas can account for essentially all the dynamical mass inside $R(v_{circ})$ in some galaxies like \n1569, but the fraction varies enormously among the sample members. Although little CO emission is detected from the dwarf galaxies (e.g. Young \et 1996), the large uncertainty in the $H_2$ to CO conversion factor does allow substantial mass contributions from molecular gas (e.g. Maloney \& Black 1988). A dominant disk in the inner few kiloparsecs has two immediate implications. First, the disk mass contributes significantly to the gravitational acceleration of the kiloparsec-scale shells/disk outflows. Second, the observed turnover in the rotation curve may not be revealing much about the core radius of the halo -- a critical parameter for estimating the escape velocity. The maximum circular velocities in the galaxies do generally appear to be comparable to the expansion velocities of the supershells, so the escape velocities are greater than the projected shell speed. The expansion speeds along the minor axis of \n1569 do reach values several times the maximum rotation speed. \item The warm shells alone lift gas out of the disk at rates comparable to or greater than the current galactic star formation rates. The shells transport $10^5$ to $10^6$\msun\ of gas over kiloparsec-scale shells in 10~Myr and leave the sound speed high in a large volume of the ISM. This hydrodynamic mixing will be faster than a diffusion process, so the bubbles will clearly alter the chemical evolution of these galaxies. The timescales for blowout are shorter than the evolutionary timescales of most models for Type~Ia supernova progenitors, so the mass loss begins before much of the iron from the burst has been mixed into the ISM. The composition of the ejected material will depend on the duration of the wind and the composition of the ambient ISM, so their impact on the galactic chemical evolution is interwoven with the galactic star formation history. \item Although the current kinetic energy in the large expanding structures is only comparable to the binding energy of the ISM in \n1569, bubble blowout may still extinguish star formation in particular regions of the other galaxies. If the hot spots percolate across the dwarf irregular galaxies for many rotational periods, then a substantial fraction of the interstellar gas may still be cycled through a halo and/or lost from the galaxy. If the present mass loss rates could be sustained for 6 orbital timescales, for example, then most of the interstellar HI could be removed from the disks of six of the 15 galaxies. This global gas-dynamical feedback will be discussed in the context of the galactic star formation history in a forthcoming paper. \end{itemize} | 98 | 4 | astro-ph9804165_arXiv.txt |
9804 | astro-ph9804023_arXiv.txt | Annihilation of high energy, $\sim 10^{21}$eV, neutrinos on big bang relic neutrinos of $\sim 1$eV mass, clustered in the Galactic halo or in a nearby galaxy cluster halo, has been suggested to generate, through hadronic $Z$ decay, high energy nucleons and photons which may account for the detected flux of $>10^{20}$eV cosmic-rays. We show that the flux of high energy nucleons and photons produced by this process is dominated by annihilation on the uniform, non-clustered, neutrino background, and that the energy generation rate of $\sim 10^{21}$eV neutrinos required to account for the detected flux of $>10^{20}$eV particles is $>10^{48}{\rm erg/ Mpc}^3{\rm yr}$. This energy generation rate, comparable to the total luminosity of the universe, is $\sim4$ orders of magnitude larger than the rate of production of high energy nucleons required to account for the flux of $>10^{19}$eV cosmic-rays. Thus, in order for neutrino annihilation to contribute significantly to the detected flux of $>10^{20}$eV cosmic-rays, the existence of a new class of high-energy neutrino sources, likely unrelated to the sources of $>10^{19}$eV cosmic-rays, must be invoked. | The Fly's Eye \cite{Fly} and AGASA \cite{AGASA} experiments confirmed the existence of a break in the energy spectrum of high energy cosmic rays at $\sim5\times10^{18}{\rm eV}$, for which evidence existed with weaker statistics in the data of previous experiments (Haverah Park, Yakutsk, Sugar, see e.g. \cite{Watson} for a review). Fly's Eye data also strengthen the evidence for a change in primary composition from predominantly heavy nuclei below the break to predominantly light nuclei above the break. These features strongly suggest, when coupled with the lack of anisotropy that would be expected for cosmic-rays (CRs) of Galactic origin, that below $\sim 10^{19}{\rm eV}$ the CRs are mostly heavy ions of Galactic origin, and that an extra-Galactic component of protons dominates above $\sim10^{19}{\rm eV}$. This conclusion is further supported by the fact that the CR energy spectrum is consistent with a cosmological distribution of sources of protons, with injection spectrum $dN/dE\propto E^{-2.2}$ typically expected for Fermi acceleration \cite{cosmology}. In particular, there is evidence for the existence of a Greisen-Zatsepin-Kuzmin (GZK) ``cutoff'', i.e. for the suppression of CR flux above $\sim5\times10^{19}{\rm eV}$ expected due to interaction of protons with the microwave background radiation \cite{GZK}. The evidence for GZK suppression is strengthened by recent AGASA data \cite{AGASA1}. In Fig. 1, the CR spectrum reported by the Fly's Eye and the AGASA experiments \cite{Fly,AGASA1} is compared with the flux expected for a homogeneous cosmological distribution of sources, each generating a power law differential spectrum of high energy protons $dN/dE\propto E^{-2.2}$ (For the model calculation we have used a flat universe with zero cosmological constant, $H_0=75{\rm km}\ {\rm s}^{-1}$, and time independent energy generation rate per comoving unit volume $5\times10^{44}{\rm erg/Mpc}^3{\rm yr}$; The spectrum is insensitive to the cosmological parameters and to source evolution, since most of the cosmic rays arrive from distances $<1$Gpc \cite{cosmology}). The deficit in the number of events detected above $5\times10^{19}{\rm eV}$, compared to a power-law extrapolation of the flux at lower energy, is consistent with that expected due to a cosmological GZK suppression. However, with current data the ``cutoff'' is detected with only $2\sigma$ significance \cite{Tokyo}. The number of events detected above $10^{20}$eV is consistent with that expected based on the cosmological model presented in Fig. 1 (There is an apparent ``gap'' between the highest and second highest energy events detected by the Fly's Eye \cite{gap}. However, assuming that the cosmological model is valid, the probability that such an apparent ``gap'' would be observed is $\sim15\%$ \cite{cosmology,gap}). Nevertheless, the detection of $>10^{20}$eV events does pose challenges to most models of CR production. The high energies rule out most of the acceleration mechanisms so far discussed \cite{Hillas}, and since the distance traveled by such particles must be smaller than $100{\rm Mpc}$ \cite{dist} due to their interaction with the micro-wave background, their arrival directions are inconsistent with the position of astrophysical objects, e.g. jets of powerful radio galaxies \cite{Biermann}, that are likely to produce high energy particles \cite{obj}. Cosmological $\gamma$-ray bursts (GRBs) are likely sources of high-energy CRs, which may account for the CR flux above $10^{19}$eV as well as for the $>10^{20}$eV events \cite{GRBs}. This model recently gained support form GRB afterglow observations \cite{AG}. Other models for the production of ultra-high energy CRs were suggested, where the highest energy events are produced by the decay of super-massive elementary particles related to grand unified theories (see, e.g., \cite{Berezinsky} for recent review). Sources of such particles may be topological defects, left over from a phase transition associated with the symmetry breaking of the grand unified theory \cite{TD}. While no firm prediction exists of the CR flux in these theories, a generic feature of the super-massive particle decay scenarios is that the injection spectrum is much harder than expected for Fermi acceleration. Therefore, this scenario can account only for the flux of $>10^{20}$eV particles, and can not simultaneously explain the origin of $10^{19}$--$10^{20}$eV CRs. It has recently been suggested that annihilation of high energy, $\sim 10^{21}$eV, neutrinos on big bang relic neutrinos of $\sim 1$eV mass, clustered in the Galactic halo or in a nearby galaxy cluster halo, may generate high energy nucleons and photons which may account for the detected flux of $>10^{20}$eV cosmic-rays \cite{Weiler}. The existence of $>10^{21}$eV neutrino flux was argued plausible based on the argument that the mechanism producing the observed high energy, $>10^{19}$eV, particles, most likely protons, also produces charged pions of comparable energy, which subsequently decay to produce neutrinos. It was suggested that the generation spectrum extends well beyond $10^{20}$eV, and that while nucleons produced by a distant source, e.g. a powerful radio galaxy, lose their energy interacting with the micro-wave background, high-energy neutrinos propagate without energy losses and may annihilate on relic neutrinos, producing $>10^{20}$eV nucleons and photons at small distances that would allow them to propagate to Earth. In Sec. 2 we derive the energy density of high energy neutrinos required to account for the observed rate of $>10^{20}$eV air showers. The implications of our results are discussed in Sec. 3. | 98 | 4 | astro-ph9804023_arXiv.txt |
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9804 | hep-ph9804336_arXiv.txt | The broken-symmetry electroweak vacuum is destabilized in the presence of a magnetic field stronger than a critical value. Such magnetic field may be generated in the phase transition and restore the symmetry inside the bubbles. A numerical calculation indicates that the first-order phase transition is delayed but may be completed for a sufficient low value of the Higgs mass unless the magnetic field is extremely high. | It has been found that very strong magnetic fields are capable of destabilizing the electroweak vacuum by forming a vector boson $W^{+}W^{-}$ condensate and restoring the symmetry \cite{amb}. The required field can only be thought to have existed at the very beginning of the universe and one of the possibilities is that it was generated during the electroweak phase transition \cite{baym,cheng}. This primordial field may have been subsequently the seed of the present galactic magnetic field \cite{enq}. One may wonder whether the restoration of symmetry caused by this strong magnetic field can delay the electroweak phase transition. In particular, if it is of first-order the magnetic field might avoid its completion through the bubble mechanism. The simplest way to see why a strong magnetic field can destabilize the electroweak vacuum is to consider the energy of a charged spin-one particle interacting with a uniform magnetic field along the 3-axis \begin{equation} E_{N}^{2}=p_{3}^{2}+m_{0}^{2}+\left( 2N+1\right) eB-geB\quad . \end{equation} For the lowest Landau level $N=0$ if the gyromagnetic factor $g$ is $2$ as occurs in the $W$ case, it is clear that the effective mass will become zero for \begin{equation} B_{c}=\frac{m_{W}^{2}}{e}\simeq 10^{24}G\quad . \label{eq2} \end{equation} This expression is analogous to that of the critical electric field required to create pairs through tunneling. If one wishes to calculate the decay probability of the vacuum, one must evaluate \begin{equation} Z=<0|e^{-iHt}|0>=e^{-it\left( E_{vac}-i\frac{\Gamma }{2}\right) }\quad . \end{equation} In Euclidean metric the one-loop amplitude for a scalar field depends on $ \det \left( -D_{E}^{2}+m^{2}\right) $ with $D_{E\mu }=\partial _{\mu }-ieA_{\mu }$, being $D_{E4}=iD_{0}$. Using the Schwinger proper time method \cite{schw} one obtains \begin{equation} \ln Z=\int_{0}^{\infty }\frac{ds}{s}tre^{-\left( -D_{E}^{2}+m^{2}\right) s}\quad . \label{e2} \end{equation} For constant electromagnetic fields the trace is known to give the vacuum energy density \cite{schm} \begin{equation} \rho =-\int_{0}^{\infty }\frac{ds}{s}\frac{e^{-m^{2}s}}{\left( 4\pi s\right) ^{2}}\left[ \frac{es\sqrt{E^{2}-B^{2}}}{\sin \left( es\sqrt{E^{2}-B^{2}} \right) }-1\right] \quad , \label{e1} \end{equation} where the $-1$ comes from subtracting $\rho \left( A=0\right) $. This integral has a logarithmic divergence for $s=0$ which can be absorbed renormalizing fields and charge \cite{schw}. In Eq.(\ref{e1}) for $E>B$ the integral has poles in the $s$-axis which give origin to an imaginary part corresponding to pair creation. We will be instead interested in the case of $E=0$ and constant magnetic field for which \begin{equation} \rho =-\int_{0}^{\infty }\frac{ds}{s}\frac{e^{-m^{2}s}}{\left( 4\pi s\right) ^{2}}\left[ \frac{esB}{\sinh \left( esB\right) }-1\right] \end{equation} that has no poles. For the spin-1 $W_{\mu }$ case we adopt the view that the only modification to $\ln Z$ is the interaction of spin with magnetic field $2e\mathbf{B}\cdot \mathbf{s}$ in the exponent of Eq.(\ref{e2}). Now the trace must be performed on momentum and spin states where the latter involves this added interaction to give \begin{equation} \ln Z=\int_{0}^{\infty }\frac{ds}{s}\left( e^{-2eBs}+e^{2eBs}+1\right) tre^{-\left( -D_{E}^{2}+m_{W}^{2}\right) s}\quad . \end{equation} Since the remaining trace is equal to the scalar case, the relevant part of the vacuum energy density is \begin{equation} \rho =-\int_{0}^{\infty }\frac{ds}{s}\frac{e^{-m_{W}^{2}s}}{\left( 4\pi s\right) ^{2}}\left[ \frac{esB}{\sinh \left( esB\right) }2\cosh \left( 2eBs\right) \right] \quad . \end{equation} This expression has no poles but diverges for $s\rightarrow \infty $ when $ B>B_{c}=m_{W}^{2}/e$ due to the gyromagnetic factor $2$ of the $W$ boson, which would not occur either for $g=1$ or for the $s=1/2$ case. This divergence is an indication of the vacuum instability for large magnetic field. The decay rate should be evaluated in the more realistic situation of $B$ increasing with time, with the consequent generation of an electric field. In our calculation of the next section we will not take into account the evolution with time of the magnetic field but we will consider that when its value is larger than the critical one in the region of a bubble containing the broken-symmetry vacuum, the bubble will be destroyed. | We have seen that the highest possible magnetic field together with the most favorable law for having homogeneous field in regions of increasing size might have cosmological consequences through the non-completion of the first-order electroweak transition through a bubble mechanism. Therefore the usual electroweak baryogenesis due to bubble expansion would be affected. But one must notice that also the homogeneous increase of $\varphi $ can produce a matter-antimatter asymmetry. This is because there will be a baryonic chemical potential related to the time variation of the $CP$ violating phase $\theta $. The resulting baryonic density will depend on the variation $\Delta \theta $ in the interval when the sphalerons are active due to the smallness of $\varphi $. For a weakly first-order transition, an advantage of this mechanism compared to the bubble one is that the baryonic density would not be erased in the broken phase because here the value of $ \varphi $ is larger due to the delay of the phase transition. One may remind that this problem is also avoided by the baryogenesis in cosmic strings but paying the price of a suppression factor in the active volume. However, it is unlikely that such a strong and large size primordial magnetic field has occurred, and for more acceptable fields the effect would be only a small decrease of the temperature for the completion of a first-order transition. We have studied the influence of the magnetic field on the phase transition using the easiest model, i.e. the standard model and not the MSSM where presumably the first-order phase transition can occur for not too light Higgs mass \cite{care}. The fact that we obtain the completion of the first-order transition without magnetic field for $m_{H}\simeq 70GeV$, not far below the experimental bound is probably due to the definition that it occurs when the bubbles touch each other without taking into account their scattering. But we believe that the general conclusions on the magnitude of the effect do not depend on the details of the used electroweak model. Regarding further developments of this calculation, it would be important to evaluate the vacuum decay rate caused by a time dependent magnetic field in order to consider more carefully its effect on the bubbles instead of taking the simplification of assuming their disappearance as soon as the magnetic field is larger than the critical value. \strut \strut \strut \textbf | 98 | 4 | hep-ph9804336_arXiv.txt |
9804 | astro-ph9804090_arXiv.txt | We have computed finite temperature corrections to the electron-hadron scattering cross sections. These are based upon the renormalized electron mass and the modified density of states due to the presence of a background thermal bath. It is found that the electron-hadron thermal transport scattering cross section can be much larger than the zero temperature one. In the case of electron-neutron transport scattering, we find $\sigma_{ne}(T) / \sigma_{ne} (T=0) \simeq 5$ at $T \simeq 0.1 \, MeV$. \vskip 0.3cm | Finite temperature effects on elementary processes are significant from the point of view of cosmology and astrophysics. The early universe is usually described as a hot gas of particles in nearly thermodynamical equilibrium. Temperature effects enter through the statistical distribution functions. These can renormalize the masses and the wave functions. These renormalized masses and wave functions can then affect scattering processes and decay rates. Several authors \cite{donoghue} have generalized the electron-mass and wave-function renormalization to all temperatures and densities. Dicus et al. \cite{dicus} and independently, Cambier et al. \cite{cps} included the finite temperature effects on weak reaction rates in calculations of standard big-bang nucleosynthesis $(BBN)$. They obtained the corrected light-element abundance and found that the corrections are only of order of a few percent. After that, Saleem \cite{saleem} included the effects of the electron mass shift at finite temperature on $BBN$ and Baier et al. \cite{baier} examined the finite temperature radiative corrections to the weak neutron-proton decay rates. More recently, Fornengo et al. \cite{jwkim} have considered the finite temperature effects on the neutrino decoupling temperature which is important in the evolution of the early universe. In the present work we consider finite temperature corrections to electron-hadron scattering which is important for baryon inhomogeneous cosmologies. Baryon inhomogeneities might have been produced during the cosmological quark-hadron phase transition in the early universe \cite{witten}. If such inhomogeneities were present, then the different diffusion lengths for neutrons and protons could lead to the formation of high-baryon density proton-rich regions and low-baryon density neutron-rich regions. The light element nucleosynthesis yields from such regions can differ significantly from those of standard homogeneous big-bang nucleosynthesis \cite{AHS}. In view of the importance of using light-element yields from the $BBN$ to constrain the baryon-to-photon ratio as well as various cosmological and particle physics theories, such inhomogeneous models must be examined seriously. It is therefore important to quantify the effects of baryon diffusion as accurately as possible. In this regard Applegate, Hogan and Scherrer (AHS) \cite{AHS} have calculated the diffusion rate of baryons through the electron-positron plasma in the early universe. Subsequently, several authors used their results in calculations of inhomogeneous $BBN$ \cite{applegate,mathews}. In $AHS$ it was suggested that the diffusion coefficients could be derived from the mobility of the heavy particles, and that the mobility is determined from the distribution functions of the background plasma and the transport cross section. It is important, therefore, to carefully quantify the values of the distribution functions and the transport cross sections. However, in all previous baryon diffusion coefficient calculations, vacuum transport scattering cross sections have been used. Therefore, in order to estimate the baryon diffusion coefficients more precisely, in the present work we take into account the finite temperature effects in the calculation of baryon diffusion coefficients at temperatures $\lsim \; MeV$. Specifically, we calculate the transport scattering cross section of elastic electron-hadron scattering at finite temperature. Here we shall treat hadrons as particles which have an internal structure and an anomalous magnetic moment (although we do not have a good field theory for the magnetic moments of protons or neutrons at finite temperature). Also, we assume that their internal structure is not affected by finite temperature since their mass is more than about 1000 times the temperatures of interest. The plan of the paper is as follows. In Section \ref{sec:finite}, we discuss how to include finite temperature effects in the calculation. In particular, we will briefly discuss the effective mass of an electron in the MeV temperature range. In Section \ref{sec:scattering}, we evaluate the electron-hadron transport scattering cross section at finite temperature. Finally, we summarize our results and discuss some astrophysical applications. We shall employ units in which $\hbar = k_B = c = 1$, except when specific units must be attached to a result. | We have calculated temperature-dependent electron-hadron transport cross sections. These are important, for example, in the calculation of baryon diffusion coefficients at finite temperature. The major motivation here has been to investigate whether finite temperature effects can significantly change the baryon transport cross section $\sigma_t$. In this work, we have treated hadrons as particles which have an internal structure and an anomalous magnetic moment. Also, we have assumed that their internal structure is not affected by the finite temperature since their mass is more than about 1000 times the temperature of interest. Two major features of the finite temperature effects on the light particles have been included in the calculation: (1) finite temperature Dirac spinors which are recast into the form of an effective electron mass ; (2) finite temperature modifications to the phase space distribution of the electrons. We find that, for $m_0 < T$, both $\sigma_{ne}(T)$ and $\sigma_{pe}(T)$ approach the ultrarelativistic limit (where the electron mass can be ignored). In the case of electron-proton scattering, we have compared it with the Coulomb scattering cross section at finite temperature. In particular, for the case of electron-neutron transport scattering, we find $\sigma_{ne}(T) / \sigma_{ne} (T=0) \simeq 5$ at $T \simeq 0.1 \, MeV$. In conclusion, the baryon diffusion coefficients which affect baryon inhomogeneities during big-bang nucleosynthesis could be changed significantly by our temperature dependent electron-hadron transport cross sections. Up to the time of weak decoupling ($T \simeq 1 \, MeV$) there is little change in the cross sections. However, during the epoch of nucleosynthesis ($T \lsim 0.2 MeV$) when baryon diffusion is most important, the transport cross sections increase as the temperature decreases. On the other hand, baryon diffusion at low temperature is strongly affected by proton-neutron scattering for which these finite temperature effects are insignificant. Clearly, a study of the effects of these new cross sections on the baryon diffusion coefficients and inhomogeneous primordial nucleosynthesis is desired. These will be the subject of a subsequent paper. \vspace{0.5cm} {\bf Acknowledgements.} \noindent The author (ISS) would like to thank P. Marronetti and Prof. S. Rhie for their helpful comments. ISS also acknowledges the support by the Korea Research Foundation (KRF) for the Post-Doctoral Fellowship at University of Notre Dame. This work supported in part by DOE Nuclear Theory Grant DE-FG02-95ER40934. | 98 | 4 | astro-ph9804090_arXiv.txt |
9804 | astro-ph9804059_arXiv.txt | To determine the magnification of an extended source caused by gravitational lensing one has to perform a two-dimensional integral over point-source magnifications in general. Since the point-source magnification jumps to an infinite value on caustics, special care is required. For a uniformly bright source, it has been shown earlier that the calculation simplifies if one determines the magnification from the area of the images of the extended source by applying Green's theorem so that one ends up with a one-dimensional integration over the image boundaries. This approach is discussed here in detail, and it is shown that it can be used to yield a robust and efficient method also for limb-darkened sources. It is also shown that the centroid shift can be calculated in a similar way. | For fitting light curves for the ongoing microlensing events, there is a need for robust and efficient methods for calculating the magnification of extended sources, which are not limited to point-lenses. Among the observed events, the presence of binary lenses is a reality (Dominik \& Hirshfeld~\cite{MLMC1let},\cite{MLMC1}; Udalski et al.~\cite{OGLE7}; Alard et al.~\cite{DUO2}; Bennett et al.~\cite{MLMC9}), and planetary events involve a special case of a binary lens. In addition, for some configurations, the light curve for a limb-darkened source will differ significantly from that of a uniformly bright source. The limb-darkening effect has recently been observed in the galactic microlensing event MACHO 97-BLG-28, which involves both an extended source and a binary lens, by the PLANET collaboration (Albrow et al.~1998a,b\nocite{M28Planet1}\nocite{M28Planet2}); the fitting has been done by myself using the algorithm described in this letter. If one wants to integrate the point-source magnification in two dimensions one has to take special care of the position of the caustics, where the point-source magnification becomes infinite. While this integration can be performed easily for a point-mass lens (e.g. Schneider et al.~\cite{SEF}, p. 313; Witt \& Mao~\cite{WM}; Sahu~\cite{Sahu}; Dominik~\cite{DoDiss}), this would be a difficult task for a general lens (e.g. a binary lens), especially at a cusp singularity. In contrast, the area of the images of the extended source and therefore its magnification remains continuous when the source hits a caustic. The determination of the extended source magnification from the boundaries of the image areas has been used to analyze the images of background galaxies behind a cluster of galaxies (Dominik~\cite{DoDipl}). The image boundaries can be obtained with a contour plot of an implicit function describing the source boundary in the lens plane (Schramm \& Kayser~\cite{SK}). This method has been expanded with routines for correcting, testing and finally analyzing the contour line in order to produce an efficient and safe algorithm (Dominik~\cite{DoAstro}). In that paper, it is noted that it is easy to analyze the images from the contour line data, and an example is given, where quantities such as the area, width, length and curvature of the image have been determined. Concerning microlensing light curves, it has been noted by Bennett \& Rhie (\cite{BenRhie}) that it is advantageous to integrate in the lens plane rather than in the source plane to determine the magnification of an extended source. For uniformly bright sources, Gould \& Gaucherel (\cite{GouGau}) proposed applying Green's theorem so that only one integration along the image boundary must be performed rather than two over the image area. This approach is identical to that used earlier (Dominik~\cite{DoDipl},~\cite{DoAstro}). The contour plot method is the most convenient way to obtain data points on the image boundary from which the area can be calculated. In Sect.~2, this general approach is described, Sect.~3 gives details for a uniformly bright source, and Sect.~4 shows how this approach can also be used for limb-darkened sources, in which case an easy-to-perform two-dimensional integration remains. In Sect.~5, the calculation of the centroid shift is discussed. | 98 | 4 | astro-ph9804059_arXiv.txt |
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9804 | astro-ph9804329_arXiv.txt | s{ Knowledge of the constants that describe the current cosmological world model, $H_0$, $t_0$, and the three $\Omega$s is central to physical cosmology. Although there is a vast range of suggested tests and existing constraints much of the recent discussion of cosmological constants involves three time variable photospheres: Cepheids, SNIa, and the CMB last scattering surface. These concluding remarks for the Moriond XXXIII meeting are made at a time when many of the established methods have made careful, interesting, statements about the values of various cosmological constants based on data of small random errors. The flood of new data over the next few years will lead to a satisfying increase in the precision of both direct and model dependent estimates of the main cosmological parameters. } | Commenting on the progress being made in estimating the values of the cosmological parameters has perils not unlike those of critiquing a great artwork, perhaps an opera, while it is being staged for the first time. The problem is that there is no working consensus for the values of the cosmological constants. There are rather wide compromise ranges which accommodate most of the derived values. However the values of the cosmological constants at one end of a compromise range are completely incompatible with those at the other, both in physical meaning and in stated measurement error. An aspect of cosmological parameter determination that is fascinating for any interested scientist is that it continues to reward an integrated view of the entire subject from trigonometric parallaxes to photometric response functions to radio interferometric mapping to X-ray plasma analysis to the outer limits of particle theory, to name but a few. Martin Rees has described Cosmology as the ``Grandest of the Environmental Sciences'' \cite{rees} which serves to remind us of the difficulty of relating the observations of the universe to the simple cosmological models of interest. Although the FRW model and an interest in its parameters have been around for a long time, it was not until the 1980's when development of solid state detectors of very low noise and very high quantum efficiency allowed virtually every waveband used by astronomy to greatly increase both the quality and the abundance of data. The objects now examined range from nearby white dwarfs to galaxies at redshifts beyond 5, to precise measurements of the CMB radiation all over the sky. The following discussion discusses the broad comprise ranges for the basic cosmological parameters. In most cases the data have small random errors. The problems for all methods come in calibration and model uncertainties. For instance Cepheids distances have random errors of about 1\%, nevertheless the total error in the Hubble constant is generally quoted to be at least 10 times larger as a result of potential calibration and systematic errors. Even those apparently large error budgets are hard won from vast efforts to control, measure and remove systematic errors. The true error ranges of these are continuing to shrink significantly and steadily, such that sometime in the next decade our understanding will undergo a phase transition, hopefully a crystallization not a meltdown. | The current situation is precisely why the measurement of the cosmological parameters is the primary activity of many astronomers and astrophysicists. The exciting likelihood is that most of the major cosmological constants will be known to a satisfying degree of on the time scale of a decade. The Moriond meetings provide an ideal format for frank discussion of extremely controversial cosmological issues. I thank the organizers for the splendid job they did in managing to bring us all together and providing a never ending flow of food and stimulus for mind and body. | 98 | 4 | astro-ph9804329_arXiv.txt |
9804 | astro-ph9804144_arXiv.txt | The proton burning process $p+p\rightarrow d +e^+ +\nu_e$, important for the stellar evolution of main-sequence stars of mass equal to or less than that of the Sun, is computed in effective field theory using chiral perturbation expansion to the next-to-next-to-leading chiral order. This represents a model-independent calculation consistent with low-energy effective theory of QCD comparable in accuracy to the radiative $np$ capture at thermal energy previously calculated by first using very accurate two-nucleon wavefunctions backed up by an effective field theory technique with a finite cut-off. The result obtained thereby is found to support within theoretical uncertainties the previous calculation of the same process by Bahcall and his co-workers. | The proton fusion reaction \be p + p \rightarrow d + e^+ + \nu_e \label{pp}\ee which plays an important role for stellar evolution and -- as the dominant neutrino source -- for the solar-neutrino problem, has quite a long history of investigation. Indeed the reaction rate of this process (hereafter called the $pp$ rate) was first calculated by Bethe and Critchfield (\cite{bethe38}). Salpeter (\cite{salpeter}) recalculated the $pp$ rate using the {\em effective range approximation} and argued that the relevant nuclear matrix element squared could be estimated with an accuracy of the $\sim$5 \% level. (The $pp$ rate itself was subject to much larger uncertainty, $\sim$20 \%, because of the limited precision with which the Fermi coupling constant was known at that time.) Bahcall and May (\cite{bahcall69}) examined the dependence of the $pp$ rate on explicit forms of the two-nucleon wavefunctions generated by two-parameter nuclear potentials of various forms adjusted so as to reproduce the scattering length and effective range (for the $pp$ channel) and the low-energy properties of the deuteron (for the $np$ channel). The $pp$ rate was found to vary by $\sim 1.5$ \% corresponding to the changes in the deuteron wavefunction, and by $\sim 1.2$ \% due to the change in the $pp$ wavefunction. The most updated work along this line was done by Kamionkowski and Bahcall (\cite{bahcall94}) employing deuteron wavefunctions obtained from much more accurate potentials such as the Argonne $v_{14}$, $v_{18}$, Urbana $v_{14}$, super-soft-core (SSC) and Reid soft-core potentials. Changes in the $pp$ rates arising from the different potentials were found to be $\sim 1\ \%$. Thus it seems that the presently available calculated $pp$ rate is robust and needs no further scrutiny, the famous solar neutrino problem remaining unresolved from this angle and hence persisting as one of the outstanding unsolved problems in astrophysics (Bahcall \cite{unsolved}). There are however two reasons for revisiting this issue. One is that while the calculated $pp$ rate seems to have converged to a ``canonical" value given in (Kamionkowski \& Bahcall \cite{bahcall94}, hereafter KB), there lingers the unsettling feeling that the strong interaction involved in nuclear physics of the two-nucleon systems is infested with uncontrollable uncertainties associated with model dependence in the treatment, making it difficult to assess the accuracy achieved. Thus it is not unexpected that this canonical value will be -- as has been in the past -- challenged. Indeed it has recently been argued by Ivanov et al. (\cite{ivanov}) that the nonrelativistic potential models used in the previous works could be seriously in error. They show that in their version of a {\it relativistic field theory model}, the $pp$ rate comes out to be as big as $2.9$ times the previous estimates.\footnote{They use a procedure that seems to disagree with other physical properties of low-energy $pp$ systems, as was pointed out by Bahcall and Kamionkowski (1997). It has also been pointed out by Degl'Innocenti et al. (\cite{Degl}) that such a large deviation from the value used by KB would be inconsistent with helioseismology in the Sun.} Should their new result turn out to be correct, it would have profound consequences on theories of stellar evolution in general and on the solar neutrino problem in particular. In a nutshell, the issue comes down to whether or not a more general framework such as relativistic field theory would invalidate the calculation made in the traditional nonrelativistic potential models. The claim of the authors in (Ivanov et al. \cite{ivanov}) is that it indeed does. Our aim is to address this issue using a low-energy effective field theory of QCD that has found a quantitative success in other nuclear processes. The second reason is really more theoretical, independent of the above important astrophysical issue. Along with the thermal $np$ capture, the proton fusion process is the simplest nuclear process amenable to an accurate calculation -- something rare in hadronic physics -- and it is of interest on its own to test how well a calculation faithful to a ``first-principle approach" can tackle this problem. In particular, we are interested in checking how accurately the effective field theory approach, found to be stunningly successful for the $np$ capture $n+p\rightarrow d+\gamma$, low-energy NN scattering and static properties of the deuteron (Park, Min, \& Rho \cite{pmr_PRL}; Park et al. \cite{pkmr}), fares with the proton fusion problem, the weak interaction sector of the Standard Model. The strategy we shall adopt here is quite close to that used for the $np$ capture process (Park et al. \cite{pmr_PRL}). We shall use chiral perturbation theory to the next-to-next-to-leading order (NNLO) in chiral counting; as defined precisely below, this corresponds to $O(Q^3)$ relative to the leading-order term. This is roughly the same order as considered for the $np$ capture. However the relative importance of terms of various chiral orders is somewhat different here. As we explain later, in the present case, the corrections to the leading order are not suppressed by what is called the ``chiral filter" \footnote{The chiral filter phenomenon is explained in detail in (Park, Min, \& Rho, \cite{pmr_Report}). Crudely stated, it refers to the general feature that whenever one soft-pion exchange is allowed by kinematics and selection rules, it should give a dominant contribution with higher-order (or shorter-range) terms strongly suppressed. A corollary to this is that whenever one soft-pion exchange is not allowed, all higher-order terms {\it can be} important, making chiral perturbation calculation generically less powerful.} and so the accuracy with which these can be calculated is not as good as in the $np$ capture case. Even so, using the argument developed in (Park et al. \cite{pkmr}), we shall suggest that the procedure used here provides a model-independent result in the same sense as in (Park et al. \cite{pmr_PRL}, Park et al. \cite{pkmr}). To streamline the presentation, we first give our result and then discuss (as briefly as possible) how we arrive at it in the rest of the paper. Apart from the meson-exchange contributions which are of order of ${\cal O}(Q^3)$ and which for the reason mentioned above and further stressed in our concluding section, are the main uncertainty to the order considered, our chiral perturbation theory result in terms of the reduced matrix element $\Lambda$ defined in (Bahcall \& May \cite{bahcall69}) is \bea \Lambda_{\chi PT}^2 = (1 \pm 0.003) \times 6.93 \label{ChPT} \eea where the uncertainty is due to experimental errors.\footnote{ Our theoretical uncertainty is, if very conservatively estimated, about $0.1\ \%$.} The above result is to be compared with the value obtained by Kamionkowski and Bahcall (KB) \be \Lambda_{\rm KB}^2 = \left(1 \pm 0.002^{+0.014}_{-0.009}\right) \times 6.92\, . \label{bahcallS}\ee As we shall explain, there are some differences in details between our calculational framework and that of KB, but our final numerical result is in good agreement with that of KB and disagrees with that of Ivanov et al. (\cite{ivanov}). The paper is organized as follows. In Section \ref{2}, our strategy for carrying out a chiral perturbation calculation for two-nucleon systems is outlined. Our approach here is similar to the one used in the previous calculation of the $np$ capture process. We shall sketch a justification of this approach from the standpoint of low-energy effective field theory of QCD (with concrete supporting evidence summarized in Subsection \ref{cutoff}). Section \ref{chiralcounting} describes chiral counting of the terms appearing in the relevant weak current. In Section \ref{wavefunction}, the wavefunctions for the initial $pp$ state and the final $d$ state are specified. Our numerical results are given in Section \ref{numbers}, and a brief discussion including a comment on the main uncertainty in the calculation is given in Section \ref{discussion}. | \label{discussion} Following the procedure of chiral perturbation theory proven to be highly successful for the thermal $np$ capture process, we have calculated the $pp$ fusion rate to $\calO (Q^3)$ in chiral counting relative to the leading single-particle Gamow-Teller matrix element. {\it To the order considered}, the error involved in the calculation is small, $\lesssim 1$ \%. This result, given a justification from a cut-off effective field theory of low-energy QCD as in the case of the $np$ capture (Park et al. \cite{pkmr}), supports the canonical value of Bahcall et al. and does not support the ``relativistic field theory model" result of Ivanov et al. (\cite{ivanov}). The main caveat in this calculation is in the meson-exchange contribution which comes out to be about 4 \% when calculated to the chiral order $\calO (Q^3)$. At the next order, $\calO (Q^4)$, loops and higher-order counter terms enter, so that there is no reason to believe that they are negligible compared with the $\calO (Q^3)$ tree contributions. (For instance, it could be lowered to about $1\sim 2$ \% instead of $\sim 4$ \% found here). This aspect is different from the case of the $np$ capture where the chiral filter mechanism assures the dominance of the tree-order pion exchange-current contribution. Here the absence of the chiral filter phenomenon {\it can} allow higher-order (loop) terms to figure equally importantly as the tree-order terms. {}From this viewpoint it is not surprising that a model calculation of the terms of $\calO (Q^4)$ and higher based on the vector-meson exchange and form factors (Bargholtz \cite{bar79}) indicates that there can be a considerable suppression of the tree-order correction. Although such a reduction in the exchange-current contribution seems to go in the right direction for beta decays of higher-mass nuclei (Carlson \cite{carlson}), it is probably unsafe to import the result of such model calculations into our work. To go to $\calO (Q^4)$ or above in our theory at which heavy mesons and form factors come in, a large number of Feynman graphs of the same chiral order have to be computed on the same footing to assure chiral symmetry, and such a calculation has not been done yet. In the absence of consistent calculations, our attitude is that we should not attach any error estimates on terms not accounted for to the chiral order computed. Calculating the higher-order terms will be left for a future exercise. | 98 | 4 | astro-ph9804144_arXiv.txt |
9804 | astro-ph9804002_arXiv.txt | Time-resolved eclipse spectroscopy of the nova-like variable UX~UMa obtained with the HST/FOS on 1994 August and November is analyzed with eclipse mapping techniques to produce spatially resolved spectra of its accretion disc and gas stream as a function of distance from disc centre. The inner accretion disc is characterized by a blue continuum filled with absorption bands and lines which cross over to emission with increasing disc radius, similar to that reported by Rutten et~al (1994) at optical wavelengths. The comparison of spatially resolved spectra at different azimuths reveals a significant asymmetry in the disc emission at UV wavelengths, with the disc side closest to the secondary star showing pronounced absorption by an `iron curtain' and a Balmer jump in absorption. These results suggest the existence of an absorbing ring of cold gas whose density and/or vertical scale increase with disc radius. The spectrum of the infalling gas stream is noticeably different from the disc spectrum at the same radius suggesting that gas overflows through the impact point at disc rim and continues along the stream trajectory, producing distinct emission down to $0.1\; R_{L1}$. The spectrum of the uneclipsed light shows prominent emission lines of Ly$\alpha$, N\,{\sc V} $\lambda 1241$, Si\,{\sc IV} $\lambda 1400$, C\,{\sc IV} $\lambda 1550$, He\,{\sc II} $\lambda 1640$, and Mg\,{\sc II} $\lambda 2800$, and a UV continuum rising towards longer wavelengths. The Balmer jump appears clearly in emission indicating that the uneclipsed light has an important contribution from optically thin gas. The lines and optically thin continuum emission are most probably emitted in a vertically extended disc chromosphere + wind. The radial temperature profiles of the continuum maps are well described by a steady-state disc model in the inner and intermediate disc regions ($R \leq 0.3 R_{L1}$). There is evidence of an increase in the mass accretion rate from August to November (from \.{M}$= 10^{-8.3\pm 0.1} \;{\rm to}\; 10^{-8.1\pm 0.1}\; M_{\odot} \; yr^{-1}$), in accordance with the observed increase in brightness. Since the UX\,UMa disc seems to be in a high mass accretion, high-viscosity regime in both epochs, this result suggests that the mass transfer rate of UX~UMa varies substantially ($\simeq 50$ per cent) on time scales of a few months. It is suggested that the reason for the discrepancies between the prediction of the standard disc model and observations is not an inadequate treatment of radiative transfer in the disc atmosphere, but rather the presence of additional important sources of light in the system besides the accretion disc (e.g., optically thin continuum emission from the disc wind and possible absorption by circumstellar cool gas). | Accretion discs are an important phenomenon in astrophysics, invoked to solve a wide range of astrophysical problems ranging from planetary formation to quasar energetics (Frank, King \& Raine 1992). Although considerable effort in both observation and theory has been invested over the past decade, the structure and underlying physics of accretion discs remains poorly understood. Major unsolved problems include the nature of the viscosity mechanism -- responsible for the spiraling inward of the disc material -- (the angular momentum problem), the fate of the kinetic energy expended at the inner edge of the accretion disc (the boundary layer problem), the vertical structure of the disc (the Balmer decrement problem), and the outflow of matter in connection with a disc wind (possibly a solution to, or at least an element of, the boundary layer problem). Progresses in solving these issues has been hampered because most of the previous observational constraints provided only the spectrum of the total light from the disc. A better understanding of the physics of accretion discs requires spatially-resolved studies. Cataclysmic Variables (CVs) are mass-exchanging binary systems containing a white dwarf and a late-type star (Warner 1995). If the white dwarf is not strongly magnetized ($B < 10^{6}$~G) an accretion disc is formed. Accretion discs in non-magnetic CVs cover a range of accretion rates and viscosity states. For example, {\em dwarf novae} undergo large outbursts ($\Delta m = 3-5$~mag, typical duration of 5-10 days) which reflects changes in the structure of the discs -- from a cool, optically thin, low viscosity state to a hot, optically thick, high viscosity state -- and which are usually parameterized as a large change in the mass accretion rate ( \.{M}$= 10^{-11} \; M_\odot \; yr^{-1} \mapsto 10^{-9} \; M_\odot \; yr^{-1}$. See, e.g. Pringle, Verbunt \& Wade 1986). On the other hand, {\em nova-like} variables seem to be permanently in a high viscosity state, presumably as a result of the fact that the accretion rate is always high. Because the nature of the other constituents in these systems -- the white dwarf and the normal star -- are reasonably well understood, and because orbital variations often provide considerable insight into the system geometry, non-magnetic CVs are the ideal laboratories for understanding accretion discs. Eclipsing systems are particularly useful since the occultation of the accretion disc by the late-type star provides information about the disc's spatial structure through the eclipse shape. The eclipse mapping method (Horne 1985, 1993; Rutten, van Paradijs \& Tinbergen 1992; Baptista \& Steiner 1993) assembles the information contained in the eclipse shape into a map of the disc surface brightness distribution. When applied to time-resolved spectroscopy through eclipses this technique delivers the spectrum of the disc at any position on its surface. Information on the radial dependence of the temperature and vertical temperature gradients (for optically thick regions), or temperature, surface density and optical depth (where the disc is optically thin) can be obtained by comparing such spectra with the predictions of models of the vertical disc structure. The spatial structure of the emission-line regions over the disc can be similarly mapped from data of high spectral resolution. Furthermore, by studying the time-variations in the structure of accretion discs of dwarf novae undergoing outbursts it may be possible to uncover the nature of the (so far unknown) viscosity mechanism which drives accretion discs. UX~UMa is a well known, bright ($V \simeq 12.5$) eclipsing nova-like variable with an orbital period of 4.72~hr. Eclipse mapping in broad-bands (Horne 1983; Rutten et~al. 1992) shows that its accretion disc is optically thick and is close to a steady state at a mass accretion rate of $\simeq 10^{-8} \; M_\odot \; yr^{-1}$. The broad-band mapping was extended to spectrally-resolved mapping in the optical range by Rutten et al. (1993, 1994). Their results show that the continuum becomes fainter and redder with disc radius -- reflecting a radial temperature gradient -- and reveal that the Balmer lines are seen in absorption in the inner disc but in emission in the outer disc. Baptista et al. (1995) performed a similar study using HST data in narrow spectral windows about the C\,{\sc IV} 1550 and He\,{\sc II} 1640 line regions. This study showed that the UV continuum reasonably follows the $T \propto R^{-3/4}$ law for steady mass accretion, confirming the results from the optical analysis. The C\,{\sc IV} and He\,{\sc II} line profiles are dominated by emission from the disc wind. Spatially-resolved spectra reveal that these lines appear as narrow absorption features at disc centre and change with increasing radius to broad emission in the outer disc regions besides showing large uneclipsed components. This behaviour is similar to that found for the Balmer lines and suggests that these optical lines may also have a wind component. In this paper, we report on the ultraviolet (UV) and optical mapping of the accretion disc and gas stream of UX~UMa, based on observations made with the {\it Faint Object Spectrograph} (FOS) on the {\it Hubble Space Telescope} (HST). The reader is referred to Knigge et~al. (1998a) for an initial description of these observations, with emphasis on the spectral properties of the integrated spectra of the accretion disc, the bright spot and the uneclipsed light. Sect.\,\ref{observa} describes the data and its reduction. The extraction of narrow-band light curves and their analysis with eclipse mapping techniques are described in Sect.\,\ref{analise}. Sect.\,\ref{results} presents and discusses spatially resolved spectra of the accretion disc and the gas stream region as a function of distance from disc centre, the spectrum of the uneclipsed light, and the radial temperature distribution in the ultraviolet. The possible influence of the assumed eclipse geometry on the results is addressed in Sect.\,\ref{geo}. Sect.\,\ref{discuss} discusses the implications of the results in the context of disc atmosphere models. The results are summarized in Sect.\,\ref{conclusao}. | \label{discuss} Knigge et~al. (1998a) show that disc models constructed as ensembles of stellar atmospheres provide poor descriptions of the observed integrated spectrum of UX~UMa. The disc model spectra are too blue at ultraviolet wavelengths and overpredict the magnitude of the Balmer jump. These problems are not new. The difficulties in fitting integrated spectra of nova-likes and dwarf nova in outburst with disc model spectra have a long history (e.g., Wade 1984, 1988; La Dous 1989; Long et~al. 1991, 1994; Knigge et~al. 1997). In discussing possible explanations for these problems, Knigge et~al. (1998a) postulated the presence of a significant amount of optically thin material in the system in order to reconcile the disc models with the observed spectrum. Our spatially resolved study confirms their suggestion by revealing that the integrated spectrum of UX~UMa has indeed a substantial contribution from optically thin emission, most probably associated to the uneclipsed parts of the disc chromosphere + wind. A calculation by Knigge et~al. (1998a) indicated that the addition of an optically thin component with $T= 3 \times 10^4$\,K, $n_H= 5 \times 10^{12}\; {\rm cm}^{-3}$, and vertical extension $H= 9.7 \times 10^9$ cm would be enough to bring the combined disc model plus optically thin emission into good agreement with the observed PRISM spectrum. The predicted fluxes of their optically thin component raises from $\simeq 2$ mJy at 2000 \AA\ to $\simeq 5.2$ mJy at 3600 \AA, being at the level of $\simeq 3$ mJy at 4500 \AA\ -- in good accordance with the fluxes of the uneclipsed component in Fig.\,\ref{ffig6}. For this case, their inferred mass accretion rate is $5 \times 10^{17}\; g\,s^{-1}$ or $10^{-8.1}\; M_\odot\: yr^{-1}$, in excellent agreement with our result. Thus, the reason for the discrepancies between the prediction of the standard disc model and observations is not an inadequate treatment of radiative transfer in the disc atmosphere (or standard models of vertical structure), but rather the presence of additional important sources of light in the system besides the accretion disc (e.g., optically thin continuum emission from the disc wind and possible absorption by circumstellar cool gas). Following the same line of reasoning, if disc winds are a common characteristic of all non-magnetic, high state cataclysmic variables, one might expect their disc chromospheres to contribute a non-negligible amount of optically thin emission to the total light of the system. Under this hypothesis, the discrepancy between disc models and the integrated spectrum observed in other non-magnetic, high state systems may be removed by the inclusion of a proper optically thin component to their total light. These results underscore the importance of spatially resolved studies in disentangling the different components of the integrated spectra of cataclysmic variables. In this particular case, it helped to clarify the situation regarding the apparent discrepancy between disc atmospheres models and the observed spectra. | 98 | 4 | astro-ph9804002_arXiv.txt |
9804 | astro-ph9804287_arXiv.txt | Randich and Schmitt [1995, A\&A 298, 115] found that the coronal activity of solar-type and low mass stars in Praesepe is significantly lower than that of stars in the Hyades cluster. This result is quite surprising since the Hyades and Praesepe have approximately the same age and metallicity and are often thought to have originated in the same Giant Molecular Cloud complex. We have carried out several tests in order to find a possible explanation for this result. We have measured radial velocities of two groups of Praesepe stars (a dF-dK sample and a dM sample) and have measured H$\alpha$ as a chromospheric activity index for the dM sample. Based on analyses of these data, we conclude that the Praesepe catalog used in the X-ray analysis does not contain a significant number of non-members, and thus that membership problems do not seem to be the cause of the Randich and Schmitt result. The comparison of the H$\alpha$ equivalent widths for the M dwarfs in Praesepe with those in the Hyades indicates that, at least for stars in this mass range, the Praesepe stars are as active or more active than their Hyades counterparts. The similarity of chromospheric emission allows us to reject differences in the rotational velocity distribution as the origin of the dissimilar Lx luminosity functions. We have also analyzed a few ROSAT PSPC pointings of Praesepe in order to obtain a new and independent estimate of the X-ray luminosities and upper limits for a small sample of Praesepe members. This analysis suggests that the previous ROSAT/PSPC analysis produced slightly optimistic X-ray upper limits; however, the differences between the old and new upper limits are not large enough to explain the dichotomy in the X-ray properties of Praesepe and the Hyades. Therefore, our examination of the available data does not provide a clear reason to explain why the X-ray luminosity functions of the two clusters are different. Part of the explanation could be found in the binaries. Speculatively, these clusters could have different orbital period distributions, with more short period binaries among the Hyades, which would show larger coronal activity. | Open clusters play a key role in the understanding of different time-dependent stellar properties such as the evolution of rotation, stellar activity, and the lithium abundance. The comparison between different clusters which have the same age allows us to prove if this approach is correct or if other effects, such as e.g. a different metal content, are also important. In this report, we examine the Hyades and Praesepe clusters. During the last 15 years, it has been demonstrated that X-ray emission is a `normal' characteristic of late type stars. As with other stellar properties which depend on rotation (in this case through the dynamo effect, Parker 1955), the emission strength decays with age, as shown by the comparison between open clusters of different ages such as the Pleiades (Caillault and Helfand 1985; Micela et al. 1985; Micela et al. 1990; Stauffer et al. 1994), and the Hyades (Stern et al. 1981; Pye et al. 1994; Stern et al. 1994; Stern et al. 1995). The Hyades cluster has been extensively studied at X-ray wavelengths. The ROSAT All-Sky Survey (RASS) detected members down to Log~Lx=1--2$\times$10$^{28}$ erg~s$^{-1}$, with detection rates of 90\% for spectral type dG, 40\% for dK and 30\% for dM stars (Stern et al. 1995). They also found that X-ray luminosity functions (XLDF) of K and M-type dwarfs are significantly affected by the presence of a large number of binary systems in the cluster, in the sense that an important fraction of the stars with strong X-ray emission were binaries. A study of the X-ray properties of Praesepe has been carried out by Randich and Schmitt (1995). The Hyades and Praesepe have similar age and metallicity although the Hyades are slightly more metal rich. Moreover, their kinetic properties are quite close (Eggen 1992) and they could have been born in the same molecular cloud. Randich and Schmitt (1995) presented the results from ROSAT PSPC Raster Scan images in a 4$^\circ$$\times$4$^\circ$ region. Their detection rates for Praesepe were 33\%, 14\% and 13\% for dG, dK and dM stars, respectively. As a consequence, the X-ray luminosity functions of Praesepe in each spectral range are dominated by the upper limits (UL). Since a large fraction of the Praesepe Raster Scan was characterized by a sensitivity similar to that of the Hyades RASS observation, the difference in the detection rates means that the bulk of the Praesepe population is underluminous in X-rays with respect to the Hyades. The goal of this papers is to try to disentangle this problem, looking for possible reasons of the disparate behavior of these coeval clusters in X-rays. We present the Praesepe data studied here and the reduction process in Section 2, where in Section 3 we analyze the data and perform a comparison with the Hyades cluster. Section 4 contains the more important conclusions derived from this study. | We have tried to establish the reasons of the different X-ray properties of late type stars in the Hyades and Praesepe. We have studied two different samples of stars: dF-dK Praesepe stars having detected and upper limits for their X-ray luminosities and dM Praesepe stars which have been not detected by ROSAT. The measured radial velocities for both samples show that contamination by spurious members cannot account for the differences in the level of coronal activity, since all stars (but one) studied here, and presumably most of the stars in the Randich and Schmitt (1995) sample, are real members. Using simultaneously color-magnitude diagrams and the measured radial velocities, we have discovered new binaries in Praesepe for the dF-dK stars. The comparison of the fraction of binaries in Praesepe and the Hyades shows that it could be slightly different in both clusters. Since the observed levels of coronal activity, assumed equal sensitivity in the observations, are lower in Praesepe, one would expect a smaller binarity rate in Praesepe than in the Hyades. Moreover, we have shown that the detection rate for the binaries is much higher in the Hyades than in Praesepe. This could be interpreted as an effect of a difference in the distribution of the orbital periods in both clusters. Finally, the study of the statistical properties of the H$\alpha$ spectral line for the dM stars in both clusters shows that in fact Praesepe presents higher chromospheric activity for this kind of stars than the Hyades. This result is also surprising, since none of the Praesepe M dwarfs were detected in X-ray, whereas many of the Hyades M dwarfs are coronally active. For this reason, one possible explanation for the differences in the X-ray properties between both coeval clusters, the existence of different distributions of the rotational velocities, seems unlikely. We have found several Praesepe dM stars which have a remarkable strong H$\alpha$ emission and very low Lx upper limits, an unexpected situation. All these data could indicate the possibility of a difference between the sensitivity of the ROSAT All-Sky Survey for the Hyades and the ROSAT PSPC observations of Praesepe. However, our re--analysis of several ROSAT pointings shows that the previous assignation of upper limits is essentially correct (although there is a suggestion that the initial estimates for the upper limits were too low by a factor of two). We propose differences in the orbital period distribution as a partial explanation of the dichotomy of the Lx properties. Extensive studies of different properties which characterize late-type stars, such as rotational velocities and periods, lithium abundances and additional activity indicators should be made in a large variety of open clusters in order to have a comprehensive perspective of the evolution of this type of stars. In particular, a similar comparison to that performed here with other clusters of the same age, such as Coma, could also contribute towards an understanding of the differences in the X-ray properties of coeval clusters. New X--ray data from AXAF or XMM could help to solve this problem. | 98 | 4 | astro-ph9804287_arXiv.txt |
9804 | astro-ph9804078_arXiv.txt | The synchrotron reflection scenario recently proposed to explain $\gamma$-ray flares observed from blazar jets is studied. Our analysis takes into account the angular distribution of the beamed radiation, the finite extent of the scattering region, and light travel-time effects. We compare energy densities and powers for synchrotron, SSC, reflected synchrotron (RSy), and external Compton (EC) scattering processes. If the width of the scattering layer is much larger than $\Gamma R^\prime_B$, where $\Gamma$ and $ R^\prime_B$ denote the bulk Lorentz factor and comoving-frame radius of the plasma blob, respectively, then the ratio of the RSy and synchrotron energy densities $\sim 4 \, \Gamma^3 n_{\rm BLR} \sigma_{\rm T} R^\prime_B$, where $n_{\rm BLR}$ is the mean particle density in the broad line region (BLR). Our results imply that Thomson-thick scattering regions of narrow extent must be present for the synchrotron reflection mechanism to operate effectively. This process seems unlikely to cause flares in lineless BL Lac sources, where X-ray and TeV flares are common and the BLR is thought to be weak or absent. We sketch time profiles of flares for various scenarios, including a model where the blob is energized by sweeping up surrounding material. | More than 50 blazar-type AGNs have been detected with high confidence by EGRET to emit $\gamma$-rays above 100~MeV (Matox et al. \markcite{Mattox97}1997). These sources are identified with flat-spectrum radio sources classified as BL-Lac objects or quasars. Many of these objects exhibit variability on all wavelengths, with some of the most rapid variability, on time scales of hours to days, observed at the highest $\gamma$-ray energies (e.g., Bloom et al. \markcite{bloom97}1997; Wagner et al. \markcite{wagner95}1995; Mukherjee et al. \markcite{mukerjee97}1997). The large apparent luminosities in combination with the short variability time scales provide evidence for the widely accepted relativistic jet model for AGNs (for recent reviews, see Schlickeiser \markcite{schlickeiser96}1996 and Hartman et al. \markcite{Hartman97}1997), according to which the radio--$\gamma$-ray emission from blazars is emitted via nonthermal synchrotron radiation and Comptonization of soft photons by energetic particles in relativistic outflows powered by accreting supermassive black holes. Soft photons which are Compton-scattered to produce the $\gamma$-ray emission include internal synchrotron photons (e.g., Marscher \& Gear \markcite{mg85}1985, Maraschi et al. \markcite{maraschi92}1992, Bloom \& Marscher \markcite{bm96}1996) and accretion-disk radiation which enters the jet directly (Dermer \& Schlickeiser \markcite{ds93}1993) and after being scattered by surrounding BLR clouds and circumnuclear debris (e.g., Sikora, Begelman \& Rees \markcite{sbr94}1994; Blandford \& Levinson \markcite{bl95}1995; Dermer, Sturner, \& Schlickeiser \markcite{dss97}1997; Protheroe \& Biermann \markcite{pb97}1997). It has recently been proposed (Ghisellini \& Madau \markcite{gm96}1996; hereafter GM96) that the beamed synchrotron radiation, after scattering off a cloud near the jet trajectory and reentering the jet, can be a source of copious soft photons and lead to a pronounced flare of very short duration as the relativistic jet plasma passes through the cloud. Wehrle et al. (\markcite{Wehrle98}1998) argue that this mechanism might explain the February 1996 $\gamma$-ray flare observed from 3C 279. A detailed analysis which correctly accounts for causality effects and the finite width of the scattering layer was not performed by GM\markcite{gm96}96, and such a treatement is required before blazar flare spectra and light curves can be modeled. Here we examine this model in more detail for a simple geometry of the BLR in the limiting regime where the nonthermal jet electron distribution does not evolve. In \S 2 we describe the model. Numerical calculations of the magnetic-field and photon energy densities and synchrotron and Compton powers are presented in \S 3. Application to blazar flares is made and the effects of blob energization on time profiles of flares are indicated. We summarize in \S 4. | Photon energy densities and radiative powers due to different processes were numerically calculated as a function of the distance of the blob from the accretion disk for a wide range of parameters. Figure 2 shows an example of our series of simulations, using parameters representative of a flat-spectrum radio quasar (FSRQ) and a BL Lac object (BL). Here we let $p = 3$ (FSRQ; solid curves) and $p = 2.7$ (BL; shaded curves). This choice yields fairly flat $\nu F_{\nu}$ spectra which can be compared with the peaks of the broadband $\nu F_{\nu}$ spectral energy distributions of blazars. The electron density and blob radius are chosen so that the resulting total luminosities are in accord with typical values of the apparent luminosities of FSRQs and BLs, and so that $< 100$ GeV $\gamma$ rays are not absorbed by $\gamma$-$\gamma$ pair production on the synchrotron photons intrinsic to the source. For the generic FSRQ, we assume that the BLR has a radial Thomson depth $\tau_{\rm T, BLR} = 0.2$ and occupies a spherical shell located between 0.05 and 0.5~pc from the central engine. For simplicity, we assume that the accretion disk radiates isotropically with a luminosity of $10^{46}$~erg~s$^{-1}$. This choice of parameters gives a Compton power which is $\sim 10$ times greater than the synchrotron power. Figure 2 shows that although the RSy radiation energy density $u'_{RSy}$ is larger than $u'_{Sy}$ when the blob is located within the BLR, the RSy radiative power is about equal to the SSC power. This is because the backscattered jet synchrotron photons are boosted in energy by a factor of $\Gamma^2$ relative to the energy of the synchrotron photons in the comoving frame; thus Klein-Nishina effects reduce the radiative power resulting from Compton scattering of RSy radiation more strongly than for the SSC radiation. In our simulations using plausible parameters compatible with the assumed spherical shell geometry of the BLR, we find that the Compton power due to the synchrotron mirror mechanism is at most comparable to the SSC power. The efficiency of this process is improved when $\Gamma \gg 10$ and a very narrow ($\Delta r_{\rm BLR} \lesssim r_{\rm in} / (2\Gamma^2)$), Thomson-thick BLR is located very far ($r_{\rm in} \gg \Gamma^2 R'_B$) from the central engine. The synchrotron reflection process might therefore operate in BLR clouds which are thought to surround Seyfert AGNs and FSRQs. BLR clouds, as understood through photoionization models (see, e.g., Wandel \markcite{Wandel97}1997 and references therein), consist of dense ($n\sim 10^{10}-10^{11}$ cm), Thomson-thick ($\tau \sim 1-10$) regions covering a small ($\sim 10$\%) fraction of the central engine. For such a model to be feasible, however, the conditions regarding duty cycle and power outlined by Dermer \& Chiang (\markcite{dc98}1998) must be met. We defer presentation of results in the regime $R^\prime_B \gtrsim \Delta r_{\rm BLR}/\Gamma$ to future work; the shell geometry used here is excessively artificial in this limit. The absence of strong emission lines in X-ray selected BLs and (to a lesser extent) radio-selected BLs suggests that the BLR is considerably more dilute in BLs than in FSRQs, and that BLs have mean Thomson thicknesses $\tau_{\rm T,BLR} \ll 0.1$ (see Scarpa \& Falomo \markcite{sf97}1997 and references therein). (On the other hand, the strength of the central ionizing photon source might be much less in BLs than FSRQs.) Superluminal motion observations also indicate that typical values of $\Gamma$ for BLs lie in the range between $\sim 3$ and 7 (see the review by Urry \& Padovani \markcite{up95}1995). The ability of the synchrotron reflection process to produce gamma-ray flares therefore seems more difficult in BLs than in FSRQs, yet TeV flares often coincident with X-ray flares have been detected from three BL Lac objects (e.g., Punch et al. \markcite{punch92}1992, Macomb et al. \markcite{Macomb95}1995; Catanese et al. \markcite{Catanese97}1997, \markcite{Catanese98}1998). For the assumed BL parameters in Figure 2, the synchrotron reflection flare could hardly be detected. If an accretion disk steadily radiates photons (see B\"ottcher \& Dermer \markcite{bd95}1995 for a treatment of time-variable disk radiation), then light-travel time effects can be neglected. The accretion disk supplies an abundant supply of soft photons, which can enter the jet directly (ECD) and after being scattered by the BLR (ECC). The comoving photon energy densities from the ECD process can dominate that from the ECC and synchrotron processes when $z \lesssim 10^{-2}$~pc, but declines $\propto z^{-3}$ and $\propto z^{-2}$ farther out (see Dermer \& Schlickeiser \markcite{ds93}1993; B\"ottcher, Mause \& Schlickeiser \markcite{bms97}1997). The ECC photon energy density increases slowly with $z$ when $z < r_{\rm in}$, and begins to decrease for $r_{\rm in} \lesssim z \lesssim r_{\rm out}$. Outside the BLR, when $z\gtrsim r_{\rm out}$, the energy density asymptotically approaches the limiting behavior $u^\prime_{\rm ECC} \propto z^{-2}$. When $z \lesssim r_{\rm in}$, the ECC process dominates over the SSC process provided that \begin{equation} {\tau_{\rm T, BLR} \over r_{\rm in}^2} \gtrsim {2 \over 3} {c \, B^2 \tau_B \over L_D \, \Gamma^2} \> \big( {p - 1 \over 3 - p}\big) \> {\gamma_2^{3 - p} - \gamma_1^{3 - p} \over \gamma_1^{1 - p} - \gamma_2^{1 - p}}, \end{equation} where $\tau_B = R^\prime_B \, n_{\rm e,jet} \, \sigma_T$ is the Thomson depth of the blob. Using the parameters adopted in Figure 2 but letting $\tau_{\rm T, BLR}$, $r_{\rm in}$, and $L$ vary, we find that the ECC photon energy density dominates the synchrotron photon energy density in the comoving blob frame when $ r_{\rm in}({\rm pc}) \lesssim 0.4 \> \tau^{1/2}_{\rm T, BLR} \, L_{46}^{1/2}$ and $ r_{\rm in}({\rm pc}) \lesssim 0.1 \> \tau^{1/2}_{\rm T, BLR} \, L_{44}^{1/2}$ for the FSRQ and BL parameters, respectively, where $L_n = L_{\rm disk} / (10^n~{\rm erg \> s}^{-1})$. It should be noted that Klein-Nishina effects are not included in estimate (6). The bottom panel of Figure 2 illustrates the time profiles of a flare calculated using the FSRQ parameters for the ECC (thick solid curve) and RSy (thick dot-dashed curve) processes. Note that the observer's time element is linearly related to $z$ for a blob moving with constant velocity. The ECC process gives a fast-rise, power-law-decay--type light curve, and the RSy mechanism gives a gradual rise of the $\gamma$-ray flux and a sharp drop as the blob leaves the BLR. A flare produced by the RSy process could be identified by a rapid decline of $\gamma$-rays which is not accompanied by a corresponding decrease of the synchrotron emission. For comparison, we also sketch a flare time profile produced by the ECD process, and time profiles produced by sweeping energization of the blob. In this process, the bulk kinetic energy of the outflowing plasmoid is converted into internal nonthermal particle energy by sweeping up BLR material (see, e.g., Panaitescu \& M\'esz\'aros \markcite{pm98}1998; Dermer \& Chiang \markcite{dc98}1998; Chiang \& Dermer \markcite{Chiang98}1998). The pair of light solid and dot-dashed curves illustrate ECC and RSy flares, respectively, are modeled assuming that the nonthermal electron energy is proportional to the amount of swept-up matter which is then added to a nonthermal lepton distribution accelerated at the base of the jet. Blob deceleration is assumed to be negligible here. The upper curves of the two pairs are modeled assuming no radiative losses, and the lower curves of the two pairs illustrate the effects of radiative losses on the nonthermal leptons by crudely multiplying the upper curves by a decaying exponential with a $1.5\times 10^{18}$~cm decay length. The model RSy light curves, with and without sweeping energization, are similar to $\gamma$-ray light curves observed in the 1991 and 1996 flares of 3C 279 (Hartman et al. \markcite{Hartman96}1996; Wehrle et al. \markcite{Wehrle98}1998). This may be a consequence, however, of the highly idealized BLR geometry used in the calculation. More symmetical flaring profiles observed from PKS 0528+134 (Collmar et al. \markcite{collmar97}1997), PKS 1622-297 (Mattox et al. \markcite{Mattox97}1997), and PKS 1406-076 (Wagner et al. \markcite{Wagner95}1995) might be more easily explained by the ECD or ECC processes. The declines of the X-ray and optical fluxes correlated with the EGRET $\gamma$-ray fluxes in the February 1996 3C 279 and the 1406-076 flares could, however, rule out the RSy mechanism since the synchrotron component is not directly affected by the reflection process. The inclusion of electron energy evolution and the relaxation of the assumption of a constant velocity blob must be treated to strengthen such conclusions. Future flare modeling must treat the passage of the jet through the BLR clouds, which themselves are in Keplerian motion around the central black hole. The passage of a jet through such a region will display a complicated signature when monitored by different telescopes with different sensitivities and imaging capabilities, which can only be decoded when full account is taken of the processes considered here. The efficiency of the different Compton-scattering scenarios, including the RSy mechanism, in this more realistic system is presently under investigation by the authors. | 98 | 4 | astro-ph9804078_arXiv.txt |
9804 | astro-ph9804308_arXiv.txt | We consider the implications of the detection of spiral structure in the accretion disc of the binary IP Pegasi. We use numerical simulations of the development of a disc outburst to construct predicted Doppler tomograms, which are found to be in close agreement with the observations if the spiral pattern arises as a transient feature when the disc expands viscously at the start of the outburst. The good agreement of such viscous disc simulations with the data is consistent with models in which most of the angular momentum transport in the disc originates in internal stresses rather than globally excited waves or shocks. Future detailed observations of the development of transient spiral features offer the potential to measure the dependence of the disc viscosity on the local physical conditions in the disc. | Recent observations of the dwarf nova IP Pegasi provide convincing evidence for spiral structure in the emission from an accretion disc in a binary system (Steeghs, Harlaftis \& Horne 1997). During an outburst, changes in the profile of spectral lines with binary phase were inverted using the technique of Doppler tomography (Marsh \& Horne 1988) to reveal a loosely wrapped, two-armed spiral pattern in the disc emission. No such structure is observed in the quiescent disc (Marsh \& Horne 1990). These observations provide a potential new constraint on the angular momentum transport processes operating in accretion discs. Two mechanisms are known that can provide a source of viscosity in ionized, non-self-gravitating accretion discs in binary systems; turbulence driven by the non-linear development of the Balbus-Hawley instability (Balbus \& Hawley 1991; Tout \& Pringle 1992; Stone et al. 1996; Brandenburg et al. 1996); and spiral waves or shocks driven by the gravitational perturbation of the secondary (Sawada, Matsuda \& Hachisu 1986; Spruit 1987; Rozyczka \& Spruit 1989; Savonije, Papaloizou \& Lin 1994). It is obvious that the second scenario leads to a spiral pattern of disc emission, but even if tidally induced shocks are unimportant for the angular momentum transport budget in the steady-state they might still be observable in outburst, when the enhanced viscosity forces the disc to expand into a region where the strength of the tidal forces is greater (Papaloizou \& Pringle 1977; Lin \& Pringle 1976). We note that although it is generally believed that the spiral shock mechanism is inefficient in the relatively cool discs found in cataclysmic variables (Livio 1994; Savonije, Papaloizou \& Lin 1994), there are considerable theoretical uncertainties in both mechanisms, and additional observational input is highly desirable. In this Letter, we compare simulations of accretion disc evolution with the observations of IP Peg. Our goal is to test whether viscous disc simulations, which are predicated on the existence of an internal origin for the disc viscosity, are consistent with the strong spiral structure observed in the data. We describe our calculations in Section 2, and present in Section 3 model Doppler tomograms for comparison with the observations. Section 4 summarises our conclusions, and outlines the theoretical expectations for spiral structure in other disc systems. | In this Letter we have presented a simulation of the evolution of the accretion disc for the parameters of the binary IP Pegasi in outburst. We find as our main result that a spiral pattern is formed in the outer disc during outburst as the enhanced viscous stresses push the disc edge into a region of strong gravitational perturbations from the secondary. The spiral structure obtained in the simulation is two-armed, non-resonant, and much more prominent in outburst as compared to quiescence. Comparing the results with the observations of Steeghs, Harlaftis \& Horne (1997), we find that there is excellent agreement with the azimuthal extent, velocity range, and asymmetry of the observed pattern. The observations of IP Pegasi and other cataclysmic variables find no clear evidence for spiral patterns in quiescent discs. This is consistent with theoretical expectations provided that the quiescent discs are cool, viscous, and have not expanded close to the tidal radius. Observations following the decay of spiral structure after the outburst has ended would be valuable in understanding the interplay of these factors. However the current evidence continues to support the conclusion of Savonije, Papaloizou \& Lin (1994) that angular momentum transport by spiral shocks is insufficiently effective to provide the bulk of the angular momentum transport in the relatively cool, thin discs found in cataclysmic variables. Spiral shocks {\em are} likely to be important over a wide range of radii in the hotter discs found in X-ray binaries (Owen \& Blondin 1997), and in the quasi-spherical accretion flows postulated for the advectively dominated regime, although internal sources of viscosity are also likely to be more efficient in those thicker disc geometries. The current observations can be modelled adequately using a highly simplified three-dimensional model of an outburst caused by a thermal disc instability, in which the only inputs are the change in $\alpha_{\rm SS}$ and $c_s$ between quiescence and outburst. Future observations, extending over the rise to outburst and during the decline, may be able to provide stronger constraints on the assumed disc model. In particular, since the spiral pattern arises as a result of the imbalance between internal viscous stresses and well-understood gravitational torques, such observations can probe the variation of the disc viscosity with the local physical conditions in the disc. Eclipsing systems such as IP Peg are particularly promising in this regard as the radial run of quantities such as the effective temperature can readily be derived simultaneously from eclipse mapping. The detection of spiral structure in the accretion disc of IP Peg, which has relatively feeble outbursts, implies that similar or stronger features may be expected in most dwarf novae. Theoretically, we note that qualitative differences are expected in systems with low mass ratio (roughly, $q < 1/4$). In these binaries both the observations of superhumps in the light curve, and numerical simulations (Murray 1996, 1998), suggest that a strong $m=1$ mode is excited in outburst. These features are {\em not} stationary in the corotating binary frame, making detailed investigation via Doppler mapping harder. It would also be worthwhile investigating whether spiral structure is observable in magnetic systems such as intermediate polars, where non-axisymmetry might be induced at the {\em inner} edge of the disc as a result of magnetic torques from the white dwarf. {\em Note added:} Godon, Livio \& Lubow (1998) have recently presented calculations showing that a steady tidally induced spiral pattern does not match the observations of IP Peg. This is consistent with our finding that consideration of the transient behaviour of the disc during outburst is required. | 98 | 4 | astro-ph9804308_arXiv.txt |
9804 | astro-ph9804293_arXiv.txt | The Carina Nebula is an extremely bright southern \HII\ region embedded in a giant molecular cloud and contains some of the most massive stars known in our Galaxy. We are undertaking a multi-wavelength study of the Carina Nebula in order to examine the detailed kinematics and distribution of the molecular and ionised gas, and to look for further evidence of ongoing star formation. Here we present the results of the initial molecular cloud observations which were made by observing the \CO\ emission with the Mopra antenna. The observations reveal the clumpy morphology of the molecular gas, and allow us to identify many interesting regions for follow-up observations. | The Carina \HII\ region/molecular cloud complex is an excellent region for studying the interaction of massive stars with their parental Giant Molecular Cloud (GMC). The nebula covers an area of $\approx$ 4 deg$^{2}$ and is bisected by a prominent V-shaped dark lane. There are over 14 star clusters in this region which have been studied extensively over the past twenty years. For excellent reviews see Feinstein (1995) and Walborn (1995). The most influential clusters of the nebula are the two OB clusters, Tr 14 and Tr 16. These clusters contain numerous O-type stars, including three O3 stars each, making them two of the most massive star clusters in our galaxy. Tr 14 is a compact cluster situated to the north-west of the nebula, adjacent to the western dust lane. Tr 16 is an open cluster centred northwards of the vertex of the dark lane. It contains one of the most massive stars known: $\eta$ Car. Here we will adopt the popular view (e.g. Tovmassian 1995, Walborn 1995) that Tr 14 and Tr 16 are at a common distance of about 2.2 kpc and that Tr 14 is younger than Tr 16. Considering the extensive studies on the stellar content of the Carina Nebula, in particular $\eta$ Car and its surrounding Homunculus nebula, relatively little work has been done on the extended nebula in the last fifteen years. Early radio continuum observations revealed that the nebula contains a large ionised region with two peaks, Car I and Car II (Gardner \& Morimoto 1968). Higher resolution radio continuum data show that both Car I and Car II are made up of a number of filamentary arcs and rings and are everywhere thermal (Retallack 1983, Whiteoak 1994). Car II is located to the north of $\eta$ Car and Car I is located towards the western dark lane, just west of Tr 14. The dynamics of the ionised gas in this region have been studied via hydrogen recombination line emission (Gardner et al. 1970, Huchtmeier \& Day 1975) and H$\alpha$ and [\NII] emission observations (Deharveng \& Maucherat 1975). The results show line splitting towards the Car II region which has been interpreted as an expanding shell of ionised gas. The dark lanes consist of molecular gas and dust that are associated with the nebula (Dickel 1974). H$_{2}$CO and OH absorption measurements identified two optical depth maxima which were located towards these lanes (Gardner, Dickel \& Whiteoak 1970, Dickel \& Wall 1974). Extended far-IR emission is confined there also (Harvey Hoffmann \& Campbell 1979, Ghosh et al. 1988). There are two main CO emission regions towards the nebula; a northern and southern cloud (de Graauw et al. 1981, Whiteoak \& Otrupcek 1984). Both regions are part of a much larger GMC which has a projected length of 130 pc and a mass in excess of 5$\times$10$^{5}$ \Msun (Grabelsky et al. 1988). The area between the southern and northern CO clouds is centred on the Keyhole Nebula, a dense dark cloud northwest of $\eta$ Car. Here the molecular gas exists in dense clumps of typical mass 10 \Msun\ that are separated both in space and velocity (Cox \& Bronfman 1995). The picture used to describe the Carina complex is one in which the massive star clusters, Tr 14 and Tr 16, are interacting strongly with the molecular cloud from which they formed. It is generally accepted that the photons from Tr 14 and Tr 16 are responsible for the ionised emission of Car I and Car II respectively, and that their strong stellar winds are producing the general expansion of the nebula. We are undertaking a multi-wavelength study of the Carina Complex in order to study the detailed kinematics and distribution of the molecular and ionised gas and to look for further evidence of ongoing star formation. Here we present the results of initial observations of the \CO\ emission. CO emission is thermalised in both low- and high-density gas and therefore is suitable for tracing the overall distribution and velocity structure of the molecular cloud. It also can pinpoint any `CO hot-spots'. These are warm molecular cores where stars could possibly be forming. | We have presented data from the first stage of an extensive study of the GMC associated with the Carina nebula. The data consists of observations of \CO\ emission which have been used to trace the overall GMC as well as pinpoint and CO `hot-spots' or dense regions where stars could possible form. The observations are at a higher resolution than previous studies and reveal the clumpy nature of the northern and southern cloud regions. They also show the positional coincidence between the far-infrared emission and the strong CO emission towards the Car I region. This supports a blister-type model for this region. Further observation of different transitions are being made to better constrain the temperature and density of the molecular gas in this interesting region. | 98 | 4 | astro-ph9804293_arXiv.txt |
9804 | gr-qc9804051_arXiv.txt | Linde's proposal of a Euclidean path integral with the ``wrong'' sign of Euclidean action is often identified with the tunneling proposal for the wave function of the universe. However, the two proposals are in fact quite different. I illustrate the difference and point out that recent criticism by Hawking and Turok does not apply to the tunneling proposal. | 98 | 4 | gr-qc9804051_arXiv.txt |
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9804 | astro-ph9804016_arXiv.txt | We report on a near-infrared, long-baseline interferometric search for luminous companions to the star 51~Pegasi conducted with the Palomar Testbed Interferometer. Our data is completely consistent with a single-star hypothesis. We find no evidence to suggest a luminous companion to 51~Pegasi, and can exclude a companion brighter than a $\Delta$K of 4.27 at the 99\% confidence level for the 4.2-day orbital period indicated by spectroscopic measurements. This $\Delta$K corresponds to an upper limit in the companion M$_K$ of 7.30, in turn implying a main-sequence companion mass less than 0.22 M$_{\sun}$. | The recent inference of a planetary-mass gravitational companion to the star 51~Pegasi (HD 217014) from apparent radial velocity variation by Mayor \& Queloz (1995) has subjected this otherwise unremarkable star to remarkable scrutiny. The Mayor and Queloz result was quickly verified by several groups with similar or higher-resolution spectroscopic techniques (c.f.~\cite{Marcy97}). However, there has been no other evidence for a companion, e.g.~precision photometric monitoring has failed to show evidence for eclipses (\cite{Henry97}), and there is a significant lack of x-ray flux from the system compared to binary systems with similar periods (\cite{Pravdo96}). Further, 51~Peg's G5V spectral classification has become mildly controversial (e.g.~\cite{Houk95}, who argues for a G2-3V), as has its physical size (e.g.~\cite{Hatzes97,Henry97}). A planetary-mass companion in a 4.2 day orbit around a solar-mass 51~Peg would have an orbital semi-major axis of approximately 0.05 AU (\cite{Marcy97}), slightly more if the companion were more massive. At a distance of 15.4 $\pm$ 0.2 pc (\cite{Perryman96}), the approximate maximum primary-companion angular separation would be 3.5 millarcseconds (mas). Such an angular separation is well below resolution limits for current conventional imaging technology, but is accessible to optical and near-infrared interferometry. As only the lower mass limit is set by the spectroscopic results, it is possible the companion is significantly more massive -- perhaps even a low-mass star. We have therefore studied 51~Peg with the Palomar Testbed Interferometer (PTI) in an attempt to detect the putative companion if it is indeed sufficiently luminous. PTI is a 110m-baseline interferometer operating at K-band (2 -- 2.4 $\mu$m) located at Palomar Observatory, and described in detail elsewhere (\cite{Colavita94}). The minimum PTI fringe spacing is roughly 4 mas at the sky position of 51~Peg, making a (sufficiently) luminous companion readily detectable. | We find no evidence to suggest that the putative 4.2-day period companion to 51~Peg is detectable in our data; all of the datasets we have analyzed indicate that 51~Peg is at least as stable as our two calibration sources. The 1997 PTI data on 51~Peg is sufficiently stable that we can place significant limits on $\Delta$K and consequently M$_K$ of a 4.2-day period companion. We find upper limits in $\Delta$K of 4.78, 4.53, and 4.27 for the 4.2-day period companion to 51~Peg at 68\%, 95\%, and 99\% confidence levels respectively. These $\Delta$K limits imply companion M$_K$ limits of 7.81, 7.56, and 7.30, corresponding to upper limits on the mass of a putative main sequence companion at 0.17, 0.20, and 0.22 M$_{\sun}$ at the 68\%, 95\%, and 99\% confidence levels respectively (\cite{Henry93}). Our results cannot exclude the possibility of a very low-mass star in a face-on orbit as the 51~Peg companion, but such a star would have to be of spectral type M5V or later. | 98 | 4 | astro-ph9804016_arXiv.txt |
9804 | astro-ph9804220_arXiv.txt | We numerically investigate stellar and gas dynamics in star-forming and dissipative galaxy mergers between two disk galaxies with specific orbital configurations. We find that violent relaxation combined with gaseous dissipation in galaxy merging transforms two disk galaxies into one S0 galaxy with polar-rings: Both the central S0-like host and the polar-ring component in a polar-ring galaxy are originally disk galaxies. We also find that morphology of the developed polar-rings reflects both the initial orbit configuration of galaxy merging and the initial mass ratio of the two merger progenitor disk galaxies. Based upon these results, we discuss the origin of the fundamental observational properties of polar-ring galaxies, such as the prevalence of S0 galaxies among polar-ring galaxies, the rarity of polar-ring galaxies among S0 galaxies, the dichotomy between narrow polar-rings and annular ones, shapes of polar-ring warps, and an appreciably larger amount of interstellar gas in the polar-ring component. | Polar-ring galaxies are generally considered to be dynamically peculiar systems in which the outer rings composed of gas and stars are aligned roughly in a perpendicular orientation with respect to the major axis of the central host galaxies (Schweizer, Whitmore, \& Rubin 1983; Whitmore et al. 1990; Sackett 1991). A growing number of observational studies have been recently accumulated which can provide valuable information about the origin of these peculiar polar-ring galaxies. Nearly all of the central host are morphologically normal S0 galaxies, some of which are confirmed to be rapidly rotating by kinematical studies (Schechter \& Gunn 1978; Schechter, Ulrich, \& Boksenberg 1984; Whitmore et al. 1990; Whitmore 1991). Approximately only 0.5 percent of all S0 galaxies have observable polar-rings, which suggests that a particular mechanism is required for the formation of polar-rings in S0 galaxies. The ring component also shows rapid rotation comparable to that of the main host galaxies, implying that two dynamically different system coexist in these polar-ring galaxies. An appreciable amount of HI gas, which is sometimes comparable to the total mass of the host, is closely associated with the stellar ring component (e.g., Shane 1980; Schechter et al. 1984; Richter, Sackett, \& Sparke 1994; Arnaboldi et al. 1997; Galletta, Sage, \& Sparke 1997). The morphology of the polar-rings is basically divided into two broad classes (Whitmore 1991): A narrow ring which is not extended in size (e.g., ESO 415-G 26) and an annulus which is a disk-like component with the central part cut out (e.g., NGC 4650a). Peculiar morphology is observed in some polar-ring galaxies (e.g., the Helix galaxy, NGC 2685, and double ringed system, ESO 474-G 26), which further implies considerably complicated physical processes in polar-ring formation and simultaneously provides a clue to the understanding of the origin of polar-ring galaxies (Sackett 1991). Roughly two-thirds of these polar-rings show obvious galactic warps whose shapes look like `integral sign' and/or `banana'(Whitmore 1991). Statistical studies on the distribution of the angle between the ring component and the central host reveal that these two components strongly prefer to be orthogonal with each other. These peculiarities both in the kinematics and morphology observed in polar-ring galaxies have attracted a number of theoretical interests, which are divided basically into two categories: One is the origin of the polar-rings and the other is the nature of dark matter halo surrounding polar-ring galaxies. Although there are a large number of important studies addressing the three dimensional shapes of dark matter halo in the galaxies (e.g., Whitmore, Mcelroy, \& Schweizer 1987; Reshetnikov \& Combes 1994; Sackett et al. 1994; Combes \& Arnaboldi 1996), we here restrict ourselves to the mechanisms which would naturally explain the formation of the polar-ring galaxies with spherical haloes. It is generally believed that the formation of the polar-rings is the results of a `secondary event' involving a pre-existing S0 galaxy (e.g., Steiman-Cameron \& Durisen 1982; Sparke 1986; Quinn 1991; Rix \& Katz 1991; Reshetnikov \& Sotnikova 1997). Specifically, the host S0 galaxy is supposed to have acquire the material constituting the ring component by capturing the gas during tidal interaction with neighbor galaxies. The subsequent gravitational interaction combined with the gaseous dissipation then spreads the captured gas and forms the polar-rings around the host galaxy. One of the promising models along this orthodox scenario is the `preferred plane model' in which the differential precession of the rings and the gaseous dissipation cooperate to play a vital role in leading the acquired gas to settle into the stable polar orbit and finally to form the polar-rings (Tohline \& Osterbrock 1982; Durisen et al. 1983; Schweizer et al 1983). A number of numerical simulations have already confirmed in what physical conditions the polar-rings are more likely to form and continue to exist for a relatively longer time-scale (Habe \& Ikeuchi 1988; Christodoulou et al. 1992; Katz \& Rix 1992). Indeed these previously proposed models have provided a potential success in reproducing the polar-rings in S0 galaxies, however, these seem to be incapable of giving sufficiently conclusive and persuasive answers to the following seven questions on the origin of the polar-ring galaxies (Sackett 1991; Whitmore 1991; Arnaboldi et al. 1997; Galletta et al. 1997): (1) Why are nearly all the central host galaxies morphologically classified as S0 ? (2) Why are polar-ring galaxies so rare among S0 galaxies ? (3) Why do some polar-ring galaxies have a narrow ring and some have annuli ? (4) Why are the mass and angular momentum of the ring component comparable or sometimes larger than those of the host ? (5) Why are the rings so `polar' ? (6) Why some polar-rings have considerably peculiar morphology such as helical and double-ringed shapes ? (7) Why is there an appreciably greater amount of interstellar gas in the polar-ring component ? In particular, (1), (4), and (7) could not be explained simply by the previous theoretical models, implying either that more elaborated and sophisticated models along the above orthodox scenario should be considered or that the alternative model should be proposed for the explanation of the above questions. The purpose of this paper is to explore the origan of polar-ring galaxies and to propound a new mechanism which more naturally and reasonably explains the aforementioned observational trends of polar-ring galaxies. In the present study, we consider that the dissipative galaxy merging between two disks is a promising mechanism that quite reasonably answers the above seven questions. Therefore we investigate how the dissipative galaxy merging transforms two disks into one early-type S0 galaxy with polar-ring. Furthermore, we investigate how the orbit configuration of galaxy merging and initial mass ratio of the two progenitor disk galaxies can affect the morphology of polar-rings developed after galaxy merging. In this paper, the galaxy merging with specific orbit configurations and sufficient amount of gaseous dissipation is demonstrated to play a vital role in forming both the cental S0-like host and the surrounding polar-ring component in polar-ring galaxies. This paper is an extended version of Bekki (1997) in which the basic mechanism of polar-ring S0 galaxy formation in galaxy mergers is briefly summarized. The layout of this paper is as follows. \S 2 describes numerical models for dissipative galaxy merging. \S 3 gives the results obtained in the present study. In \S 4, we mainly discuss whether or not the model proposed in the present paper can become a new promising model which naturally and reasonably explains the observational properties of polar-ring S0 galaxies. | The present numerical study provides a new mechanism by which both the central S0-like host and the ring component in a polar-ring galaxy are simultaneously formed. Although uncertainties of the numerical treatment of gas dynamics and star formation still remain, it appears that our model has succeeded in reproducing $some$ polar-ring galaxies and explaining naturally a number of important observational properties of them. In the proposed model, the formation of polar-ring galaxies is essentially ascribed to the details of dynamics of dissipative galaxy merging with specific orbital configurations. Specifically, the central host of a polar-ring galaxy is the galaxy which has been inevitably transformed from a late-type spiral into an early-type S0 galaxy during the merging. The ring component, on the other hand, is the `galaxy' which has been dramatically transformed from a late-type spiral into a narrow ring or annuli owing to the violent gravitational interaction and gaseous dissipation during the merging. Although both specific orbital configurations and gaseous dissipation in galaxy merging are required for the formation of polar-ring galaxies in the present model, these constraints also give natural explanations to observed trends such as the prevalence of S0 among polar-ring galaxies (e.g., Whitmore 1990), the rarity of the polar-ring galaxies among S0 galaxies (e.g., Whitmore et al. 1991), and an appreciably larger amount of interstellar gas in polar-rings (e.g., Sackett 1991). Moreover it is found that the morphology of polar-rings such as a narrow ring (e.g., ESO 415 - G 26), annular rings (e.g., NGC 4650a), helical rings (e.g., NGC 2658), and double rings (e.g., ESO 474 - G 26) can reflect both the orbital parameters of galaxy merging and the initial mass ratio of two merger precursor galaxies. Thus, the present study demonstrates that a merger remnant of a gas-rich galaxy merger is one of promising candidates of polar-ring S0 galaxies. | 98 | 4 | astro-ph9804220_arXiv.txt |
9804 | astro-ph9804150_arXiv.txt | We present a new, fully covariant and manifestly gauge-invariant expression for the temperature anisotropy in the cosmic microwave background radiation resulting from scalar perturbations. We pay particular attention to gauge issues such as the definition of the temperature perturbation and the placing of the last scattering surface. In the instantaneous recombination approximation, the expression may be integrated up to a Rees-Sciama term for arbitrary matter descriptions in flat, open and closed universes. We discuss the interpretation of our result in the baryon-dominated limit using numerical solutions for conditions on the last scattering surface, and confirm that for adiabatic perturbations the dominant contribution to the anisotropy on intermediate scales (the location of the Doppler peaks) may be understood in terms of the spatial inhomogeneity of the radiation temperature in the baryon rest frame. Finally, we show how this term enters the usual Sachs-Wolfe type calculations (it is rarely seen in such analyses) when subtle gauge effects at the last scattering surface are treated correctly. | The calculation of the primary temperature anisotropy in the cosmic microwave background radiation (CMB) resulting from density perturbations has a long history, beginning with the seminal paper by Sachs and Wolfe~\cite{sachs67}. Since the original Sachs-Wolfe estimate, a wealth of detailed predictions for the anisotropies expected in various cosmological models have been worked out. The calculations are straightforward in principle, but, like many topics in cosmological perturbation theory, are plagued by subtle gauge issues~\cite{stoeger95a}. The problems of gauge-mode solutions to the linear perturbation equations and the gauge-ambiguity of initial conditions can be eliminated by working exclusively with gauge-invariant variables, as in the widely used Bardeen approach~\cite{bardeen80} and the less well known covariant approach advocated by Ellis and coworkers~\cite{ellis89a,ellis89b}. However, gauge issues still arise in connection with the definition of the temperature perturbation and the placement of the last scattering surface~\cite{stoeger95a,ellis-er97}. The latter gauge issues do not arise at first-order in numerical calculations which integrate the Boltzmann equation in a perturbed universe, since the visibility function (which determines the position of the last scattering surface) multiplies first-order variables giving only a second-order error from the use of a zero-order approximation to the visibility~\cite{LC-scalcmb}. However, this is not always the case in Sachs-Wolfe style analyses, which integrate along null geodesics back to the surface of last scattering, unless care is taken to ensure that the final result involves only first-order variables on the last scattering surface, which then only need be located to zero-order. In this paper, we present a new expression for the CMB temperature anisotropy arising from linear scalar perturbations which is fully covariant and manifestly gauge-invariant. We obtain our expression by integrating the covariant and gauge-invariant Boltzmann equation~\cite{LC-scalcmb,LC97-er} along observational null geodesics, paying careful attention to the gauge issues discussed above. Unlike some covariant results in the literature (see, for example,~\cite{LC97-er,dunsby96b}), the expression derived here can be integrated trivially, in the instantaneous recombination approximation, up to a Rees-Sciama term in universes with arbitrary matter descriptions. (The covariant results in~\cite{LC97-er,dunsby96b} can only be integrated in baryon-dominated universes, thus excluding CDM dominated universes, and other such models favoured by observation.) We base our treatment on the physically appealing covariant and gauge-invariant formulation of perturbation theory, as described in~\cite{ellis89a,ellis89b}. In this approach, one works exclusively with gauge-invariant variables which are covariantly-defined and hence physically observable in principle. The covariant method has many advantages over other gauge-invariant approaches (such as that formulated by Bardeen~\cite{bardeen80}). Most notably, the covariant variables have transparent physical definitions which ensures that predictions are always straightforward to interpret physically. Other advantages include the unified treatment of scalar, vector and tensor modes, a systematic linearisation procedure which can be extended to consider higher-order effects (the covariant variables are exactly gauge-invariant, independent of any perturbative expansion), and the ability to linearise about a variety of background models, such as Friedmann-Robertson-Walker (FRW) or Bianchi models. For universes which are baryon-dominated at last scattering, our expression for the temperature anisotropy may be compared to other gauge-invariant analytic results in the literature. We show that, with suitable approximations, the result derived here reduces to that given by Panek~\cite{panek86} and corrects a similar result given by Dunsby recently~\cite{dunsby96b}. For the baryon-dominated universe, we use numerical results for the covariant, gauge-invariant variables on the last scattering surface, obtained from a gauge-invariant Boltzmann code~\cite{LC-scalcmb}, to discuss the different physical contributions to the primary temperature anisotropy. In particular, we show that on intermediate and small scales, the ``monopole'' contribution to the temperature anisotropy is described by the spatial gradient of the photon energy density, in the energy-frame, on the last scattering surface. Since the (real) last scattering surface is approximately a surface of constant radiation temperature (so that recombination does occur there), the inhomogeneity of the radiation energy density in the energy-frame determines a distortion of the last scattering surface relative to the surfaces of simultaneity in the energy-frame. The extra redshift (due to the expansion of the universe) which the photons incur due to the distortion is seen as a ``monopole'' contribution to the temperature anisotropy on intermediate scales. There is a significant ``dipole'' contribution to the anisotropy on intermediate and small scales, which we discuss also. We end with a discussion of the gauge issues inherent in the original Sachs-Wolfe calculation of the CMB anisotropy~\cite{sachs67}, focusing on the ``monopole'' contribution to the temperature anisotropy on intermediate scales, described above. This contribution is often missed in Sachs-Wolfe type calculations through an incorrect treatment of gauge effects at the last scattering surface~\cite{ellis-er97}. (Equivalently, the term is often missed through a failure to recognise the direction-dependence of the ``expected temperature'' used to define the temperature perturbation in many calculations.) This often neglected term, which is not important on large scales, is an essential component of the Doppler peaks in the CMB power spectrum. We employ standard general relativity and use a $(+---)$ metric signature. Our conventions for the Riemann and Ricci tensors are fixed by $[\nabla_{a},\nabla_{b}] u^{c} = -{{\mathcal{R}}_{abd}}^{c} u^{d}$, and ${\mathcal{R}}_{ab} \equiv {{\mathcal{R}}_{abc}}^{b}$. We use units with $c=G=1$ throughout. | Starting from a covariant and gauge-invariant formulation of the Boltzmann equation, we have derived a new expression for the CMB temperature anisotropy under the instantaneous recombination approximation, valid for scalar perturbations in open, closed and flat universes. Our expression uses only covariantly-defined variables, and is manifestly gauge-invariant. The result is more useful in multicomponent models with scalar perturbations than earlier covariant results~\cite{ellis-er97,LC97-er,dunsby96b}. In the case of a universe which is baryon-dominated at recombination, we find a simple expression for the anisotropy which corrects a similar result by Dunsby~\cite{dunsby96b}. By making use of numerical solutions to the perturbation equations, we have discussed the conditions on the last scattering surface and their contributions to the characteristic features of the CMB power spectrum. We ended with a discussion of the original Sachs-Wolfe calculation for the temperature anisotropy. We have discussed why it is not necessary to locate accurately the last scattering surface in such calculations (because of the compensation effect), and how the extra term in the Sachs-Wolfe calculation, reported recently by Ellis and Dunsby~\cite{ellis-er97}, is missed in many calculations which employ a gauge-dependent ``expected temperature'', since the angular dependence of this temperature is often overlooked. For a universe which is matter dominated at recombination, but not necessarily adiabatic, the extra term is the spatial gradient of the radiation energy density in the energy-frame, $(\bar{{\mathcal{X}}}_{k}^{(\gamma)} Q^{(k)} /4)_{A}$. For models with isothermal surfaces of last scattering, this inhomogeneity describes a distortion of the last scattering surface relative to the surfaces of simultaneity of the energy-frame. The extra redshift incurred by this distortion is a significant component of the temperature anisotropy on intermediate and small scales. | 98 | 4 | astro-ph9804150_arXiv.txt |
9804 | astro-ph9804193_arXiv.txt | We search OGLE-I photometric database for stars, which, as defined by formal criteria adopted by OGLE-I microlensing search, showed variability during only one out of 3 or 4 observing seasons. The results include 17 previously reported microlensing events, 2 newly discovered candidate events and 15 intrinsically variable stars that have a potential of contaminating samples of microlensing events. Based on photometry obtained in 1992 and 1993 OGLE \#10 was tentatively included in the list of microlensing candidates, however its light curve in 1994 and 1995 shows many characteristics of the variable stars found in our search, and most likely it is not a microlensing event. For all stars which passed our tests, we provide $44 \times 44$ arcsec ($101 \times 101$ pixels) centered subframes from each OGLE-I frame in $I$ band. It is the first time when images used to derive photometry of microlensing events are available in convenient format to astronomical community. | The search for rare cases of gravitational microlensing in the Local Group requires monitoring of $\sim 10^6$ stars over several months in order to yield a significant rate of detections. A common implementation adopted is to construct a massive photometry database and subsequently select stars which experienced brightening of the type we expect on theoretical grounds (see Paczy\'nski, 1996, for a recent review of basic theory, current microlensing searches and results). For the vast majority of events the light curves should follow a single point mass microlensing curve. Possible departures and exceptions from this basic case are extremely interesting. For random distribution of stars the probability that a given star will be lensed twice over duration of the experiment is negligible. Large fraction of binary lenses is expected to give raise to the population of ``wide binary'' light curves with two separate amplification regions ($\sim 1\%$ of the total rate of events). However in most cases the secondary peak should be weak ($A_{\rm max} < 0.1$) and probably would only be discovered in light curves, which called attention because of the primary event (Di Stefano and Mao 1996). Photometry in very crowded fields, which are natural places to look for microlensing events, is often of limited quality. The majority of stars measured are just above the detection limit, many events have modest amplitudes, and numerous light curves are unevenly sampled. Therefore a microlensing curve may accidentally give a good fit to the brightening which is due to the intrinsic variability of the star. Certainly a repeated brightening of the same type would need to be very carefully examined before any claim of a detection of the wide binary microlensing repeater. As a result, it is generally hard to lower confusion rate without lowering the number of events returned by the procedure. For microlensing events reported by OGLE-I project (Udalski et al. 1992) the basic requirement was that variability should occur during only one observing season (Udalski et al. 1994a). It was assumed that a sample of stars selected according to the above condition contains the majority of microlensing events and relatively few variable stars of other types. In this paper we investigate how many variable stars have light curves that, given the time sampling of OGLE-I experiment, appear constant during all seasons except for just one. The quantitative information about stellar variability background, against which microlensing events are detected, allows further tuning of the methods used in automated detection of the events. | A general conclusion is that variability background in OGLE-I search was reasonably well separated from microlensing events, although OGLE \#10, a single candidate which most likely is not a microlensing event, constitutes a 5\% confusion rate. A variable star, just like any other star, may be amplified by microlensing, but an increase of brightness by 0.1 mag may be naturally explained by intrinsic variability of this star, especially that similar objects (certainly not microlensing events) apparently exist. We find two additional possible events that were not returned by the automated procedure of Udalski et al. (1994a). This is not surprising since we relax some of the selection cuts applied before. Moreover, one of those events happened near the end of the observing season while the other had very short time scale and poorly sampled light curve. Both of them are very inconspicuous. The outburst experienced by MM3 $I$ 58214 (most likely a flare or CV star) is an important case which has a potential of contaminating samples of microlensing candidates. Suppose we had no data just before the flare. With photometric accuracy comparable to OGLE-I data such variable could be taken for a fading microlensing event and uncertainty would have to be resolved by spectroscopy and/or monitoring of the star long after the event. We note a relatively large number of periodic or almost periodic variables which change amplitude. They mimic constant stars for time long enough to pass the most important criterion of the OGLE-I search, i.e., that a star should vary within a limited time interval with essentially constant flux at all other times. It is mostly due to relatively poor time coverage of the OGLE-I experiment and should not be difficult to overcome in the second phase of the project, OGLE-II (Udalski, Kubiak and Szyma\'nski 1997). Two recommendations can be made for the future. First, low amplitude events, e.g. with $A_{\rm max} < 1.5$, may be safely ignored in the calculation of the optical depth to prevent potential problems with objects similar to OGLE \#10. MACHO team is already using such cut off to avoid contamination by ``bumpers''. Second, a requirement of roughly even photometric coverage of both sides of the magnification peak allows filtering out stars with (usually asymmetric) outbursts. OGLE-I events used in the optical depth determination satisfy such condition, however some of the events discovered by the EWS do not. The full set of data used in this paper, including $101 \times 101$ pix subframes extracted from every $I$ band image of each object in Tables~1 and 2, is available for public. Images in FITS format as well as photometric data in standard Johnson system may be retrieved via anonymous ftp from {\tt astro.princeton.edu (128.112.24.45)} -- directory {\tt /ogle/var\_background} and {\tt sirius.astrouw.edu.pl (148.81.8.1)} -- directory /ogle/var\_background. See {\tt README} file for details. | 98 | 4 | astro-ph9804193_arXiv.txt |
9804 | astro-ph9804299_arXiv.txt | We present X-ray spectra obtained with \sax\ (Satellite per Astronomia X) of 10 BL Lac objects, selected from the Einstein Medium Sensitivity and Slew Surveys. We find that in about half of the objects a fit in the 0.1-10 keV range with a single power law and free absorption yields values of $N_{\rm H}$ larger than the Galactic ones. In most of these cases, however, broken power law fits with $N_{\rm H}$ fixed at the Galactic values yield an alternative, better description of the data and indicate a steepening of the spectrum with increasing energy. One object (1ES1101-232) is detected up to $\sim$ 100 keV. Its spectral energy distribution (SED) peaks in the medium energy X-ray band. For each object we compute the peak frequency of the SED from multifrequency data. The spectral indices $\alpha_x$ in the 2-10 keV band ($F_\nu \propto \nu^{-\alpha_x}$) are smaller (i.e. flatter spectrum) for objects with higher peak frequencies. We therefore confirm and extend to higher energies the behavior already known for X-ray selected BL Lac objects in the ROSAT band. We do not find spectral indices smaller than 1; however, the flat distribution of $\alpha_x$ and the correlation between $\alpha_x$ and peak frequency found from our data suggest that a number of objects may exist, which in the quiescent status have flatter spectrum and peak frequency in the hard X-ray range. | \noindent BL Lacertae objects are a rare type of Active Galactic Nuclei (AGN) characterized by strong and variable emission of non-thermal radiation across the entire electromagnetic spectrum, from radio waves to high energy $\gamma$-rays. In three cases (Mkn 421: Punch et al. 1992; Mkn 501: Quinn et al. 1996; 1ES 2344+514: Catanese et al., 1998) the emission has been detected up to TeV energies. BL Lacertae objects comprise the most violent (highly and rapidly variable, highly polarized) and most elusive (extremely difficult to find in optical surveys) sources amongst AGN. Unlike most other AGN they do not show evidence (by definition) for strong emission lines or large Infra-Red or UV excesses. The emission from radio to $\gamma$-rays can be explained as due to synchrotron radiation up to a certain maximum frequency (that ranges approximately from $10^{13}$ to $10^{17}$ Hz), above which a sharp turnover occurs until a second component due to Compton scattered radiation dominates, making these objects detectable up to the highest energies so far accessible (see e.g., Ulrich, Maraschi and Urry, 1997). The extreme properties of BL Lacs require that the matter emitting the radiation moves at relativistic speeds in the direction of the observer. The spectral change from synchrotron to Compton radiation is crucial for the understanding of the physics of BL Lacs. However, up to now this has been inferred only from the comparison of X-ray measurements carried out with different instruments and very often at different epochs. The wide energy band of \sax\ offers for the brightest objects the best opportunity to directly detect without ambiguity this spectral change and to study the X-ray spectra at the same epoch over a large interval. To this end we have undertaken a program that aims at studying in detail the X-ray spectrum of a large and well defined subsample of soft X-ray selected BL Lacs. This sample includes mostly objects that are expected to show strong spectral curvature and spectral breaks, since the synchrotron break should occur just before or in the \sax\ band. We aim at measuring in detail the shape of the most energetic part of the synchrotron emission, and trying to establish where and how the Compton component becomes dominant. We also intend to look for the correlation between spectral slope and break energy found in ROSAT data (Padovani $\&$ Giommi, 1996; Lamer, Brunner \& Staubert, 1996). | We have analyzed the spectra for 10 X-ray selected BL Lacs observed with the Narrow Field Instruments on board the \sax\ satellite. The sources are detected from $\sim 0.2$ up to $\sim$ 10 keV (and in one case up to $\sim 100$ keV with the PDS instrument) with a very smooth appearance. The spectrum is generally well fitted by either a single power law, or by a broken convex power law that most probably represents the steepening after the synchrotron peak, whose position is determined also by using simultaneous optical observations. Variability is not present during the short \sax\ exposure; analysis of ROSAT data shows for most of the sources little variability (within $\sim 30\%$) with respect to the \sax\ flux and spectral indices consistent with the \sax\ ones. The spectral energy distributions, which include literature data, instead show variability in all bands. The X-ray spectral indices $\alpha_x$ range between 1 and 1.5 with a flat distribution and a mean value $\langle \alpha_x \rangle = 1.31\pm0.06$. The scatter in the distribution is due to an anti-correlation we have found between $\alpha_x$ and the frequency of the peak of the emission, $\nu_{peak}$. This extends to the \sax\ band a correlation which had been discovered in the ROSAT band for this class of objects. The fact that sources with harder X-ray spectra have higher $\nu_{peak}$ is expected if the \sax\ band is still dominated by synchrotron emission, which is also consistent with the spectral energy distributions of our BL Lacs. Furthermore, we have no evidence of a spectral flattening (indicating the arising of the Compton component) in the present spectra, but future PDS detections, that are possible with exposure times slightly longer than those obtained here, might help in this respect. The large fraction (at least 2 out of 10) of HBL selected in the soft X-ray band found with a flat ($\alpha_x \sim 1$) X-ray slope (i.e., they are near the peak of the synchrotron emission) and the distribution of $\alpha_x$ values support the view that objects with even higher spectral peaks in their quiescent status indeed exist, and might be found in large numbers if we devise the correct strategy (e.g., samples at harder X-rays, TeV sources, etc.) Moreover, these sources are good candidates to be TeV {\it emitters}. In fact, in the sources with the flattest $\alpha_x$ the peak of the synchrotron component is localized in the soft X--ray range. Electrons emitting at 1 keV by the synchrotron process have Lorentz factors $\gamma\sim 2.5\times 10^5 (\nu_{peak,1~keV}/B\delta)^{1/2}$, where $\nu_{peak} = 2.42 \times 10^{17} \nu_{peak,1~keV}$ Hz, $B$ is the value of the magnetic field in Gauss and $\delta$ is the usual Doppler factor. Through the inverse Compton mechanism, they can emit up to $E\sim \gamma m_ec^2\delta\sim 130 (\nu_{peak,1~keV}\delta/B)^{1/2}$ GeV. If the magnetic and radiation energy densities are equal (as it is, approximately, in the three BL Lacs already detected in the TeV band), the flux level of the synchrotron and inverse Compton peaks is roughly equal. Low redshift sources are therefore good candidates to be {\it detected} in the GeV--TeV band, while the high energy emission of the more distant sources could be absorbed in $\gamma$--$\gamma$ interactions with the background IR photon field, whose intensity is still uncertain. Indeed, a cutoff in the high energy spectrum could be used to determine the IR background (see, e.g. Stecker \& De Jager, 1997). This \sax\ project is still ongoing. We expect therefore to increase considerably the sample of soft X-ray selected BL Lacs for which we measure the spectrum in the 0.2-10 keV range and possibly above. With a larger complete sample, and combining the results with other complementary \sax\ projects, we expect to be able to draw a clearer picture of the relationship between the local X-ray slope and the overall energy distribution of this class of sources, in order to derive firmer conclusions on the behaviour at hard X-ray energies and on the mechanisms of the emission. | 98 | 4 | astro-ph9804299_arXiv.txt |
9804 | astro-ph9804066_arXiv.txt | About 25\% of the optical extragalactic sky is obscured by the dust and stars of our Milky Way. Dynamically important structures might still lie hidden in this zone. Various approaches are presently being employed to uncover the galaxy distribution in this Zone of Avoidance (ZOA). Results as well as the different limitations and selection effects from these multi-wavelengths explorations are being discussed. Galaxies within the innermost part of the Milky Way --- typically at a foreground obscuration in the blue of $A_{\rm B} \ga 5^m$ and $|b| \la \pm5\degr$ --- remain particularly difficult to uncover except for H\,{\sc i}-surveys: the Galaxy is fully transparent at the 21cm line and H\,{\sc i}-rich galaxies are easy to trace. We will report here on the first results from the systematic blind H\,{\sc i}-search ($v \leq 12700$ km\,s$^{-1}$) in the southern Zone of Avoidance which is currently being conducted with the Parkes Multibeam (MB) Receiver. | To understand the dynamics within the local Universe -- the mass distribution and the local velocity field with its peculiar and streaming motions -- a detailed map of the 3-dimensional galaxy distribution is highly desirable. However, the dust extinction and confusion with stars in the disk of our Galaxy make this very difficult for $\sim$25\% of the sky, and the following questions remain unanswered: Could a nearby Andromeda-like galaxy have escaped detection to date, hence change our understanding of the internal dynamics and mass derivations of the Local Group (LG), and the present density of the Universe from timing arguments (Peebles 1994)? Is the dipole in the Cosmic Microwave Background Radiation (direction and amplitude) entirely explained by the gravity on the LG from the irregular mass/galaxy distribution? As the nearest galaxies ($v<300$ \kms) generate 20\% of the total dipole moment (Kraan-Korteweg 1989) nearby individual galaxies are equally important as massive groups, clusters and voids. Is the mass overdensity in the Great Attractor (GA) region -- postulated from a large-scale systematic flow of galaxies towards ($\ell,b,v)\sim(320\degr,0\degr,4500$\kms) (Kolatt \etal\ 1995) -- in the form of galaxies, hence does light trace mass? Does the Supergalactic Plane, other superclusters, walls and voids connect across the Milky Way and might other large-scale structures (LSS) have gone undetected due to this 'zone of avoidance'? | The combination of the complementary multiwavelength surveys allow a new probing of LSS in the 'former' ZOA. The \HI\ surveys are particularly powerful at the lowest latitudes. But future merging of ZOA data with catalogs outside the ZOA will have to be done with care to obtain 'unbiased' whole-sky surveys. From the sensitivity attained with the first 2 scans of the ZOA MB-survey it can be maintained that no Andromeda or other \HI-rich Circinus galaxy is lurking undetected behind the extinction layer of the southern Milky Way.\\ {\sl Acknowledgements} --- The help of the HIPASS ZOA team members R.D. Ekers, A.J. Green, R.F. Haynes, P.A. Henning, R.M. Price, E. Sadler, and L. Staveley-Smith is gratefully acknowledged. | 98 | 4 | astro-ph9804066_arXiv.txt |
9804 | astro-ph9804072_arXiv.txt | Multicolour images of the starbursting metal poor blue compact galaxy ESO~338-IG04 have been obtained with the {\it Wide Field Planetary Camera 2} on board the {\it Hubble Space Telescope}. In the images we find numerous point-like sources concentrated towards the main body of the galaxy, which we identify as globular cluster candidates. We show that these objects are physically associated with the galaxy and that they are spatially extended. Given their high intrinsic luminosities, these objects cannot be individual stars. Using photometric evolution models we show that the objects constitute a rich population of massive star clusters with ages ranging from a few Myr to $\sim$ 10 Gyr, and masses ranging from $10^4$ to more than $10^7 {\cal M_{\odot}}$. There are peaks in the age distribution of the clusters: one with objects $\le30$ Myr, one at $\sim 100$~ Myr, one at $\sim 600$~ Myr, one to two at $2.5-5$ Gyr and one at $\sim10$ Gyr. The youngest objects are predominantly found in the crowded starburst region. They have properties which agree with what is expected for young globular clusters, although it cannot be excluded that some of them may be dissolved or disrupted. For objects older than a few times 10 Myr, the only plausible explanation is that these are globular clusters. The galaxy presently appears to be involved in a merger, which is the probable cause of the present globular cluster formation. The presence of a numerous intermediate age (2.5 to 5 Gyr) population of globular clusters, suggests that a previous merger might have occurred. As the starburst fades, this galaxy will become very rich in globular clusters. Transforming all objects to an age comparable to that of Milky Way globular clusters reveals a luminosity function similar to the Galactic. We suggest that this galaxy is the result of a merger between a dwarf elliptical and a gas rich dwarf. The possibility of dating the globular clusters offers an efficient way of studying the history of violent star formation in this and similar galaxies. | Globular clusters (GCs) are generally old stellar systems and are believed to be the first objects to form in the process of the formation of a galaxy. This is supported by age estimates from observations of GCs in the Milky Way and other nearby galaxies. The Large Magellanic Cloud (LMC) is however known to host several young blue "populous" clusters, which could be young globular clusters. In the recent years blue globular cluster candidates have been found in some interacting/merging galaxies such as NGC~3597 (Lutz 1990, Holtzman et al. 1996), NGC~7252 (Whitmore et al. 1993), "The Antennae" (Whitmore \& Schweizer 1995) and NGC~3921 (Schweizer et al. 1996). This has strengthened the idea that GCs can form not only when galaxies form, but also when galaxies are reformed in the process of galaxy mergers (Schweizer 1986, Ashman \& Zepf 1992, Whitmore 1996). This could possibly also circumvent the problem, faced by the idea that elliptical galaxies form by the merging of late type galaxies, that ellipticals have higher specific frequencies of globular clusters (van den Bergh 1984, 1994). Understanding how and when GCs can form will thus not only help us in understanding these systems, but also aid in understanding galaxy formation and evolution. Studies of actively star-forming galaxies, e.g. Henize 2-10 (Conti et al. 1994), NGC~1244 (Hunter et al. 1994), NGC~1705 (O'Connell et al. 1994), NGC~1569 (De Marchi et al. 1997) and M82 (O'Connell 1995) have revealed the presence of star clusters with luminosities comparable to R126, the central cluster of 30 Doradus in the LMC, and higher. Meurer et al.(1995) used the HST for an ultraviolet study of nine starburst galaxies and found bright "super star clusters" in all. These clusters are generally bluer than similar objects found in interacting galaxies, and are preferentially found in the central regions of the starbursts. It has been proposed (e.g. Conti et al. 1994) that these blue objects might be forming globular clusters. This is however still an open question since it is not clear if these systems will survive as gravitationally bound systems. Uncertainties in the value of the (universal ?) stellar initial mass function (IMF) makes the mass estimates, mainly based on ultraviolet data, highly uncertain, especially since some studies only includes one spectral bandpass. Many of these galaxies appear to be involved in some form of interaction. One could interpret these observations as evidence that massive star clusters can form in regions of very active star formation, such as giant extragalactic HII-regions (GEHRs) like 30 Doradus. Kennicutt and Chu (1988) made a statistical investigation of data on extragalactic blue populous clusters and giant HII-regions. They concluded that the young clusters in LMC are not luminous enough to evolve into globular clusters comparable to the massive galactic ones, and that far from all GEHRs will produce young populous clusters that could become GCs. Blue compact galaxies (BCGs) are characterised by their blue colours ($B-V \le 0.5$), strong nebular emission lines, indicative of the formation of relatively hot (massive) stars, and low chemical abundances. The derived star formation rates could, considering the gas supply, only be sustained for a small fraction of a Hubble time. This together with the low metallicities (with IZw18 being the extreme in this sense) once lead to the idea that BCGs might be genuinely young objects now experiencing their first star formation epoch (Sargent and Searle 1970). Now, most BCGs are believed to be old, experiencing recurrent bursts of star formation intervened by long quiescent periods. Still we do not yet understand what triggers these bursts of star formation. One possible trigger mechanism is tidal or direct interactions with companion galaxies or gas clouds. In this paper we will present conclusive evidence for young and old globular clusters in the blue compact galaxy ESO~338-IG04. Section 2 describes the observations and Sect. 3 describes how photometry was performed on the globular cluster candidates. In Sect. 4 we show that the objects are physically associated with the galaxy and that they are spatially resolved. In Sect. 5 we discuss how the observed photometric properties can be interpreted in terms of age and mass of the objects. Section 6 includes a further discussion on the nature of the cluster candidates, and Sect. 7 contains the conclusions. \subsection{General properties of the target galaxy} ESO~338-IG04, also known as Tol~1924-416, resides at a distance of 37.5 Mpc ($v_{hel}=2813$ km~s$^{-1}$, $H_0= 75$ km s$^{-1}$ Mpc$^{-1}$; this value will be used throughout the rest of the paper) at the celestial coordinates $\alpha_{1950}=19^{\rm h}~24^{\rm m}~30^{\rm s}~ \delta_{1950} = -41\degr~34\arcmin~00\arcsec$. It is intrinsically bright ($M_V=-19.3$) and blue ($B-V=0.4$), (Bergvall and \"Ostlin 1998). The oxygen abundance is $12 \%$ of the solar value (Bergvall, 1985). It's physical size is small, $8.5 \times 4.5$ kpc measured at $\mu _V=25 {\rm mag}~{\rm arcsec}^{-2}$, the non dwarfish luminosity being due to the active star burst. It has an integrated HI-mass of $3\times10^9 {\cal M_{\odot}}$ (\"Ostlin et al. 1997b). The optical velocity field of ESO~338-IG04 show rotation aligned with it's apparent major axis, though with several large scale irregular features, and it is most easily understood as a merger between two galaxies or a galaxy and a gas cloud (\"Ostlin et al. 1997a). This is supported by the optical tail, extending towards the east. It has a spectroscopically confirmed companion galaxy (Bergvall unpublished; \"Ostlin et al. 1997a) which lies at a projected distance of 70 kpc, see Fig. 3. The companion is a somewhat fainter ($M_V = -17.9$), star forming galaxy which shows regular rotation (\"Ostlin et al. 1997a). Bergvall noticed, in ground based images, an apparent concentration of faint blobs around the galaxy, which he interpreted as globular cluster candidates. We have re-observed this galaxy with the HST and will in the subsequent sections show that there is substantial evidence that this galaxy has formed massive globular clusters at several epochs. \begin{figure} \picplace{8.5cm} \caption[]{A Digitized Sky Survey (DSS) image (obtained with $Skyview$) showing the target galaxy and its companion. North is up and east is left. The total angular size of the field is ~$9\arcmin \times 9\arcmin$. ESO~338-IG04 is at the upper left, the companion at the lower right. } \end{figure} In July 1997, we obtained spectra of a few of the brightest globular cluster candidates, using the ESO New Technology Telescope (NTT), at La Silla, Chile. The results from this spectroscopic study will be presented in a future publication. | Multi-colour photometry with HST/WFPC2 of the metal poor blue compact galaxy ESO~338-IG04 (Tol~1924-416) has revealed a rich population of faint point-like sources in and surrounding the galaxy. Aperture photometry has been performed on these, and the photometric results have been transformed into ages and masses using a spectral evolutionary synthesis model. Special care was taken to assure that the sources are physically associated and not chance projected background or foreground sources. The results can be summarised as follows: \begin{enumerate} \item The objects discussed in this paper have absolute magnitudes, ~$M_v$, in the range -7 to -15 and $(v-i)$ colours ranging from less than -1 to almost 2. \item The vast majority of the found objects are spatially resolved star clusters, physically associated with the galaxy. The outer objects follow the luminosity distribution of the galaxy closely (Fig. 9). \item The number of interlopers, i.e. foreground stars, background galaxies and super giants in the target galaxy, in the presented sample is conservatively estimated to be less than ten objects. This leaves us with more than 112 detected star clusters in ESO~338-IG04. \item Using the photometric evolution model we show that the objects, in general, are well fitted by Salpeter and Miller-Scalo IMFs, and that the resulting age distribution is insensitive to the adopted IMF. The ages of the clusters range from a few Myr to more than 10 Gyr. \item The objects have inferred masses ranging from $10^4$ to more than $10^7 {\cal M_{\odot}}$. \item The above listed properties show that the objects are massive globular clusters of varying age. \item There are several peaks in the globular cluster age distribution. This shows that there have been several globular cluster forming events in this galaxy. These appear to have occurred ~$\sim$ 11 Gyr ago, 2.5-5 Gyr ago, 600 Myr ago, 100 Myr ago and "now". The present event has lasted for a few times 10 Myrs. \item More than half of the total expected number of GCs have an age in the range 2 to 5 Gyr. This peak consists of two subpopulations with either different age ($\Delta_{\rm AGE} \sim 2.5$~Gyr) or metallicity ($\Delta_{\rm [Fe/H]} \sim 1$~dex). \item The specific frequency, $S_N$, is high in this galaxy. Taking into account that the fading starburst will decrease the luminosity of the galaxy by 1.5 magnitudes in 1 Gyr, this galaxy will be come very rich in globular clusters, even if the newly formed GCs don't survive. The predicted specific frequency of globular clusters is much larger than for late type galaxies and comparable to that for giant ellipticals. \item We suggest that a merger is responsible for the current starburst and cluster formation. In view of the high specific frequency of old globular clusters suggests that an elliptical galaxy is a main ingredient in this system. To provide gas for star and cluster formation a gas rich dwarf must be the other main ingredient. A major formation event 2.5-5 Gyr ago suggests that a merger involving one gas rich component also occurred at that time. \item In view of the high specific frequency and the present merger, we suggest that this galaxy will evolve into a moderately luminous elliptical galaxy, unless disturbed by possible future interactions. \item This investigation for the first time gives strong support for a rich population both old and newly formed GCs in a metal poor blue compact galaxy. \item Detection and dating of GCs in moderately distant BCGs is feasible using the HST and offers an efficient way of studying the violent star formation history, if also the effects of metallicity and extinction can be handled. \end{enumerate} | 98 | 4 | astro-ph9804072_arXiv.txt |
9804 | astro-ph9804244_arXiv.txt | Recent work suggests that rich clusters of galaxies commonly have large populations of dwarf (ie. low luminosity) members, that is their luminosity function (LF) turns up to a steep slope at the faint end. This population, or more particularly the relative numbers of dwarfs to giants, appears to be very similar for clusters of similar morphology, but may vary between cluster types. We have previously suggested that dwarfs may be more common in less compact, spiral rich clusters. Similarly we have found evidence for population gradients across clusters, in that the dwarf population appears more spatially extended. In the present paper we summarise the current evidence and propose, in analogy to the well-known morphology - density relation, that what we are seeing is a dwarf population - density relation; dwarfs are more common in lower density environments. Finally we discuss recent semi-analytic models of galaxy formation in the hierarchical clustering picture, which may give clues as to the origin of our proposed relation. | Much recent work has been devoted to the question of the galaxy luminosity function (LF) within rich clusters, particularly with regard to the faint end which has become accessible to detailed study through various technical and observational improvements (see e.g., Driver et al. 1994; Biviano et al. 1995; Bernstein et al. 1995; Mohr et al. 1996; Wilson et al. 1997; Smith, Driver \& Phillipps 1997 = Paper I; Trentham 1997a,b). For the most part these studies concur that the LF becomes steep (Schechter (1976) slope $\alpha \leq - 1.5$) faintwards of about $M_{B} = -17.5$ or $M_{R} \simeq -19$ (for $H_{0}$ = 50 km s$^{-1}$ Mpc$^{-1}$), and Paper I suggested that such a dwarf rich population might be ubiquitous. In a subsequent paper (Driver, Couch \& Phillipps 1997a = Paper III) we have examined the luminosity distribution in and across a variety of clusters, examining the possible dependence of the dwarf population (in particular the ratio of dwarfs to giants) on cluster type and position within the cluster. In the present paper we summarise the evidence to date for the (dis)similarity of the dwarf population in different environments. | As with the corresponding morphology density relation for giant galaxies, the cause of our population - density relation could be either `nature' or `nurture', ie. initial conditions or evolution. Some clues may be provided by the most recent semi-analytic models of galaxy formation, which have been able to account in a general way for the excess of (giant) early type galaxies in dense environments (e.g., Baugh, Cole \& Frenk 1996). The steep faint end slope of the LF appears to be a generic result of hierarchical clustering models \footnote{ And was considered a problem until observational evidence for steep LFs increased!} (e.g., White \& Frenk 1991; Frenk et al. 1996; Kauffmann, Nusser \& Steinmetz 1997), so is naturally accounted for in the current generation of models. The general hierarchical formation picture envisages (mainly baryonic) galaxies forming at the cores of dark matter halos. The halos themselves merge according to the general Press-Schechter (1974) prescription to generate the present day halo mass function. However the galaxies can retain their individual identities within the growing dark halos, because of their much longer merging time scales. The accretion of small halos by a large one then results in the main galaxy (or cluster of galaxies, for very large mass halos) acquiring a number of smaller satellites (or the cluster gaining additional, less tightly bound, members). Kauffmann et al. (1997) have presented a detailed study of the distribution of the luminosities of galaxies expected to be associated with a single halo of given mass. The LFs are somewhat disjoint owing to the specific halo masses modelled; especially for the low mass halo there is a preferred luminosity for the central galaxy plus a tail to lower luminosities. For a realistic mix of halo masses, these would no doubt be smoothed to look more like conventionally observed LFs. Nevertheless, we can still easily compare the numbers of dwarf galaxies per unit giant galaxy luminosity (rather than the amplitude of the giants' LF) between halos of different mass. The Kauffmann et al. models mimic a "Milky Way system" (halo mass $5 \times 10^{12} M_{\odot}$), a sizeable group (halo mass $5 \times 10^{13} M_{\odot}$) and a cluster mass halo ($10^{15} M_{\odot}$). Using their figure 2 (which also emphasises the identical faint end slopes predicted for all the different environments), we choose to quantify the number of dwarfs by $N_{-18}$, the number of dwarfs per system in the $M_{B} = -18$ bin. Because of the very similar slopes, the choice of bin or range of bins does not affect our conclusions, so this is the equivalent of the total number of dwarfs used in Figure 1. To quantify the giant population we choose the total light of galaxies of $M_{B} = -20$ or brighter, in units of $L_{*}$ galaxies (taking $M_{B}^{*} = -21$), which we call $N_{-21}$. Using this definition, rather than the actual number of galaxies brighter than some value (as in our observational data) allows for the discretization of the LFs for small halos. The results are summarized in Table 1. The ratio of these two values $N_{-18}$ and $N_{-21}$ then quantifies the relative dwarf galaxy populations. Roughly speaking, for smooth LFs with a shape similar to that observed, we should multiply these values by about 5, giving a range from about 1 to 3, to compare with our observational DGRs. \begin{table*} \begin{center} \begin{tabular} {lccccccccr} Halo Mass ($M_{\odot}$) & $N_{-18}$ & $N_{-21}$ & $N_{-18}/N_{-21}$ \\ \tableline $5 \times 10^{12}$ & 0.2 & 0.34 & 0.6 \\ $5 \times 10^{13}$ & 2.2 & 3.8 & 0.6 \\ $1 \times 10^{15}$ & 40 & 190 & 0.2\\ \end{tabular} \end{center} \caption{Dwarf numbers as a function of halo mass. \label{tbl-1}} \end{table*} We see that the Milky Way and small group halos have similar numbers of dwarf galaxies per unit giant galaxy light, whereas the dense cluster environment has a much smaller number of dwarfs for a given total giant galaxy luminosity. Thus the predictions of the hierarchical models (which depend, of course, on the merger history of the galaxies) are in general agreement with our empirical results if we identify loose clusters and the outskirts of rich clusters with a population of (infalling?) groups (cf. Abraham et al. 1996), whereas the central dense regions of the clusters originate from already massive dark halos. By inputting realistic star formation laws etc., Kauffmann et al. can further identify the galaxies in the most massive halos with old elliptical galaxies, and those in low mass halos with galaxies with continued star formation. This would imply the likelihood that our dwarfs in low density regions may still be star forming, or at least have had star formation in the relatively recent past (cf. Phillipps \& Driver 1995 and references therein). Note, too, that these galaxy formation models would also indicate that the usual (giant) morphology - density relation and our (dwarf) population - density relation arise in basically the same way. Finally, we can see that if these models are reasonably believable, then we need not expect the field to be even richer in dwarfs than loose clusters; the dwarf to giant ratio seems to level off at the densities reached in fairly large groups. To summarise, then, we suggest that the current data on the relative numbers of dwarf galaxies in different clusters and groups can be understood in terms of a general dwarf population versus local galaxy density relation, similar to the well known morphology - density relation for giants. Low density environments are the preferred habitat of low luminosity galaxies; in dense regions they occur in similar numbers to giants, but at low densities dwarfs dominate numerically by a large factor. This fits in with the general idea that low luminosity galaxies are less clustered than high luminosity ones (particularly giant ellipticals). Plausible theoretical justifications for the population - density relation can be found within the context of current semi-analytic models of hierarchical structure formation. | 98 | 4 | astro-ph9804244_arXiv.txt |
9804 | gr-qc9804086_arXiv.txt | Inflation of cosmic gauge and global strings is investigated by numerically solving the combined Einstein and field equations. Above some critical symmetry-breaking scales, the strings undergo inflation along the radial direction as well as the axial direction at the core. The nonsingular nature of the spacetimes around supercritical gauge and global strings is discussed and contrasted to the singular static solutions that have been discussed in the literature. | Cosmic strings are linelike topological defects that may form as a result of a phase transition in the early universe. If a string is associated with a magnetic field, it is called a gauge string, otherwise it is a global string. They have attracted much attention because of their cosmological importance: deficit angle in the spacetime geometry and a candidate for the seed of structure formation in the early universe. It was proposed that topological defects can inflate if the symmetry-breaking scale satisfies $\eta \gtrsim \eta_c \sim {\cal O}(m_p)$ in Refs.~\cite{Linde,Vilenkin}. This was later verified in numerical simulations by Sakai {\it et al.}~\cite{Sakai}. They found, in particular, that the critical value of $\eta$ for domain walls and global monopoles is $\eta_c \simeq 0.33m_p$. Then what about cosmic strings? There is no reason why we exclude cosmic strings out of the topological inflationary category. Recently, it was numerically proved that a (2+1) dimensional gauge string can inflate by de Laix {\it et al.}~\cite{Tanmay}. In this paper, we shall numerically solve the combined Einstein and field equations for a gauge and a global string in (3+1) spacetime dimensions. For the gauge string, we find that the core inflates if $\eta \gtrsim 0.25m_p$ with unit winding number in the critical coupling case (Bogomol'nyi limit). For the global string, $\eta \gtrsim 0.23m_p$. The critical values decrease as the winding number increases. For the gauge string, the critical value also decreases slightly as the coupling of the gauge field to the scalar field becomes weaker than the self coupling of the scalar field. The asymptotic spacetime of a gauge string is known to be conical~\cite{VilenkinC}. This spacetime exhibits a deficit angle $\Delta =8\pi G\mu$, where $\mu\sim \eta^2$ is the mass per unit length of the string. When the symmetry-breaking scale is sufficiently large, the deficit angle exceeds $2\pi$ and analyses of the static solution show that the spacetime possesses a physical singularity outside the core of the string~\cite{Gott,Ortiz,Laguna}. However, from our numerical results we know that supermassive strings are dynamical and undergo inflation at the core. Therefore, we believe that the static treatment of supermassive strings loses its validity and that we should treat them in a time-dependent way. For global strings, the singularity exists regardless of the symmetry-breaking scale. Many people have tried to find a static solution of a global string and they found that there also exists a physical singularity outside the core of the string~\cite{Cohen,Sikivie,Gregorys,Gibbons}. What was suggested to remove this singularity is again a time-dependent treatment of the string. Gregory~\cite{Gregoryt} introduced a specific metric which has an axial time dependence and showed that this spacetime is nonsingular. In our work, we follow the evolution of supermassive gauge and global strings in a general time-dependent metric and show that no singularity develops in the spacetimes around the strings. In the next section, we solve the Abelian Higgs model of a gauge string and discuss its inflation and spacetime geometry. Sec.~III is devoted to global strings. Our conclusions are summarized in Sec.~IV. In Appendix, we show the equations in detail and the numerical algorithms. | We have investigated inflation in cosmic strings. In the core region, the strings undergo inflation radially as well as axially when $\eta \gtrsim \eta_c$. With unit winding number ($n=1$) the critical values for inflation were found to be $\eta_c \approx 0.25m_p$ for a gauge string in the Bogomol'nyi limit ($\beta=1$) and $\eta_c =0.23m_p$ for a global string. The critical values decrease as $n$ and $\beta$ increase. We have explained this $\eta_c$ variation in terms of the core size of defects. The core of defects inflates when its size becomes bigger than the horizon scale: for larger $n$ and $\beta$, strings have bigger cores, and the global string has a bigger core than the gauge string for a given $\eta$. Regardless of the symmetry-breaking scale $\eta$, around the center of defects the de Sitter expansion is established since the scalar field stays about the top of the potenital ($\phi\approx 0$). However, this is not sufficient for the cores of defects to inflate. Inflation requires another condition which is the core size being comparable to the horizon scale so that the core can be dynamical due to the gravitational effect. Or equivalently, the potential $V(\phi)$ needs to be flat enough at $\phi\approx 0$ so that the field $\phi$ can spend enough time about the top of the potential. For this condition to be satisfied, the symmetry-breaking scale $\eta$ needs to be sufficiently large. This description also explains why we have somewhat lower critical values of $\eta$ for strings than those for domain walls and global monopoles ($\eta_c\approx 0.33m_p$). Strings have bigger cores at the same symmetry-breaking scale than the other defects. For supermassive gauge strings and all scale global strings, we have had troublesome physical singularities outside the core when we treat them in a static way. The elegant exit to nonsingular spacetimes is to introduce a time-dependent treatment. From the numerical simulations we could show that there is no singularity developing around time-dependent supermassive strings. | 98 | 4 | gr-qc9804086_arXiv.txt |
9804 | astro-ph9804008_arXiv.txt | }[1]{{\footnotesize \noindent {\bf Abstract} #1 \\}} \renewcommand{\author}[1]{\subsubsection*{#1}} \newcommand{\address}[1]{\subsubsection*{\it#1}} \setlength{\textheight}{20cm} \setlength{\textwidth}{13.5cm} \begin{document} \chapter*{Galactic Environments of the Sun and Cool Stars\footnote{To be published in {\it Planetary Systems -- The Long View}, eds. L. M. Celnikier and J. Tran Than Van, Editions Frontieres, 1998}} \author{Priscilla C. Frisch} \address{University of Chicago, Dept. Astronomy and Astrophysics, 5640 S. Ellis Ave., Chicago, IL 60637} \abstract{ The importance of understanding the current and historical galactic environments of cool stars is discussed. The penetration of interstellar gas into a stellar astrosphere is a function of the interaction of the star with the interstellar cloud surrounding the star, and this factor needs to be understood if an efficient search for life-bearing planets is to be made. For the Sun, both current and historical galactic conditions are such that if a solar wind were present, it would have excluded most inflowing interstellar matter from the inner regions of the heliosphere for the past few million years. Variations in heliosphere size over the recent historical path of the Sun are estimated, along with estimates of astrosphere sizes for selected nearby stars. Considering only possible effects due to encounters with interstellar clouds, stable planetary climates are more likely for inner than outer planets.} | The Sun moves through space at a velocity of about 17 pc per million years. This motion, combined with interstellar cloud motions driven by stellar evolution, yield a constantly changing galactic environment for the Sun and solar system. This environment affects the interplanetary environments of both outer and inner planets in the solar system, including Sun--Earth coupling mechanisms. By analogy, the interactions between other cool stars and the galactic environment of that star needs to be understood as part of the process of identifying planets conducive to ``higher'' life forms. The interstellar cloud surrounding the Sun at this time (known as the ``local interstellar cloud'', LIC), is warm, low density, and partially ionized: T$\approx$7,000 K, n(H$^{\rm \circ})$$\approx$0.2 cm$^{- 3}$, and n(e$^{-}$)$\approx$0.1 cm$^{-3}$. The standard assumption for diffuse interstellar clouds is that n(p$^{+}$)=n(e$^{-}$). On the scale of typical cloud densities, the LIC is rather tenuous, and notably lower density than the 1 au solar wind density (see Fig. \ref{densities}). This accounts for the ability of the solar wind today to exclude most interstellar material from 1 au. \begin{figure}[ht] \vspace*{4in} \special{psfile=density.eps voffset=-107 vscale=60 hscale=60 } \caption{Typical densities for material in our Galaxy. \label{densities}} \end{figure} In this paper the basis for understanding the relation between the properties of stellar wind envelopes around cool star systems and the physical properties of the surrounding interstellar clouds are examined. The author believes that the historical galactic environment of a star would have a direct impact on the stability of planetary atmospheres, and therefore on the distribution of intelligent life forms. This conclusion rests partly on the observation that the solar system has been in a region of space virtually devoid of interstellar matter over the past several million years \cite{fy86}, \cite{journey}. Additional reviews on interstellar matter (ISM) within the solar system can be found in the book {\it The Heliosphere in the Local Interstellar Medium} \cite{rudi}. For more information on the properties of local ISM (LISM) see \cite{fr95}. For more information on nearby G-star space motions and environments, and the use of astrospheres as a test for interstellar pressure, see \cite{fr93}. | 98 | 4 | astro-ph9804008_arXiv.txt |
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9804 | astro-ph9804187_arXiv.txt | We evaluate in a homogeneous way the optical masses of 170 nearby clusters ($z\le 0.15$). The sample includes both data from the literature and the new ENACS data (Katgert et al. 1996, 1998). On the assumption that mass follows the galaxy distribution, we compute the masses of each cluster by applying the virial theorem to the member galaxies. We constrain the masses of very substructured clusters (about $10\%$ of our clusters) between two limiting values. After appropriate rescaling to the X-ray radii, we compare our optical mass estimates to those derived from X-ray analyses, which we compiled from the literature (for 66 clusters). We find a good overall agreement. This agreement is expected in the framework of two common assumptions: that mass follows the galaxy distribution, and that clusters are not far from a situation of dynamical equilibrium with both gas and galaxies reflecting the same underlying mass distribution. We stress that our study strongly supports the reliability of present cluster mass estimates derived from X-ray analyses and/or (appropriate) optical analyses. \vspace*{6pt} \noindent {\em Subject headings: } galaxies: clusters: general - galaxies: distances and redshifts - X-rays: galaxies - cosmology: observations. | The knowledge of the properties of galaxy clusters plays an important role in the study of large scale structure formation. In particular, the observational distribution of the abundance of galaxy clusters as a function of their mass places a strong constraint on cosmological models (e.g., Bahcall \& Cen 1993; Borgani et al. 1997; Gross et al. 1998; White, Efstathiou \& Frenk 1993). Moreover, recent studies stress the need for having reliable estimates of cluster masses to constrain the ratio between the baryonic to total mass and the consequent value of $\Omega_0$ (e.g., White \& Frenk 1991; White et al. 1993b). Indeed, the estimate of cluster masses is not an easy task, in spite of the various methods which are available. The application of the virial theorem to positions and velocities of cluster member galaxies is the oldest method of cluster mass determination (e.g., Zwicky 1933). More recent methods are based on the dynamical analysis of hot X-ray emitting gas (e.g., Cowie, Henriksen, \& Mushotzky 1987; Eyles et al. 1991) and on gravitational lensing of background galaxies (e.g Grossman \& Narayan 1989). Mass estimates derived from the dynamical analysis of gas or member galaxies which are based on the Jeans equation or its derivations, such as the virial theorem, assume that clusters are systems in dynamical equilibrium (e.g., Binney \& Tremaine 1987). This assumption is not strictly valid; in fact, although clusters are bound galaxy systems, they have collapsed very recently or are just now collapsing, as is suggested by the frequent presence of substructures (e.g., West 1994). However, some analyses suggest that the estimate of optical virial mass is robust against the presence of small substructures (Escalera et al. 1994; Girardi et al. 1997a; see also Bird 1995 for a partially different result), although it is affected by strong substructures (e.g., Pinkney et al. 1996). Similar results come from studies based on numerical simulation (e.g., Schindler 1996a; Evrard, Metzler, \& Navarro 1996; Roettiger, Burns, \& Loken 1996) for X-ray masses estimated with the standard $\beta$-model approach (Cavaliere \& Fusco Femiano 1976), although some authors have claimed there is a systematic mass underestimation (e.g.; Bartelmann \& Steinmetz 1996). Dynamical analyses based on galaxies have the further drawback that the mass distribution or (alternatively) the velocity anisotropy of galaxy orbits should be known a priori. Unfortunately, the two quantities cannot be disentangled in the analysis of the observed velocity dispersion profile, but only in the analysis of the whole velocity distribution which, however, requires a large number of galaxies (of the order of several hundreds; e.g. Dejonghe 1987; Merritt 1988; Merritt \& Gebhardt 1994). Without some information from the relative distribution of dark and galaxy components, the virial theorem places only order-of-magnitude constraints on the total mass (e.g., Merritt 1987). The usual approach is to apply the virial theorem by assuming that mass is distributed like the observed galaxies (e.g., Giuricin, Mardirossian, and Mezzetti 1982; Biviano et al. 1993). This assumption is supported by several pieces of evidence coming both from optical (e.g., Carlberg, Yee, \& Ellingson 1997) and X-ray data (e.g., Watt et al. 1992; Durret et al. 1994; Cirimele, Nesci, \& Trevese 1997), as well as from gravitational lensing data, which, however, suggest a smaller core radius (e.g., Narayan \& Bartelmann 1997). The mass estimates derived from gravitational lensing phenomena are completely independent of the cluster dynamical status, but a good knowledge of cluster geometry is required in order to go from the projected mass to the cluster mass (e.g., Fort 1994). Moreover, strong lensing observations give values for the mass contained within very small cluster regions ($\lesssim$ one hundred of kpc) and weak lensing observations are generally more reliable in providing the shape of the internal mass distribution rather than the amount of mass (e.g., Squires \& Kaiser 1996). Up to now, few studies have dealt with wide comparisons between mass estimates obtained by different methods for the same cluster. Wu \& Fang (1996; 1997) found that masses derived from gravitational lensing analyses are higher than those from X-ray analyses by a factor of 2, but agree with those from galaxy analyses. Indeed, mass estimates from lensing seem to agree with X-ray estimates when clusters are relaxed (e.g., Allen 1997). However, Wu \& Fang's works concern only clusters which lie at moderate redshifts and show gravitational lensing phenomena which could be enhanced in the presence of substructures (e.g., Miralda-Escud\'e 1993; Bartelmann, Steinmetz, \& Weiss 1995). For nearby clusters, there is a trend to obtain larger masses from galaxy analyses than from X-ray analyses (e.g., Cowie et al. 1987; Mushotzky et al. 1995; David, Jones, \& Forman 1995), but acceptable agreement exists in some individual cases (e.g., for the Coma cluster, Watt et al. 1992). The classical approach of the virial theorem based on measurements of discrete velocities bears re-examining owing to the large new data sets which are now becoming available for nearby clusters, i.e. the ESO Nearby Abell Clusters Survey (ENACS) by Katgert et al. (1996, 1998). Moreover, the fair level of consistency among recent estimates of velocity dispersion of member galaxies resulting from different membership assignment procedures (cf. Fadda et al. 1996, hereafter F96, and Mazure et al. 1996) makes us confident of the robustness of our approach. The aim of this work is to obtain reliable mass estimates. These mass estimates will be used in the computation of the mass function of nearby clusters (Girardi et al. 1998). The paper is organized in the following manner. We describe the data sample and our selection procedure for cluster membership assignment in \S~2. We briefly describe the methods used to compute cluster masses by using member galaxies in \S~3. By assuming that mass follows the galaxy distribution, we compute virial mass estimates in \S~4, and we verify their consistency with the results of the Jeans equation in \S~5. The strongly substructured clusters are analyzed in \S~6. We compare our mass estimates with those derived from X-ray analyses in \S~7. We discuss our results in \S~8. We give a brief summary of our main results and draw our conclusions in \S~9. Unless otherwise specified, we give errors at the 68\% confidence level (hereafter c.l.) A Hubble constant of 100 $h$ \ks $Mpc^{-1}$ is used throughout. | The main points of this work may be summarized as follows: i) We evaluate in a homogeneous way the optical masses of 170 nearby clusters ($z\le 0.15$). This sample, which is the largest set of clusters up to now analyzed in the literature, includes both data from the literature and the new ENACS data (Katgert et al. 1996, 1998). ii) On the assumption that mass follows the galaxy distribution, we compute the masses of each cluster by applying the virial theorem to the member galaxies and we verify our results by using the Jeans equation. iii) Our mass estimates are smaller than previous optical estimates. This fact is due both to our better membership assignment procedure and to the application of the correction due to the presence of the surface term in the virial theorem (recently stressed by Carlberg et al. 1997a). iv) After appropriate rescaling to the X-ray radii, we compare our optical mass estimates to those derived from X-ray analyses, which we have compiled from the literature (for 66 clusters). We find a good overall agreement. v) The above agreement is expected on the basis of two common assumptions: a) that mass follows the galaxy distribution, b) that clusters are not far from a situation of dynamical equilibrium with both gas and galaxies reflecting the same underlying mass distribution. It should be pointed out that Carlberg et al. (1997a) have recently drawn similar conclusions for a sample of distant clusters (the CNOC sample). In particular, we find evidence for a galaxy distribution which is colder and less extended than the gas distribution. Several recent studies have casted doubts on cluster mass estimates and attempted to lower the cluster baryon fraction by reducing the cluster masses (e.g., Gunn \& Thomas 1996; Wu \& Fang 1996). We stress that our study strongly supports the reliability of present cluster mass estimates derived from X-ray analyses and/or (appropriate) optical analyses. Hence, it is even more difficult to reconcile present data with a $\Omega_0=1$ Universe (e.g. White et al. 1993b). Our cluster masses are suitable for statistical studies. In particular, we did not reject a priori those clusters with a poor number of selected members, which usually have a small mass, in order to avoid having a final cluster sample biased towards more massive systems. | 98 | 4 | astro-ph9804187_arXiv.txt |
9804 | astro-ph9804323_arXiv.txt | \he\ \lya\ \lm 304/\ha\ \lm 1640 emission lines are mainly produced by recombination, and their canonical ratio of $\sim 10$ may be a sensitive reddening indicator. We obtain the high S/N optical spectra of two quasars and combine them with the far-UV spectra that show the \he\ \lm 304 emission. For HS~1700+64, the \he\ \lm 1640 emission is not detected, and an upper limit to it sets the ratio greater than 20. This may not be inconsistent with the theoretical value when all observational uncertainties are taken into consideration. For Q0302-003, the ratio is very low, on the order of unity. The most plausible cause for such a low ratio is extinction in the EUV band by very fine grains of dust. Q0302-003 has a prominent narrow component of $\rm FWHM \sim 2000 \ \kms$ in its major emission lines, and it appears that reddening is associated only with the line-emitting region. We suggest that the geometry of the line-emitting region in high-z quasars resembles that in the low-luminosity active galaxies, with the presence of dust mostly in the outer part. | The first UV spectroscopic observation of a quasar (3C~273, \cite{davidsen77}) enabled a measurement of the \lya/\ha\ ratio in an active galactic nucleus (AGN). The low ratio of $\sim 1$ was in line with an independent study using the composite spectrum that was derived from various quasars (\cite{baldwin}), but was in sharp conflict with a canonical ratio of $\sim 10$ predicted by standard photoionization models (\cite{deo}). This ``\lya/\ha\ problem'' raised serious concerns about the validity of photoionization as the main line-emission mechanism in AGN and prompted extensive theoretical interest in the following years. Improved photoionization models (\cite{ferland}, and references therein) invoke large column densities and moderate degrees of ionization. In a partially ionized zone, the low escape probability for \lya\ photons makes a high population of excited states, and collisional excitation from these levels enhances Balmer lines. Calculations using a reasonable AGN continuum lead to an enhancement of Balmer lines by a factor of $\sim 2$, and hence may not fully explain the observed low \lya/\ha\ ratio. Another explanation introduces intrinsic reddening in the line-emitting region (\cite{dust}). The wavelength-dependent extinction reduces the intensities of observed UV lines, thus lowering the \lya/\ha\ ratio. The observed Pa$\alpha$/\ha\ ratio, however, appears to be too low for a straightforward full account as a reddening effect (\cite{rick}). Significant evidence exists for dust in the narrow-line region of Seyfert galaxies (\cite{deo}; \cite{n93}) and for a decrease in extinction with increasing luminosity toward quasars (\cite{cdz}; \cite{rudy}). However, \cite{wills} suggested that in intermediate-redshift quasars there may be significant reddening in the narrow-line region. Accurate assessments of the reddening effect depend on the use of good line pairs whose intrinsic ratios are fairly stable. Since singly ionized helium is hydrogenic, the \he\ I(\lya\ \lm 304)/I(\ha\ \lm 1640) ratio should therefore be the same as that for hydrogen. The \he\ emission is produced mainly by recombination, because its excitation level of 40 eV is considerably higher than the average thermal energy in the line-emitting region. The wavelength of \he\ \lm 304 emission coincides with that of O$^{++}$ transition $2p^2\ ^3P_2 - 2p\ 3d\ ^3P_2^0$, allowing Bowen fluorescence radiation (\cite{eastman}; \cite{netzer}). This radiation mechanism, however, does not appreciably affect the \he\ ratio itself. The \he\ \lm 304 emission should be extremely sensitive to reddening. While there are no hurdles in the theoretical aspects, it has taken some 20 years to advance from measuring this ratio in hydrogen to that in helium. | The \ratio\ ratio is quite different in these two quasars. Indeed the line profiles in Q0302-003 are narrower, making it easier to identify the weak \he\ \lm 1640 feature. The S/N level is high enough that a \he\ \lm 1640 feature should be detected even with a line width of $\sim$ 12 000 \kms. In Fig. 1 the profile of an assumed \he\ \lm 1640 feature is plotted, with an intensity 20\% of the \he\ \lm 304, which should have been detected. It appears that difference is not simply attributable to line widths. The intensity of \he\ \lm 304 emission is affected by the Lyman line and continuum absorption by numerous intervening absorbers along the line of sight. This can be corrected if a high-resolution spectrum at longer wavelengths yields a list of absorption lines. Our estimate, based on the statistical result of \cite{valley}, suggests an optical depth of 0.2 at a rest-frame wavelength of 300 \AA\ for a z=3.3 quasar. Therefore, this Lyman-Valley correction is not very significant. The UV and optical observations are not simultaneous, and both quasars are probably variable. A comparison of the UV spectra of HS 1700+64 obtained between 1991 and 1995 finds a significant discrepancy in flux level, and that between the optical spectra taken between 1994 and 1996 shows that the \lya\ equivalent width varies by a factor of 2. The \civ\ equivalent width of Q0302-003 has increased by $\sim 50\%$ as compared with the data of Sargent, Steidel \& Boksenberg (1989). Furthermore, the photometric quality of our optical spectra is questionable. A typical light loss with a small slit during an optical spectroscopic observation is $\sim 15\%$. These factors add uncertainties to the \ratio\ ratio. The derived line ratio is also subject to the reddening formulation. If we use the formula of Burstein \& Heiles (1978), this ratio would be even lower. We have carried out photoionization calculations (\cite{cloudy}) with various parameters. With a broad range of the density, column density, flux and shape of the ionizing continuum, the \ratio\ ratio varies within a narrow range between 9 and 11. It is therefore not practical to attribute the observed low value to special conditions in the line-emitting region. Note that a part of the \he\ \lm 304 emission may receive a contribution from \oiii\ \lm 305 emission that is produced by Bowen fluorescence mechanism (\cite{eastman}). If this were the case, the actual \ratio\ ratio would be even lower. The low \ratio\ ratio in Q0302-003 may signal internal reddening in the line-emitting region. Dust grains with dimensions of $\sim 3 \times 10^{-6}$ cm are believed to produce Galactic extinction (\cite{ext}) which generally follows a $1/\lambda$ law. If intrinsic extinction is produced by even smaller grains, and the wavelength dependence of the extinction law applies to wavelengths as short as 300~\AA, then an $E_{E - V} = 2.5$, where $E$ denotes a band around 300 \AA, is needed to explain the discrepancy between the observed and theoretical ratio of \ratio. This would translate into $E_{B-V} = 0.5$. If this is the case, significant presence of dust may be a reality in the broad-line region of some high-redshift quasars. The quantitative formalism should be more complicated than that. While the extinction curves between 1 $\mu$m and 1000~\AA\ can be approximated with a $1/\lambda$ law, very little is known about the extinction properties below 1000~\AA. \cite{hawkins} and \cite{martin} calculated the EUV extinction curve for graphite-silicate dust. Their results show {\em decreasing} extinction from 1000~\AA\ to 100~\AA. \cite{pei} suggested that a numerical formula can be applied to other galaxies, possibly to those of higher redshifts, without assuming a Galactic dust-to-gas ratio. If such extinction is real, the same extinction affects the intensity of other UV lines as well. The intrinsic hydrogen \lya/\ha\ ratio in these objects may actually be higher than the observed ones. Likely, the average ratio would be $ \sim 6$, closer than the theoretical value of $\sim 10$. This will, in turn, help understand the classical \lya/\ha\ puzzle. If significant reddening does exist in the quasar broad-line region, the observed \lya/\ha\ ratio may be corrected upward by an additional factor of $\sim 2$. Even for these two quasars, the I(\lm 1216)/I(\lm 304) ratio is very different. In HS~1700+64, this ratio is about 3, while in Q0302-003 it is about 30. Significant extinction in the EUV band can explain both the abnormal line ratios. The I(\lm 1216)/I(\lm 304) ratio in both objects is 30 or higher, consistent with photoionization models. Therefore, the likely cause for the low ratio in Q0302-003 is extinction in the broad line region by very small grains. It may not be coincidental that narrow line widths and possibly significant reddening are present in the same object. Seyfert 2 galaxies often show a higher degree of extinction (\cite{deo}), and the narrow-line region in Seyfert-1 galaxies generally exhibit a more significant reddening effect than the broad-line region (\cite{luc}). The spectrum of Q0302-003 shows a significant narrow component with FWHM $\sim 2000 $~\kms\ for major emission lines. Although this line width is not considered very narrow for Seyfert galaxies, it is for high-z quasars. Generally, narrow lines in high-z quasars (\cite{sargent}) are not as common as in Seyfert galaxies. For example, \cite{wills} found no detection of narrow line components in their radio-loud quasars of $\rm 0.26 < z < 0.77$. They suggested a significant reddening with $E_{B-V} \simeq 0.5$. We suggest that the geometry of the line-emitting region in high-z quasars resembles that in the low-luminosity active galaxies, with the presence of dust mostly in the outer part. Making an analogy of Seyfert galaxies and some low-redshift quasars, we suggest that Q0302-003 has a narrow-line region which contains a significant amount of dust. We suggest that the geometry of the line-emitting region in high-z quasars resembles that in the low-luminosity active galaxies, with the presence of dust mostly in the outer part (\cite{luc}). Does reddening apply to the EUV continuum? Recent studies (\cite{n95}; \cite{bechtold}) found that the \lya/\hb\ ratio ranges between about 1 and 40 and is approximately proportional to $f(1216)/f(4861)$, the ratio of continuum flux at adjacent points (\cite{bechtold}). Such a correlation may suggest a possible reddening effect that applies to both the continuum and lines. In such cases, the equivalent widths of concerned lines should be fairly constant. Given the significant difference in the equivalent widths of \he\ lines in our quasar samples, we see no compelling reason that the continuum emission from the central source is heavily reddened. Significant reddening would also affect the intensities of infrared lines, and future studies of these lines may provide additional evidence for fine dust in the quasar environment. | 98 | 4 | astro-ph9804323_arXiv.txt |
9804 | astro-ph9804115_arXiv.txt | The observed map of 1.809 MeV gamma-rays from radioactive $^{26}$Al (Oberlack et al, 1996) shows clear evidence of a Galactic plane origin with an uneven distribution. We have simulated the map using a Monte Carlo technique together with simple assumptions about the spatial distributions and yields of $^{26}$Al sources (clustered core-collapse supernovae and Wolf Rayet stars; low- and high-mass AGB stars; and novae). Although observed structures (e.g., tangents to spiral arms, bars, and known star-forming regions) are not included in the model, our simulated gamma-ray distribution bears resemblance to the observed distribution. The major difference is that the model distribution has a strong smooth background along the Galactic plane from distant sources in the disk of the Galaxy. We suggest that the smooth background is to be expected, and probably has been suppressed by background subtraction in the observed map. We have also found an upper limit of $1 M_{\sun}$ to the contribution of flux from low-yield, smoothly distributed sources (low-mass AGB stars and novae). | The gamma-ray created by the decay of $^{26}$Al to $^{26}$Mg was the first discovered Galactic gamma-ray line (Mahoney et al 1982). The $^{26}$Al nucleus decays by positron emission to the first excited state of $^{26}$Mg, which subsequently decays to the ground state emitting a 1.809~MeV gamma-ray. The mean lifetime of $^{26}$Al, $\tau = 1.05$ x $10^6$ years, makes the 1.809 MeV gamma ray line an excellent tracer for newly synthesized material released into the ISM over the last several million years. The main production mechanism of $^{26}$Al is proton capture on $^{25}$Mg. Astrophysical environments that can produce $^{26}$Al include hydrostatic H-burning in the convective cores of massive stars and the H-burning shells of intermediate mass stars, and explosive H burning in novae. The carbon and neon rich shells of massive stars are also a site for $^{26}$Al production both statically and explosively. In addition to its production, the fresh $^{26}$Al must be transported into the ISM before it decays in order to be observable. The explosive mechanisms present no problems, but the transport timescale in AGB stars is of similar order to the decay timescale causing a reduction in the amount of $^{26}$Al released into the ISM. The first map of Galactic 1.809 MeV gamma-ray emission from $^{26}$Al was published by Oberlack et al (1996) from COMPTEL data. This map has a 1$\sigma$ angular resolution of $1.6^{\circ}$, or $3.8^{\circ}$ FWHM. The map was produced using a Maximum-Entropy method after background subtraction. The 1.809 MeV gamma-ray map has several important characteristics, including the concentration of emission in the Galactic plane, a strong, irregular emission region toward the inner Galaxy, and a generally uneven, or clumpy, emission distribution. Along the Galactic plane there are several disconnected emission regions, some of which have been associated with O-B associations, spiral arm tangents, and the Vela SNR. Chen et al (1996) also identify several of the regions with spiral arm tangents. A recent and thorough review of the entire topic including observation, sources, and distribution of $^{26}$Al can be found in Prantzos and Diehl (1996). These observations can best be explained with sources that are spatially concentrated and rare. If the major sources of emission had small yields and a smooth Galactic distribution, the emission would be quite uniform. This is not seen in the published results (Oberlack et al 1996), which show large gaps along the Galactic plane between emission regions. We have therefore built a Monte Carlo model for the Galactic $^{26}$Al emission containing all potential astronomical sources. We have also allowed the most massive stars to form clusters that do not dissociate in the lifetime of those stars. We have made only the simplest of assumptions about Galactic structure, an exponential disk and a bulge. We have not attempted to represent any specific observed structures in the Galaxy. All non-uniformities arise from the random nature of the simulation. This produces a map that, to the eye, has a strong resemblance to the observations, with the exception of a persistent uniform background not found in the reduced observational data. We plan to use more statistically rigorous methods to measure the strength of this similarity in future work, as well as to consider some of the effects of non-uniform structure and the enhanced resolution of the INTEGRAL observatory. | 98 | 4 | astro-ph9804115_arXiv.txt |
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9804 | astro-ph9804265_arXiv.txt | The colossal power output of active galactic nuclei (AGN) is believed to be fueled by the accretion of matter onto a supermassive black hole. This central accreting region of AGN has hitherto been spatially unresolved and its structure therefore unknown. Here we propose that a previously reported `deep minimum' in the X-ray intensity of the AGN MCG$-$6$-$30$-$15, was due to a unique X-ray occultation event and that it probes structure of the central engine on scales $< 10^{14} \ \rm cm$, or $1.4\times 10^{-7}$ arcseconds. This resolution is more than a factor of $\sim 3\times 10^{6}$ greater than is possible with current X-ray optics. The data are consistent with a bright central source surrounded by a less intense ring, which we identify with the inner edge of an accretion disk. These may be the first direct measurements of the spatial structure and geometry of the accreting black-hole system in an active galaxy. We estimate a mass lower limit for sub-Eddington accretion of $3.1\times 10^{5} M_{\odot}$. If the ring of X-ray emission is identified with the inner edge of an accretion disk, we get mass upper limits of $1.9\times10^{8}$ and $9.1\times10^{8} M_{\odot}$ for a non-rotating and maximally rotating black hole respectively. We point out that our occultation interpretation is controversial in the sense that X-ray variability in AGNs is normally attributed to intrinsic physical changes in the X-ray emission region, such as disk or coronal instabilities. | \label{intro} The accretion of matter onto a supermassive black hole as a mechanism for fueling the output of active galactic nuclei (AGN) is a paradigm strongly supported by recent spectroscopic observations of the iron K$\alpha$ X-ray emission line (Tanaka \etal 1995; Yaqoob \etal 1995; Nandra \etal 1997a and references therein). The extreme Doppler and gravitational energy shifts of the line photons, together with the shape of the line are consistent with an origin in a disk rotating about a black hole (Fabian \etal 1989; Laor 1991). Strong gravitational redshifts, in which photon energies are changed by more than 10\%, occur only when matter approaches closer than $\sim 20$ gravitational radii ($= 20r_{g}; r_{g} \equiv GM/c^{2}$) from a compact object. However, it is not possible to directly map the physical structure of the system since the highest spatial resolution of X-ray optics technology is a factor $\sim 10^{6}$ too poor for even the closest AGN. Optical and radio observations provide greater resolution but the bulk of the emission at these wavelengths is not generated close enough to the central engine. So far, the highest spatial resolution observations, at radio wavelengths, have revealed a Keplerian disk in the AGN NGC 4258 down to only $\sim 60,000 r_{g}$ (Miyoshi \etal 1995; Maoz 1995). This still falls short by a factor $\sim 3000$ of mapping the black-hole region. The AGN MCG$-$6$-$30$-$15 ($z= 0.008$) was observed by the X-ray astronomy satellite {\it ASCA} ({\it Advanced Satellite for Astrophysics and Cosmology}; Tanaka, Inoue, and Holt 1994) for $\sim 4.2$ days on 23 July 1994. Results from this observation have already appeared in the literature, including the 0.5--10 keV lightcurve (Iwasawa \etal 1996, hereafter I96; Reynolds 1997; Yaqoob \etal 1997; see Figure 1) and a broad, asymmetric, variable iron K line with a strong red wing, consistent with a disk inclined at $\sim 30^{\circ}$ rotating about a black hole (Tanaka \etal 1995; I96). The X-ray luminosity exhibits erratic variability on all timescales down to $<50$ s (Matsuoka \etal 1990; Green, McHardy, and Lehto 1995; Reynolds \etal 1995; Nandra \etal 1997b). Causality arguments alone cannot put constraints on the size of the X-ray emission region since the high-frequency, lower amplitude variability may occur at localized regions of the source. Figure 1 shows an extended intensity dip at the end of the observation, from $\sim 3.3\times10^{5}$s to $\sim 3.6\times10^{5}$s. This feature has previously been dubbed as the 'deep minimum', or DM (I96). A closer inspection (Figure 2a) reveals a remarkable (albeit approximate) symmetry about the minimum luminosity. We propose that the dip was caused by an occultation of the X-ray source by optically-thick matter. This interpretation is controversial and 'non-standard', as X-ray variability in AGNs is normally attributed to intrinsic properties of the X-ray emission region, such as disk or coronal instabilities. However, so little is know about structure of the central engine in AGNs that the occultation scenario should be explored. The obscurer must be optically thick because the dip continuum spectrum only shows evidence for {\it weak} absorption, nowhere near enough to explain the observed intensity variation over the whole \asca bandpass (see Weaver \& Yaqoob 1998, hereafter WY98). The luminosity at the absolute minimum of the dip is $\sim 0.4$ of the pre-dip value and must represent persistent emission which has much smaller surface brightness than the primary source. Hereafter we will refer only to the primary X-ray emission, unless explicitly referring to the persistent emission. The proposed obscurer very likely hides the most compact and variable part of the X-ray source, since the usual rapid variability outside the dip is absent during the obscuration. Of course, it is possible that the dip is due to intrinsic variation of the X-ray source. However, the origin of X-ray variability in AGN is not understood. Models which come close to successfully reproducing the observable quantities obtained from AGN lightcurves are of the shot-noise or `rotating hot-spot' variety (e.g. Green \etal 1993; Bao and Abramowicz 1996 and references therein). However, the parameters of such models must be highly tuned in order to reproduce AGN power spectra. On the other hand, we show in this paper that a very simple-minded model can account for the temporal profile of the intensity dip in MCG $-$6$-$30$-$15 and briefly discuss the implications for AGN X-ray variability in general. | If accreting matter is exposed to the same UV/X-ray luminosity that we observe ($L \sim 4 \times 10^{43} \ \rm ergs \ s^{-1}$) then for gravitational infall to overcome outward radiation pressure requires the mass of the central black hole, $M_{\rm BH}$, to exceed $3 \times 10^{5} M_{\odot}$. Thus, if $r_{1}$ is the radius of the inner edge of the accretion disk, identified with the last stable orbit of matter, and $\Delta t$ is the time taken for the obscurer to traverse $r_{1}$ (i.e. $t_{3}-0.5[t_{2}+t_{1}$]), then $\kappa r_{g}=r_{1}$ where $\kappa=6$ or 1.24 for Schwarzschild or maximally rotating Kerr metrics respectively. But, $ r_{1}<c \Delta t$ and $r_{g} = 1.48 \times 10^{13} \ \rm (M_{\rm BH}/M_{\odot})$ cm, so $M_{\rm BH} < 1.9 \times 10^{8} M_{\odot}$, or $<9.1\times 10^{8} M_{\odot}$, for a Schwarzschild or Kerr black hole respectively. Assuming instead, the Keplerian velocity at the inner disk edge as the maximum velocity of the obscurer ($v/c = \sqrt{r_{g}/r} = \sqrt{1/\kappa}$) yields smaller mass upper limits of $M_{\rm BH}<7.7\times10^{7} M_{\odot}$ and $M_{\rm BH} <8.2\times10^{8} M_{\odot}$ for a Schwarzschild and extremal Kerr metric respectively. If the obscurer is at a distance, $d$, from the central source, $r_{1} = v \Delta t = c \sqrt{r_{g}/d}$ gives $d = c^{2}[\Delta t]^{2}/\kappa^{2}r_{g} = 6 \times 10^{15}[\Delta t]^{2}/[\kappa^{2} (M_{\rm BH}/M_{\odot})]$ cm. Thus, $M_{\rm BH} > 3 \times 10^{5} M_{\odot}$ implies $d<1.8\times 10^{16}\ \rm cm$ and $d<4.2\times10^{17}\ \rm cm$ for Schwarzschild and Kerr metrics respectively. The origin of the optically-thick blobs is unspecified but Guilbert and Rees (1988) presented some simple arguments for the existence of dense ($n > 10^{15} \rm \ cm^{-3}$), optically thick matter residing at the heart of accreting sources. The only independent estimate of the size of the `blobs' is that they should be much thicker than $10^{9}/n_{15} \ \rm cm$ ($n_{15}$ in units of $10^{15} \rm \ cm^{-3}$). The blobs must be optically thick even near their physical boundaries (i.e they must have fairly sharp edges), otherwise the dip profile would not be so well defined. Also, the blobs must be fairly stable, esepcially if they are created in the central region itself and 'propelled' up to high altitudes. The nature of the bright central X-ray source is intriguing. An inclined jet is unlikely, since even at $30^{\circ}$ the inflexions in the dip profile ($t_{2}$--$t_{3}$ and $t_{6}$-$t_{7}$) would have different durations. If the central X-ray source extracts its energy directly from the black hole then the metric is likely to be Kerr since energy cannot be extracted from a non-rotating black hole (Blandford and Znajek 1977; see also Ghosh and Abramowicz 1997). Occultations such as the one described here may occur frequently in AGN, but the relative sizes of the obscurer and source must be just right in order to observe such a clear event. Indeed the usual rapid variability or flicker, may be partly caused by the transit of optically thick bodies smaller than the source. We tested this hypothesis, again using a simple-minded model. Representing the bright central source as a circular disk with uniform emissivity (ignoring the ring due to its weaker emission), optically-thick blobs with ranges in radii (relative to the source) and velocities taken from Gaussian distributions were passed over the source. An additional parameter is required to specify the 'blob birth-rate' (i.e. the rate at which new blob trajectories are started). Such a model was used to produce predicted lightcurves for different model parameters. The power spectrum of each lightcurve was computed using the method of Papadakis and Lawerence (1993), omitting Poisson noise. In the range $\sim 10^{-5}$ Hz to $\sim 10^{-2}$ Hz, no preferred or 'universal' power-law spectral slope was found. It is possible to produce power-law spectra with slopes similar to those typically measured ($\sim -1$ to $-2$) but for most parameter values, the slopes are too steep. Thus, fine-tuning would be necessary to explain the `universal' power-law slopes found in the handful of AGN in which it can be measured (e.g. Lawerence and Papadakis 1993; Green \etal 1993). This is essentially because if the product of blob crossing-time and birth-rate is too large or too small, there will be no variability. The direct simulated lightcurves assume, of course that the {\it intrinsic} source intensity is constant, which almost certainly is not the case. Thus our model does not explain AGN variability in general but its effects are potentially important to consider in any model of AGN variability. We thank the \asca TEAM and mission operations at ISAS, Japan, for their efforts and hard work; Kim Weaver for her work on the dip spectrum, and Karen Leighly, Paul Nandra for some useful discussions. We also thank the anonymous referee. This research made use of archival data at the HEASARC, Laboratory for High Energy Astrophysics, NASA/Goddard Space Flight Center. | 98 | 4 | astro-ph9804265_arXiv.txt |
9804 | astro-ph9804053_arXiv.txt | We investigate different types of neutrino hot dark matter with respect to structure formation and anisotropies in the cosmic microwave background radiation (CMBR). The possibility of neutrino hot dark matter produced through the decay of a heavier neutrino by the process $\nu_H \to \nu_L + \phi$, where $\phi$ is a scalar particle, is discussed in detail. This type of dark matter can possibly be distinguished observationally from the standard neutrino dark matter by using new CMBR data from the upcoming satellite missions MAP and PLANCK. | 98 | 4 | astro-ph9804053_arXiv.txt |
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9804 | astro-ph9804029_arXiv.txt | The importance of the interstellar magnetic field is studied in relation to the evolutions of superbubbles with a three-dimensional (3D) numerical magnetohydrodynamical (MHD) simulation. A superbubble is a large supernova remnant driven by sequential supernova explosions in an OB association. Its evolution is affected by the density stratification in the galactic disk. After the size reaches 2--3 times the density scale-height, the superbubble expands preferentially in the $z$-direction. Finally it can punch out the gas disk (blow-out). On the other hand, the magnetic field running parallel to the galactic disk has an effect to prevent from expanding in the direction perpendicular to the field. The density stratification and the magnetic fields have completely opposite effects on the evolution of the superbubble. We present results of 3D MHD simulation in which both effects are included. As a result, it is concluded that when the magnetic field has a much larger scale-height than the density, even for a model that the bubble would blow out from the disk if the magnetic field were absent, the magnetic field with the strength of 5 $\mu$G can confine the bubble in $|z| \alt 300$ pc for $\simeq$20 Myr (confinement). In a model that the field strength decreases in the halo in proportion to $B \propto \rho^{1/2}$, the superbubble eventually blows out like a model of $B=0$ even if the magnetic field in the mid-plane is as strong as $B=5\mu$G. | A superbubble is a complex consisting of an OB association, surrounding X-ray emitting hot gas, and a corresponding HI hole/shell. Three examples are well-known in our Galaxy: Cygnus (Cash et al. 1980), Orion-Eridanus (Cowie, Songaila \& York 1979; Reynolds \& Ogden 1979), and Gum nebula (Reynolds 1976). Superbubbles are found as HI shells and holes in external galaxies, such as LMC (Meaburn 1980; Dopita, Mathewson \& Ford 1985), M31 (Brinks \& Bajaja 1986), M33 (Deul \& den Hartog 1990), M101 (Kamphuis, Sancisi \& van der Hulst 1991), and so on. Since the sizes of these objects are in the range of 100 pc -- 1kpc, this can not be explained by a single supernova explosion [the size of an ordinary supernova remnant (SNR) is $\alt 50$pc]. The amount of energy required for such a superbubble reaches $5\times 10^{51} {\rm erg} - 10^{54} {\rm erg}$ (Tenorio-Tagle \& Bodenheimer 1988). There are two models proposed for formation of superbubbles: (1) a large SNR driven by sequential supernova explosions in an OB association and (2) a complex formed by a collision of a high-velocity cloud and the galactic disk (Tenorio-Tagle 1991). Here, we confine ourselves to the first model and discuss the evolutions. For review papers of this field, see Tenorio-Tagle \& Bohdenheimer (1988), Spitzer (1990), Tomisaka (1991), Bisnovatyi-Kogan \& Silich (1995). Here, we summarize the evolution very briefly. After the size of bubble exceeds the density scale-height, the bubble becomes elongated in the direction perpendicular to the galactic disk (Tomisaka \& Ikeuchi 1986; Tenorio-Tagle, Bodenheimer \& R\'{o}\.{z}yczka, 1987; MacLow \& McCray 1987). It is shown that when the mechanical luminosity released by sequential supernova explosions is high, the expansion of the bubble to the halo is accelerated and the hot gas contained in it flows into the galactic halo (galactic fountain), finally. However, the magnetic fields running parallel to the galactic disk prevents the gas from flowing in the vertical direction (perpendicular to the fields). Tomisaka (1990) studied the adiabatic evolution of a superbubble in uniform magnetic fields and showed that the bubble driven by a mechanical luminosity of $L_{\rm SN}\simeq 3\times 10^{37} {\rm erg\,s^{-1}}$ is confined in the galactic disk by the effect of the magnetic field provided its strength is as large as $B_0=5\mu$G. This shows the superbubble is confined in the galactic disk, if (1) the magnetic fields have a large scale-height, $H_B$, and (2) $L_{\rm SN}\simeq 3\times 10^{37}{\rm erg s^{-1}}$ and $B_0\agt 5\mu$G. However, this result may be affected by the assumption of the adiabatic gas. The interstellar magnetic fields seems to play an important role in the evolution of a superbubble (see also Mineshige, Shibata, \& Shapiro 1993, Ferriere, MacLow, \& Zweibel 1991). In the present paper, a full magnetohydrodynamical calculation has been done including the radiative cooling. Further, the effect of distributions of magnetic field strength is studied. Plan of the present paper is as follows: in section 2 is given an analytical estimate to a threshold mechanical luminosity under which the superbubble is confined in the galactic disk. From this estimation, we choose the values for the parameters $L_{\rm SN}$, $B_0$, and $H_B$. Numerical method and other assumptions are also presented in $\S$2. Section 3 is for the numerical result, in which the effect of the magnetic fields is shown. In section 4, we will discuss the observability of the superbubble in external galaxies. Using the evolution obtained, we will show that a large fraction of the interstellar space is occupied with superbubbles. | \subsection{Blow-out or confinement?} As shown in the preceding section, although in Model B the mechanical luminosity is much larger than the critical luminosity of equation(\ref{eqn:Lcr0}), the vertical expansion is much decelerated by the effect of magnetic fields. This shows that $L_{\rm SN} < L_{\rm crit}$ may be only a sufficient condition for confinement of the superbubble. The expansion of a spherical shell driven by a pressure in the hot cavity $p$ is formulated as follows: \begin{equation} \frac{dMv}{dt}=4\pi\, R^2 (p-p_{\rm out}), \end{equation} \begin{equation} \frac{dM}{dt}=4\pi\, R^2\, v\, \rho_{\rm out} \end{equation} \begin{equation} \frac{dR}{dt}=v, \end{equation} \begin{equation} \frac{dE}{dt}=L_{\rm SN}-4\pi R^2\, p\, v, \end{equation} where, $M$, $R$, $v$, $p_{\rm out}$, $\rho_{\rm out}$ and $L_{\rm SN}$ represent, respectively, the mass of the shell, the radius and velocity of the shell, the interstellar pressure and its density and the energy release rate from an OB association. These equations are, respectively, the equation of motion, the mass conservation, the relation of a size $R$ to a velocity $v$, and the first law of thermal physics. If we assume $p=(2/3)\times E /(4\pi R^3/3)$, these four equations can be solved numerically. Figure 9 shows a resultant expansion law of a spherical superbubble with $L_{\rm SN}=3\times 10^{37}{\rm erg~s^{-1}}$ in a {\em uniform interstellar medium} of $n_0=0.3{\rm cm^{-3}}$. Each curve corresponds to different external pressures as $p_{\rm out}=1.7\times 10^{-12}{\rm erg~cm^{-3}}$ (solid line), $p_{\rm out}=1\times 10^{-12}{\rm erg~cm^{-3}}$ (dotted line), and $p_{\rm out}=0$ (dashed line). Weaver et al.'s (1977) solution, equation (1), agrees with the curve of $p_{\rm out}=0$. This figure shows that the interstellar pressure plays an important role especially in the late phase of the evolution. This is understood as follows: in the late phase of the superbubble the difference between the internal pressure and the outer one is small and this small difference drives the shell further. Therefore, the critical luminosity would be underestimated if we use a solution without taking the outer pressure into account. The shell of a superbubble continues to expand as long as the energy ejection continues, while a supernova remnant stops its expansion after $p_{\rm out}=p$. Thus, exactly speaking, the superbubble is never confined as long as the OB association is alive. However, a slow expansion driven by a small pressure difference as $p-p_{\rm out}$ is considered as a signature of the confinement. In this figure, we also plot an equivalent radius, which is defined using the volume occupied with a hot matter ($V_{\rm hot}$) as \begin{equation} R_{\rm equiv}\equiv \left(\frac{3V_{\rm hot}}{4\pi}\right)^{1/3}, \label{Requiv} \end{equation} for Models A and C. The equivalent radius for Model A in which the bubble is almost confined in the galactic gaseous disk shows a similar expansion law to that obtained by a thin-shell model (a solid curve). In contrast, that of Model C indicates a completely different expansion law such that after $t \agt 10$ Myr the equivalent radius increases rapidly and in $t\simeq 35$Myr $R_{\rm equiv}$ surpasses the model with $p_{\rm out}=0$. These differences seem to come from the distribution of magnetic field strength. In Model A the total pressure (thermal plus magnetic one) is almost constant as the bubble expands, because the magnetic pressure is dominant over the thermal one and magnetic fields are uniform. While, in Model C the total pressure drops according to the density distribution $\rho(z)$. This figure indicates that when the equivalent radius is well fitted by this thin shell model the bubble is nearly confined in the disk even if the gas disk has a finite scale-height. In contrast, if hot gas is ejected from the galactic disk, the equivalent radius shows a more rapid expansion than that derived by this thin-shell model. Values of equivalent radii at the epochs when numerical runs end are shown in table 1. \subsection{Observability} A shear motion in the galactic rotation and rotation itself may play a role in the evolution of a superbubble (Tenorio-Tagle \& Palou\v{s} 1987; Palou\v{s} et al. 1990; Silich 1993). Galactic shear seems to deform the shape and the Coriolis force makes the shell rotate. The characteristic time-scales of the rotation, $\tau_R$, and the shear, $\tau_S$ are estimated respectively as \begin{equation} \tau_{\rm R} \sim 1/\Omega_0 \sim 40{\rm Myr} (\Omega_0/26{\rm km~s^{-1}~kpc^{-1}})^{-1} \end{equation} and \begin{equation} \tau_{\rm S}\sim (l d\Omega/dR)^{-1} \sim 320 {\rm Myr} (l/1{\rm kpc})^{-1}(\Omega_0/ 26{\rm km~s^{-1}~kpc^{-1}})^{-1} (R_0/8.5{\rm kpc}), \end{equation} where $\Omega_0$, $R_0$, and $l$ are the angular speed of galactic rotation, distance from the galactic center, and a typical size of a superbubble, respectively. Since active SN-explosion phase continues for $\sim 50$ Myr for an OB association (McCray \& Kafatos 1987), in the late phase of $\tau_{\rm R}\la t \la 50$Myr, the effect of the Coriolis force seems to appear as a deformation force of the shell. The $\alpha\omega$-dynamo mechanism driven by superbubble was studied recently by Ferri\`{e}re (1992). The galactic rotation has a little effect on the evolution of a superbubble. Thus, if shells or holes observed in external galaxies are elongated after their inclinations are corrected, their direction seems to indicate that of the magnetic fields. There have been listed 141 HI holes in M31 by Brinks \& Bajaja (1986). These holes are observed, more or less, as elliptical. Particularly, the holes found near the major axis of M31 are important to determine the physical shape of the holes. Many of these HI holes have their major axes perpendicular to the galaxy's major axis. Since this is not explained by projection due to the inclination of M31 ($i=77\deg$), these seem to have physically a shape elongated along the azimuthal direction of the galaxy. This is not inconsistent with observations indicating that a global pattern of the magnetic field is ring-like in M31, which is measured by radio linear polarization observations (for a review, see Sofue, Fujimoto \& Wielebinski 1986), in other words, magnetic field lines run in the azimuth direction in M31. \subsection{Porosity} If hot gas contained in superbubbles occupies a large volume of the galactic disk, a picture of the interstellar medium should be changed (McKee \& Ostriker 1977). The fraction of areas covered by superbubbles younger than $\tau_{\rm active}$ is estimated with a quantity called as two-dimensional porosity which is defined as \begin{equation} Q(t < \tau_{\rm active}) \equiv r_{\rm OB}\int_0^{\tau_{\rm active}} S(t)dt, \end{equation} where $r_{\rm OB}$ is the formation rate of OB associations per unit area and $S(t)$ represents the area which covered by a hot cavity on the mid-plane of the disk $z=0$. This is identical with a two-dimensional porosity parameter calculated by Heiles (1990). He estimated galaxy-wide average of two-dimensional porosity $Q_{\rm 2D}\simeq 0.30$. $\tau_{\rm active}$ should be chosen equal to the oldest age of a superbubble which contains a hot gas inside. If we assume $\tau_{\rm active}=20$ Myr and integrate $S(t)$ for Models A and C ($S\equiv \pi x_c y_c$), these two models give respectively $2.08\times 10^6\, {\rm pc^2~Myr}$ and $1.57\times 10^6\,{\rm pc^2~Myr}$. We adopt the estimation of $r_{\rm OB}$ from a galactic type II supernova rate of $r_{\rm II}\sim 0.01{\rm yr^{-1}}$, that is, we assume that all type II SNe occur in OB associations, number of type II SNe in an OB association is constant irrespective of richness of association as $N_{\rm SN}\sim 100$ and OB associations are uniformly distributed in the galactic disk with radius $R_{\rm gal}\simeq 10$ kpc. This gives an estimation of OB association formation rate as \begin{equation} r_{\rm OB} = \frac{r_{\rm II}}{N_{\rm SN}\,\pi\, R_{\rm gal}^2}, \end{equation} \begin{equation} r_{\rm OB} \simeq 3.2 \times 10^{-7}{\rm pc^{-2}\,Myr^{-1}} \left( \frac{r_{\rm II}}{0.01{\rm yr^{-1}}}\right) \left( \frac{N_{\rm SN}}{100} \right)^{-1} \left( \frac{R_{\rm gal}}{10 {\rm kpc}}\right)^{-2}. \end{equation} This indicates the two-dimensional porosity to be equal to $Q(t < 20 {\rm Myr})\simeq 0.5-0.6$. Thus, rather large fraction of the galactic disk, $1-\exp{(-Q)}\sim 40\%-45\%$, is covered by young ($t < 20$ Myr) superbubbles. | 98 | 4 | astro-ph9804029_arXiv.txt |
9804 | astro-ph9804047_arXiv.txt | We consider magnetic field evolution of neutron stars during polar-cap accretion. The size of the polar cap increases as the field decays, and is set by the last open field line before the accretion disk. Below the polar cap we find the temperature to be so high that electron-phonon scattering dominates the conductivity. Outside the polar cap region, the temperature is such the conductivity is dominated by temperature independent impurity scattering which can be a few orders of magnitude larger than the electron-phonon conductivity. The time-scale for field decay is therefore initially given by impurity scattering dominated conductivity. When the field strength has been reduced to $\sim 10^8 ~{\rm gauss}$ the accretion is spherical and the time scale for field decay is given by the smaller electron-phonon scattering conductivity. The field strength is now reduced rapidly compared to before and this could be a reason for there being no pulsars known with field strengths below $10^8~{\rm gauss}$. We also investigate the evolution of multipoles at the neutron star surface. We find that contribution from higher-order multipoles are at most 30 \% to that of the dipole mode. | Since the discovery of pulsars there has been much discussion of observational evidence for decay of magnetic fields in neutron stars, as well as much theoretical work. As the reviews by Lamb (1991), Chanmugam (1992), and Phinney \& Kulkarni (1994) indicate, there is at present no consensus on the question of whether or not magnetic fields in isolated neutron stars can decay significantly. The general view has been that the electrical conductivity of matter in the cores of neutron stars is so high that the characteristic decay time for fields generated by electrical currents in the core is greater than the age of the Universe. Recently Pethick \& Sahrling (1995) showed that even if the conductivity in the core was small the shortest possible decay time is some two orders of magnitude longer than the decay time for configurations where the magnetic field is confined to the crust. An incorporation of general relativistic effects, Sengupta (1997), further reduces the decay rate. Still, millisecond pulsars have typical field strengths in the range $10^8~ -~ 10^{10}~{\rm gauss}$ compared to isolated radio pulsars which have typically field strengths around $10^{12}~{\rm gauss}$ and this indicates that during the spin-up phase the accretion process is reducing the field strength somehow. Millisecond pulsars are generally found in binary systems where the companion star is a white dwarf with a mass less than a solar mass, $M_{\sun}$. This system is called a Low-Mass-Binary Pulsar (LMBP) referring to the mass of the companion star. The progenitor to this system is thought to be the Low-Mass-Xray Binaries (LMXB) where a neutron star is accreting matter from a companion having a mass less than about 2 $M_{\sun}$. The accreting matter is spinning up the neutron star to millisecond periods. For details concerning this process see reviews by Phinney \& Kulkarni (1994) and Bhattacharya \& van den Heuvel (1991) among others. The evolution of the binary system after the neutron star is formed either by a supernovae or tidal capture, is assumed to occur on at least two time scales when the companion star is in radiative equilibrium. At first the companion star is evolving on a nuclear time scale slowly filling its Roche lobe. When it has been filled up the matter overflows and falls onto the companion neutron star on a thermal time scale $\tau_{th}=G M^2/R L = 5\times10^7~ (M_{\sun}/M)^2 ~{\rm yrs}$, see Bhattacharya \& van den Heuvel (1991) for details. If LMXB's are the only progenitors to LMBP's the lifetime of the LMXB, or the accretion phase of the binary system, must be of order $10^7~{\rm yrs}$. However, by using the amount of mass needed to be accreted to spin up the neutron star to a spin period $P_i$ one finds a time scale that ranges from $10^8-10^{10}\times (P_i/2~{\rm ms} )^{-4/3}~{\rm yrs}$. This discrepancy suggests that there might other progenitors to the LMBP's. For details see the review by Phinney and Kulkarni (1994). We will in this paper assume the accretion phase of progenitors to millisecond pulsars to last between roughly $10^7-10^9~{\rm yrs}$. Matter accreting onto a neutron star is expected to be disrupted by the magnetic field at some radius, $r_A$, sometimes called the Alfv\'en radius. Beyond this bare description there is no generally accepted view on how or where the matter attaches to the field lines and flows to the neutron star's surface, or on the interaction of the field with the matter outside $r_A$, despite a large number of papers on the subject, see King (1995). In the case of disk accretion there are two main approaches to the problem. In one (Ghosh \& Lamb 1978, 1979a,b; Kaburaki 1986 and Wang 1987) the field is assumed to thread a large fraction of the disk because of Kelvin-Helmholtz instabilities. The other approach (Aly 1980; Anzer \& B\"orner 1980, 1983; Scharlemann 1978) assumes the disk is a perfect conductor, completely excluding the field. In both approaches the matter is often assumed to leave the disk in i a narrow transition zone at the inner edge (near $r_A$), thereafter flowing along field lines to the neutron star. Most studies concerned with the magnetic field evolution {\it inside} the neutron star have assumed matter to accrete onto the surface spherically, e.g. Fujimoto et al. (1984), Miralda-Escud\'e et al. (1990), hereafter Mir90, Urpin \& Geppert (1995), and Konar \& Bhattacharya (1997). In this paper I consider non-spherical accretion where matter is assumed to flow onto the polar caps in a column. The cap is consequently heated up and the conductivity in the crust below the polar cap we estimate to be much smaller than the conductivity outside the accretion column. The magnetic evolution in this scenario can be divided into two stages where in stage I the global decay rate is controlled by the conductivity outside the accretion column, $\tau_B\sim 10^{8}-10^{10}~{\rm yrs}$. When the field has reached a value of about $10^8 {\rm gauss}$ the accretion is spherical and the evolution enters the second stage. Here, the whole crust is being heated up and the conductivity is dominated by electron-phonon scattering. In this stage the magnetic field decay time is a few orders of magnitude shorter than in stage I and compared to the earlier evolution the field is dissipating rapidly. I argue that this could account for the fact that no binary pulsars have been found with a magnetic field less than $10^8~{\rm gauss}$. A similar effect where the accreted flow is pushing the field lines has already been noticed by Romani (1993). The paper is organized as follows: section 2 contains a brief description of the model and the basic equations. In section 3 we estimate the length scale for temperature change at the accretion column-normal crust boundary to be less than a crust thickness. The time scale for the temperature to reach a stationary state is also shown to be much smaller than the magnetic field decay time scale. Therefore, we do not solve the energy equation explicitly but assume the conductivity as a function of angle to be close to a top-hat function at all times. Section 4 presents the numerical results for some initial depths of the magnetic field. | We have examined the influence of asymmetric accretion on the magnetic field evolution of a neutron star. The temperature structure in the crust resulting from the accretion was roughly estimated and found to vary more rapidly than the crust thickness around the edge of the accretion cap. The time scale for reaching a steady state was shown to be shorter than the time scale for magnetic field decay. Therefore, instead of doing a full blown calculation of heat conduction coupled with the magnetic field evolution we used a simple smoothed top-hat function for the temperature structure and consequently for the conductivity. The global field decay time is roughly the one given by the largest conductivity in the crust, which occurs outside the accretion cap, $\tau_{B,i} = 5.5 \times 10^{17}~(\delta R_5)^2~(\rho_{14} x/0.1)^{1/3}Z/(60~ (Q/0.01))~{\rm s}$. However, when the field strength is down to roughly $10^8$ gauss the accretion is spherically symmetric and we have $\tau_{B,ph} = 5.5 \times 10^{15}~ (\delta R_5)^2~\rho_{14}^{7/6} ~(x/0.1)^{5/3}T_8^{-2}~ {\rm s} \ll \tau_{B,i}$ and so the decay rate is much faster than initially which could provide a clue why no binary pulsars are known with field strengths less than $10^8$ gauss. Romani (1993) discussed a similar effect and found a threshold close to ours. We also found the asymmetric accretion resulted in some higher order surface magnetic multipoles and the size of these were shown to be at most 30 \% of the surface dipole field. {\bf Acknowledgements :} I would like to thank Dong Lai, Lars Bildsten and Edward Brown for useful discussions. Dong Lai is especially thanked for a careful reading of the manuscript. This work was also supported in part by the U. S. National Science Foundation under grants NSF AST93-15133 and AST94-14232, by NASA under grant NAGW-1583, and by the Swedish Natural Science Research Council. \vfill\eject | 98 | 4 | astro-ph9804047_arXiv.txt |
9804 | astro-ph9804271_arXiv.txt | The lack of bright host galaxies in several recently examined gamma--ray burst (GRB) error boxes suggests that the redshifts of cosmological GRBs may be significantly higher than previously hypothesized. On the other hand, the non--detection of multiple images in the BATSE 4B catalog implies an upper limit to the average redshift $\langle z\rangle$ of GRBs. Here, we calculate an upper limit to $\langle z\rangle$, independent of the physical model for GRBs, using a new statistical lensing method that removes distance ambiguities, and thus permits accurate computation of the lensing rate at high $z$. The upper limit on $\langle z\rangle$ depends directly on the cosmological parameters $\Omega$ and $\Lambda$. If there are no multiple images among the brightest 80\% of the first 1802 bursts in the BATSE 4B catalog, then, at the 95\% confidence level, $\langle z\rangle<$ 2.2, 2.8, 4.3, or 5.3 for ($\Omega$, $\Lambda$) values of (0.3, 0.7), (0.5, 0.5), (0.5, 0.0), or (1.0, 0.0), respectively. The 68\% upper limit to the average redshift is comparable to or less than the median redshift of GRBs in scenarios in which the GRB rate is proportional to the rate of star formation, for any cosmology. The uncertainty in the lensing rate---arising from uncertainties in the cosmological parameters and in the number density and average velocity dispersion of galaxies---will be reduced significantly in the next few years by a new generation of experiments and surveys. Moreover, the continued increase in the number of GRBs observed by BATSE will greatly constrain their redshift distribution. | Three decades after their discovery (\cite{KSO73}), the physical origin of gamma--ray bursts (GRBs) remains unresolved. Recent developments, however---including the isotropy of GRBs seen by BATSE (\cite{B96}; \cite{THBM96}) and the detection of redshift $z$=0.835 absorption and emission lines from a possible optical counterpart to GRB 970508 (\cite{M97}; \cite{R98})---strongly suggest a cosmological origin for GRBs. Initially, no--evolution fits to the log N -- log P (peak flux) distribution of BATSE GRBs suggested a typical redshift of $z\sim 1$ for dim bursts (\cite{F93}; \cite{W93}), with the break in the $-3/2$~slope of the flux distribution at $P \sim 10$ ph~cm$^{-2}$~s$^{-1}$ being interpreted as a cosmological deviation from Euclidean space at $z \sim 1$. In addition, a number of researchers reported a factor of $\sim 2$ ``time--stretching'' of dim bursts relative to bright ones (e.g., \cite{Nor94}, 1995), which was thought to imply a redshift of order unity for the dim bursts. Such a redshift would imply an extremely low rate of gravitational lensing and multiple imaging of sources detected with BATSE, perhaps as low as one multiple imaging event per 200 years (\cite{GN94}). New lines of evidence now indicate that the typical redshift of cosmological GRBs may substantially exceed unity. Particularly suggestive evidence for high GRB redshifts comes from deep HST searches of the error boxes of several of the brightest GRBs detected by BATSE and the Interplanetary network (IPN): In five cases, no obvious host galaxies for the bursts were found, down to a limiting magnitude between 3.5 and 5.5 mag fainter than what would be expected if GRBs reside in $L_\star$ galaxies and dim bursts are at $z\sim 1$ (\cite{Sch97}). This ``no--host problem'' has been analyzed further by Band \& Hartmann (1998), and implies that if GRBs reside in normal galaxies, even the {\em brightest} of them are at redshifts close to unity, with the faintest being at much higher redshift still. In addition, Fenimore \& Bloom (1995) showed that the intrinsic anticorrelation between photon energy and burst duration implies that if the observed time--stretching is caused by time dilation, then the redshift of the dimmest bursts may be as large as $z\sim 6$ (note, however, that some or all of the time--stretching may be intrinsic to the bursts; cf. \cite{SPS97}). These new lines of evidence have led researchers to explore scenarios in which the bursts are at much higher redshifts. One currently popular scenario, motivated in part by the assumption that cosmological GRBs involve remnants of massive stars, is that the GRB rate is proportional to the rate of (massive) star formation (\cite{Tot97}; \cite{WBBN98}). The star formation rate is thought to vary strongly with redshift (\cite{Ma96}), peaking at $z \sim 2$. If this is the case, the dimmest bursts may be at redshift $z\sim 6$, making them the most distant objects ever detected (\cite{WBBN98}). Other authors have noted that, if the comoving number density of GRB sources is allowed to vary with redshift, then even if the peak luminosity is fixed, the log N--log P distribution and other properties of the BATSE bursts can be fit by models with a maximum redshift up to $z\sim$10--200 (\cite{RHL95}). In these high-$z$ scenarios, the expected incidence of gravitational lensing and multiple imaging of GRB sources is much higher than it is in lower-$z$ scenarios, since in general the lensing rate increases markedly with the source redshift. The short duration of GRBs (typically tens of seconds) compared to the difference in light travel time between different ray paths (typically months, if the lens is of galactic mass; cf. \cite{M92}) means that a multiply-imaged GRB appears as two or more separate events with identical time histories and intensities that differ only by a scale factor. The image separations, of order arcseconds, are tiny compared to BATSE location errors; thus, these GRBs would appear to come from the same location. Two or more such events must be detected in order to identify a lensed GRB source; thus, the overall BATSE burst detection efficiency, $\epsilon$, is of prime importance in determining the observed incidence of lensing. The average efficiency for the 4B catalog is $\epsilon = 0.48$ (\cite{M98}), which is 40\% larger than the previously estimated value (\cite{F94}). This increase implies that the expected incidence of lensing is significantly higher than was previously believed. Lensing of GRBs was first suggested by Paczynski (1986) as a way to establish their cosmological origin. Subsequently, many authors have calculated the lensing rate of GRBs (\cite{M92}; \cite{BW92}; \cite{N93}; \cite{GN94}), assuming GRB redshifts estimated from no--evolution fits to the log N -- log P distribution, viz., $z\lta 1$. For redshifts $z \lta 1$, the gravitational lensing rate is reasonably well known for a given cosmology (e.g., \cite{TOG84}; \cite{FT91}; \cite{F92}). At redshifts $z \gta 1$, however, the lensing rate has been uncertain, mainly because of ambiguity in the angular diameter distance at high redshift (\cite{F92}). As a result, even when the cosmology and the number density and properties of lenses are fixed, the estimated lensing rate can vary by factors of several between the different prescriptions. Motivated in part by this ambiguity, Holz \& Wald (1998) have developed a numerical method to calculate lensing rates for a given cosmology, combining techniques from both ``ray-shooting'' and ``Swiss-cheese'' models. The approach resolves the angular-diameter distance uncertainties, and also correctly accounts for multiple lens encounters along the line of sight, allowing for an unambiguous calculation of lensing rates. If there are no multiple images among the bursts detected with BATSE, then an upper limit to the average redshift $\langle z\rangle$ of GRBs can be inferred, one which depends directly on the cosmological parameters (see also Nemiroff et al.\ 1994; Marani et al.\ 1998; Marani 1998). Here, we set an upper limit to $\langle z\rangle$, independent of the physical model for GRBs, using the Holz \&~Wald (1998) method to compute the lensing rate for a variety of cosmologies. The plan of this paper is as follows. In \S~2 we describe and develop the statistical lensing method, and compare its results with analytical estimates. We show that analytical estimates using the ``filled-beam" approach are reasonably accurate for source redshifts less than $\sim$2--3, but underestimate the true rate of lensing by tens of percent at higher redshifts. In \S~3 we use the numerical method to compute the lensing rate as a function of redshift for several cosmologies, and compare these results with previous estimates. We also determine upper limits on the rate of lensing and on the average redshift $\langle z\rangle$ of GRBs, assuming that no lensing events are present in the BATSE 4B catalog (\cite{M98}). We discuss the implications of these results and give our conclusions in \S~4. We follow the conventions that the Hubble constant is 100\,$h$\ km~s$^{-1}$~Mpc$^{-1}$, $\Omega$ is the present mean density in the universe in units of the closure density, and $\Lambda$ is the present normalized cosmological constant. In a flat universe, $\Omega+\Lambda=1$. | We have exhibited a new method for calculating gravitational lensing rates, given a cosmology and a distribution and evolution of the lenses. This method, which is described more fully in Holz \& Wald (1998), is free of the angular diameter distance ambiguity that is a feature of standard analytical methods. For the cosmologies and redshift ranges we have examined, we find that the lensing rate is approximately equal to the filled-beam lensing rate for $z<3$, but in excess of the filled-beam rate for $z>3$. We find that the lack of detected lensing in the current BATSE catalog places an upper limit to the median redshift of bursts that, at the 68\% confidence level, is comparable to or less than the median redshift in star-formation burst models. We now discuss these results and place them in context. In \S~4.1 we discuss the simplifying assumptions that have been made in our calculation. We find that these assumptions have, if anything, caused us to {\it underestimate} the lensing rate. In \S~4.2 we examine the uncertainties in the inputs used in our calculations, as well as the future prospects for reducing these uncertainties. In \S~4.3 we discuss the implications of the current lack of lensing for cosmological models of gamma-ray bursts, when combined with existing lower bounds to the redshift of bursts in these models. Finally, in \S~4.4 we provide a future outlook for the rapidly emerging importance of constraints inferred from detection or nondetection of lensing. \subsection{Effects of Simplifying Assumptions} We have assumed that the angular distribution of the sources of gamma-ray bursts is uncorrelated with the angular distribution of the lenses. This is well justified, because for sources at low redshift the lenses that contribute most to the lensing rate are at about half the redshift to the source, and for sources at high redshift the dominant lenses are at a redshift of approximately unity (see \S~2; see also, e.g., Turner et al.\ 1984; \cite{M92}). Hence, in a cosmological model the sources of gamma-ray bursts are separated from lenses by hundreds of megaparsecs to gigaparsecs, and therefore it is extremely unlikely that the sources and lenses are angularly correlated. We have also assumed that the phase of the orbit of BATSE is uncorrelated between the separate images of the burst, and hence that the joint probability of two images being observable is just the square of the BATSE efficiency, $0.48^2 \approx 0.23$. This assumption is reasonable if the time delay between images is significantly larger than the BATSE orbit time of $\sim 5000$ seconds. The time delay for a mass $M$ is of order $10^{-5}(M/M_\odot)$ seconds (see, e.g., Blandford \& Narayan 1992), and hence the assumption of uncorrelated phases is good for masses $M\gta 10^9\,M_\odot$. This includes almost all the effective lensing mass of galaxies, and thus this assumption is also unlikely to affect the calculated lensing rates significantly. We may have underestimated the lensing rate by assuming that only galaxies---which we have modeled as singular isothermal spheres---contribute to the lensing rate. Nemiroff et al.\ (1993) have noted that point masses, for example supermassive black holes, could also in principle contribute. To produce gamma-ray burst lensing of the type that we consider here, in which the lensing event appears as two separate bursts, the mass of such a lensing black hole must be $M\gta 10^8\,M_\odot$, because smaller masses would produce time delays less than $\sim$1000 seconds, so that overwriting or readout time would tend to prevent BATSE from detecting two separate bursts (the effect of lensing by lower-mass black holes may, however, be detectable using techniques such as autocorrelation analysis; see Nemiroff et al.\ 1994, 1998). If, however, the total lensing rate were dominated by very massive black holes, $M\gta 10^{10}\,M_\odot$, then these objects would also produce detectable image separation of lensed quasars or galaxies (the angular separation is of order $3(M/M_\odot)^{1/2} (D/1 {\rm Gpc})^{-1/2}\,\mu$arcsec for an Einstein ring, which is 0.3" for $M=10^{10}\,M_\odot$ and $D$=1 Gpc). Hence, the amount by which we have underestimated the lensing rate depends on the number of point masses in the relatively narrow range $10^8$--$10^{10}\,M_\odot$, which is the only mass range that could avoid detection in quasar lensing surveys and yet produce separately detected bursts. We may also have underestimated the lensing rate by only keeping track of the total number of images, not whether there are, e.g., two or four images in a particular event. As pointed out by Grossman \& Nowak (1994), this tends to underestimate the {\it detectable} rate of lensing because when there are several images it is more probable that at least two of them are observed by BATSE: assuming a 48\% efficiency for any particular image, the probability of observing at least one pair rises from 23\% when there are two images to 66\% when there are four images. Grossman \& Nowak (1994) estimate that, for galaxy ellipticities typical of galaxy surveys, the maximum possible enhancement of the lensing detection rate is $\sim$30\%, but that the overall enhancement is likely to be smaller. A third effect that may enhance the rate of lensing is magnification bias, which is an effect first discussed extensively by Fukugita \& Turner (1991) in the context of quasar lensing. Magnification bias occurs because sources that would have been undetectable are made visible by lensing, and hence a magnitude-limited sample contains an enhanced incidence of lensing. This can be especially important if the number of sources at a given flux rises steeply at the faint end, as it does in quasars. To detect multiple images and thus confirm strong lensing, it is necessary to detect the fainter image as well as the stronger image. For persistent sources such as quasars, this can be done by following up broad, shallow surveys by deep pointings, so that essentially all of the secondary images are detectable (e.g., as in the Hubble Snapshot Survey [Bahcall et al. 1992; Maoz \etal\ 1992, 1993a,b]). For transient sources such as gamma-ray bursts, deep follow-ups are not always possible, and hence for the secondary image to be detectable it must have a brightness in excess of the threshold of the original survey (i.e., in the case of gamma-ray bursts, the BATSE sample). This effect, plus the flatness of the faint end of the gamma-ray burst log N -- log P curve, suggests that the magnification bias for gamma-ray bursts is probably less than the magnification bias for quasars (see also \cite{GN94} for a discussion of this point). Nonetheless, magnification bias for gamma-ray bursts could in principle increase significantly the expected lensing rate, and hence decrease the upper limits on $\langle z\rangle$, compared to the conservative estimates here. We may underestimate the lensing rate because we have neglected the clustering of galaxies, the last of our assumptions. The enhanced gravitational potential in clusters tends to increase the convergence of null geodesics and hence increase the cross section for strong lensing (see also Holz \& Wald 1998). Numerical estimates suggest that the lensing rate for galaxies in a cluster will be increased by $\sim$10--20\% by this effect, but because only $\sim$10\% of galaxies are in clusters this effect is unlikely to increase the overall lensing rate significantly. Finally, however, we may have overestimated the rate of lensing by ignoring evolution in the properties of the lensing galaxies. If the comoving density in lensing galaxies was less in the past than it is now, or if these galaxies were less massive than they are today, high-redshift lenses would contribute less to the lensing rate than we have assumed. Note, however, that the peak in the lens redshift distribution is at about unity (Fig.~2), when galaxies had properties very similar to their current properties. Note also that observations of lensing of quasars (see below), which have a typical redshift similar to the proposed typical redshift $z\sim 2$ of GRBs, suggest that if there are evolutionary effects then these effects have not had a dominant effect on the lensing rate. To summarize, the net effect of our simplifying assumptions is likely to be that we underestimated the lensing rate by $\lta$10\%, and hence our results our conservative in that they give a slightly high upper bound to $\langle z\rangle$. \subsection{Uncertainties in Input Parameters} In addition to the simplifying assumptions discussed above, our calculations are clearly dependent on the input values we have assumed. One input is the dimensionless parameter $F$, which is proportional to the product of the number density of galaxies and the fourth power of their average velocity dispersion. The lensing rate scales linearly with $F$. We have used $F$=0.1, which is consistent with the recent estimates from the Century Survey (Geller et al.\ 1997). Other estimates range from $F$=0.05 to $F$=0.15, and if these other estimates are used then the range of possible maximum redshifts is increased for a given cosmology. Fundamental cosmological parameters, another input to our calculations, are also uncertain. A standard Einstein-De Sitter universe ($\Omega$=1, $\Lambda$=0) gives the lowest lensing rate, and a $\Lambda$-dominated universe gives the highest. Current observation of quasar lensing puts constraints on the allowed combinations of $F$ and cosmology. For example, five of the 351 quasars in the HST Snapshot Survey (Maoz et al.\ 1993b) were lensed. A comparison of this rate with the rate expected for models is complicated by effects such as magnification bias (see above). However, the high average redshift of the quasars ($z\sim 2$) means that many of the issues affecting lensing of gamma-ray bursts, such as evolution of clustering, also affect the lensing of quasars. Hence, quasar lensing statistics have bearing on estimated GRB lensing rates, and will become even more relevant as more extensive surveys are done. In addition, all of these uncertainties will be diminished greatly by the data that emerge from the plethora of satellites and surveys planned for the next few years. The number density and velocity distribution of galaxies (and hence $F$) will be determined with unprecedented accuracy by surveys such as 2dF (Colless 1998) and the Sloan Digital Sky Survey (Margon 1998), the latter of which is expected to see first light in 1998. Data from these surveys will also provide valuable information about $\Omega$, $\Lambda$, clustering, and the lensing rate of quasars. Complementary information about cosmological parameters will be extracted from high-redshift supernova surveys (Perlmutter \etal\ 1998) and from data gathered by the many upcoming microwave background experiments, such as MAP (Wang, Spergel, \& Strauss 1998), TOPHAT (Martin et al.\ 1996), and PLANCK (see, e.g., Bond, Efstathiou, \& Tegmark 1997 for a discussion of the constraints). The evolution of galaxy clustering is already being inferred from cluster surveys, and AXAF observations of X-ray emission from clusters of galaxies is expected to greatly enhance this understanding. Hence, one side effect of the coming data-rich era in cosmology is that calculations of lensing rates will be much less uncertain, and therefore upper limits to the redshift of a population such as gamma-ray bursts will be far more secure. \subsection{Current Implications for Cosmological GRB Models} The redshift limits presented in this paper have important consequences for cosmological models of GRBs. Lower limits on the redshift of bright GRBs are becoming stronger as the number of detected bursts rises, both because of the no-host problem (Schaefer et al.\ 1997) and because GRBs do not show any evidence of large-scale clustering (Lamb \& Quashnock 1993; Quashnock 1996). These limits suggest that if the sources of GRBs are in or near galaxies, the median redshift of the bursts detected with BATSE may be significantly greater than unity. However, the upper limits to the redshift of GRBs that follow from the lack of lensing are beginning to make such high-redshift populations less appealing. For example, the lack of lensing in the current BATSE catalog strongly rules out the \hbox{$z\sim 10$--200} models that were heretofore consistent with the properties of the BATSE bursts (Rutledge et al.\ 1995). Also, models in which the burst rate is proportional to the star formation rate (Totani 1997; Wijers et al.\ 1998), and for which $\langle z\rangle\sim 2.2$, are beginning to be constrained strongly by the lack of lensing: Our 68\% upper limit on $\langle z\rangle$ (eq.~[6]) is less than 1.9, and probably close to 1.0. If the redshift upper limits are reduced further, it may be necessary to postulate either a very short interval in which GRBs were produced, with an intrinsic brightness distribution that matches the observed log N -- log S distribution, or a population that is unassociated with any known population of objects, so that the lower limits to redshift do not apply. \subsection{Future Outlook and Summary} In the next few years gravitational lensing is likely to play a role of rapidly increasing importance in constraining cosmological models of gamma-ray bursts. As discussed in \S~4.2, the uncertainties in the calculation of the expected lensing rate will be diminished greatly, and hence the limits will be robust. In addition, as BATSE continues to detect bursts the statistics will steadily improve, and in the current high-redshift models lensing is to be expected in the next few years. In five years, the sample will be roughly double the $\sim$1800 bursts considered here, and therefore if lensing continues to be absent, upper limits on the redshift will become extremely restrictive. At the same time, the improved statistics will continue to increase the lower limits on redshift, if there is no apparent large-scale clustering (Lamb \& Quashnock 1993; Quashnock 1996). Moreover, the accurate positional estimates of BeppoSAX, HETE II, and the Fourth Interplanetary Network (IPN) are likely to provide $\sim$100 error boxes of area a few square arcminutes or less in the next five years. If instead lensing {\it is} detected, there will be many important consequences. For one, the cosmological origin of a particular GRB will be established from gamma-ray data for the first time (as opposed to from the afterglow, as may be the case for GRB~970508 [\cite{M97}]). The approximate mass of the lens, and hence its luminosity, can be estimated from the time delay between images. The sharp peak in the probable redshift of lenses (see \S~2) would then give a reasonably accurate estimate of the flux of the lens. If the GRB can be reasonably well-localized (e.g., with the IPN), this will allow a search for the lens. If the lens is found, the GRB source must essentially be angularly coincident with it, because the deflection angle is only about 1", and thus counterpart searches will be facilitated greatly. Knowledge of the redshift of the lens, combined with $H_0$ and an independent estimate of the mass of the lens (e.g., from its luminosity or velocity dispersion) will allow use of the time delay to estimate the redshift of the source. The quantifiable diversity of gamma-ray burst light curves (\cite{MN98}; \cite{Mar98}) means that false positives are unlikely, and hence even a single lensing event will generate a wealth of data. In conclusion, the increased estimates of the redshift of gamma-ray bursts that were forced by the no-host problem, in addition to the greatly improved statistics of bursts and the increase in the estimate of the BATSE burst detection efficiency from 34\% to 48\%, mean that gravitational lensing limits are now playing a prominent role in constraining gamma-ray burst models. The role of lensing will become dramatically more important and robust in the next five years. | 98 | 4 | astro-ph9804271_arXiv.txt |
9804 | astro-ph9804101_arXiv.txt | The blazar Mkn 421 has been observed, as part of the AO1 Core Program, five times from 2 to 7 May 1997. In the LECS+MECS energy band the spectrum shows convex curvature, well represented by a broken power--law. Flux variability (more than a factor 2) has been detected over the entire 0.1--10 keV range, accompanying which the spectrum steepens with the decrease in intensity. Mkn 421 has also been detected with the PDS instrument. Our preliminary analysis indicates that the PDS spectrum lies significantly above the extrapolation from the MECS, suggesting a contribution from a flatter high energy component. | Mkn 421 is one of the best known and studied BL Lac objects. It shows optical polarization, flat radio spectrum, and large variability, characteristics of the blazar class. It is bright in X--rays, where it shows prominent flares accompanied by significant spectral changes \cite{takahashi}. Among GeV--emitting blazars, Mkn 421 is unique in being the first object in which the $\gamma$--ray emission is detected to extend up to TeV energies (\cite{punch,krennrich}), at a level allowing detailed spectral and variability studies in the broadest available range of frequencies. Its GeV (EGRET) $\gamma$--ray emission connects smoothly with the E$>$0.5 TeV spectrum (e.g.~\cite{macomb}). On the contrary the X--ray spectrum, up to 10 keV, did not show any hint of the onset of the inverse Compton component responsible for GeV and TeV emission. It was then of great interest to take advantage of the full capabilities of {\it Beppo}SAX, enabling a spectral coverage up to $\gta$ 200 keV, to look for it. At the same time the fact that in the X--ray band we are observing the emission from the highest energy tail of the emitting--particle distribution could provide important clues on particle acceleration and cooling mechanisms. The Mkn~421 observations discussed here are part of the AO1 Core Program dedicated to bright blazars. The data reduction presented here has been generally performed with software released {\it before} September 1997. Data reduction with the updated software for all the on board instruments, and a more detailed analysis, will be presented in forthcoming paper with all appropriate references. \begin{figure}[t] \vspace{9pt} \psfig{file=lecs_2mecs_lc.ps,width=7.5truecm,rheight=5.3truecm} \caption{\small\sf Rebinned light curve (4000~s bins) of {\it Beppo}SAX data, LECS, and MECS divided in two energy bands [1.5 -- 4], [4 -- 10] keV.} \label{fig:light_curve} \end{figure} \begin{table*}[t] \setlength{\tabcolsep}{0.7pc} \newlength{\digitwidth} \settowidth{\digitwidth}{\rm 0} \catcode`?=\active \def?{\kern\digitwidth} \caption{{\it Beppo}SAX Observations Log} \label{tab:log} \begin{tabular*}{\textwidth}{@{}l@{\extracolsep{\fill}}ccccc} \hline {Pointing} & {LECS } & {LECS } & {MECS } & {MECS } & {MECS } \\ {Start Date} & {exp. time} & {[0.1--4 keV]} & {exp. time} & {[1.5--10 keV]} & {[$<$4/$>$4 keV]} \\ { } & {(ksec)} & {(cts/s)} & {(ksec)} & {(cts/s)} & {(cts/s)} \\ \hline { 2/V/1997 @ 04:10 } & { 4.4 } & {$ 2.806 \pm 0.020 $} & { 11.4 } & {$ 2.862 \pm 0.016 $} & { (2.14/0.72) } \\ { 3/V/1997 @ 03:24 } & { 4.3 } & {$ 1.748 \pm 0.020 $} & { 11.7 } & {$ 1.694 \pm 0.012 $} & { (1.28/0.41) } \\ { 4/V/1997 @ 03:25 } & { 4.9 } & {$ 1.362 \pm 0.017 $} & { 12.2 } & {$ 1.027 \pm 0.009 $} & { (0.81/0.22) } \\ { 5/V/1997 @ 03:32 } & { 4.9 } & {$ 1.823 \pm 0.029 $} & { 11.9 } & {$ 1.522 \pm 0.012 $} & { (1.52/0.35) } \\ { 7/V/1997 @ 04:47 } & { 6.0 } & {$ 1.612 \pm 0.016 $} & { \nodata } & { \nodata } & { \nodata } \\ \hline \end{tabular*} \end{table*} | {\it Beppo}SAX dat aof Mkn 421 shows interesting variability both in flux and in spectral shape, with a marked softening corresponding to decreasing brightness. This kind of spectral variability behavior is well known in the X--ray band for sources of the class of Mkn~421, the so called High-Frequency-Peaked BL Lacs (HBL). In general in blazars at energies just above the synchrotron peak it the relationship {\it harder-when-brighter} holds and is generally interpreted in terms of injection of fresh electrons in the highest energy end of a single population. On May 2$^{\rm nd}$ the peak of the synchrotron component could possibly fall in/or just below the LECS band, while in the lower state of May 4$^{\rm th}$ this is not longer true. Moreover, preliminary analysis of PDS data suggests the presence of a deviation from the continuously downward curvature for E $> 10$ keV, possibly being the signature of the onset of a different harder spectral component. Further temporal and spectral analysis, also with comparison of multifrequency data and comparison with model predictions is in progress. | 98 | 4 | astro-ph9804101_arXiv.txt |
9804 | astro-ph9804337_arXiv.txt | A statistical study of global galaxy parameters can help to improve our understanding of galaxy formation processes. In this paper we present the analysis of global galaxy parameters based on optical and near-infrared observations of a large sample of edge-on disc galaxies. \\ We found a correlation between the ratio of the radial to vertical scale parameter and galaxy type: galaxies become systematically thinner when going from S0's to Sc's, whereas the distribution seems to level off for later types. \\ The observed scale length ratios (and thus the radial colour gradients) largely represent the galaxies' dust content. On average the colour gradients indicated by the scale length ratios increase from type Sa to at least type Sc. For galaxy types later than Sc, the average colour gradient seems to decrease again. \\ The distribution of {\it K}-band (edge-on) disc central surface brightnesses is rather flat, although with a large scatter. However, the latest-type sample galaxies ($T > 6$) show an indication that their average disc central surface brightnesses may be fainter than those of the earlier types. This effect is probably not the result of dust extinction. | A study of the statistical properties of highly inclined, or ``edge-on'' galaxies benefits greatly from the special orientation with respect to the line of sight of such galaxies. Observations of edge-on galaxies provide us with direct measurements of the luminosity and colour distributions both perpendicular to the galaxy planes and along the galaxies' major axes at various heights above the plane. Indirectly, these luminosity distributions can be related to the galaxies' density distributions and thus their global structure. Moreover, in-depth knowledge of the dust distribution, and hence the optical depth of galaxies, is important for our understanding of galaxy evolution. \subsection{The flattening of exponential discs} \label{ratio.sec} A major advantage of studying highly inclined galaxies is that one can determine their radial and vertical scale parameters directly and independently, since the dependence of these parameters on inclination is smallest for the highest inclinations (e.g., van der Kruit \& Searle 1981a; de Grijs et al. 1997). These scale parameters provide us with information about the intrinsic shape of galaxy discs, i.e., their flattening, in a more direct way than the canonical axis ratios. Moreover, since the vertical scale height, $z_0 = 2 h_z$ (where $h_z$ is the exponential scale height), is to first order independent of radius (e.g., van der Kruit \& Searle 1981a,b, 1982a; Kylafis \& Bahcall 1987; Shaw \& Gilmore 1990; Barnaby \& Thronson, Jr. 1992; but see de Grijs \& Peletier 1997), the radial to vertical scale parameter ratio, $h_R/z_0$, can often be determined more accurately than the major to minor axis ratio. By studying the scale parameter ratio statistically, we may be able to put constraints on the disc formation processes as well as on the stability of galaxy discs (e.g., Bottema 1993). When considering the physical processes that determine the scale parameters one does not immediately expect a strong correlation between scale length and scale height. The scale height is likely determined by the internal, secular evolution of the stellar velocity dispersion (e.g., van der Kruit \& Searle 1981a; Carlberg 1987), whereas the scale length is basically the result of the composition of the protogalaxy (Fall 1983; van der Kruit 1987). However, one might expect that in a larger disc, with a greater rotation velocity, the heating of the disc stars may be more violent, thus resulting in a larger scale height. Therefore, one may expect a correlation between the rotation velocity (which can be related directly to the scale length) of a galaxy disc and the scale height, although the precise dependence is yet unknown (see, e.g., Bottema 1993). Thus, statistics on the ratio of scale length to scale height can be expected to give information on the importance of the formation processes in disc galaxies with different properties. Moreover, once the $h_R/z_0$ ratio is known, one may be able to determine the (theoretical) maximum rotation of a disc from measurements of the vertical disc dispersion (Bottema 1993). Therefore, a statistical treatment of the scale parameter ratio may put general constraints on both the kinematical properties and the global stability of galaxy discs. Bottema (1993) predicts that a constant value for the $h_R/z_0$ ratio leads to a more or less constant mass-to-light ratio of the old disc, $(M/L)_B$, under the assumption that we are dealing with exponential, locally isothermal discs with a constant ratio of vertical to radial velocity dispersion. On the other hand, if we assume a linear relationship between the old-disc absolute luminosity and the vertical velocity dispersion, Bottema (1993) shows that, for a constant $(M/L)_B$, the $h_R/z_0$ ratio decreases rapidly from faint galaxies to a constant level for normal and bright galaxies. Thus, in general, the observed velocity dispersions imply that a constant old-disc mass-to-light ratio results in an approximately constant scale parameter ratio, whereas a constant scale parameter ratio also leads to a mass-to-light ratio that is, to first order, constant. In fact, these predictions imply that all galaxy discs are governed by equal mass-to-light ratios in the old stellar populations, assuming that all galaxy discs have approximately the same colour (Bottema 1993). However, the assumption of a constant and equal mass-to-light ratio of the old-disc population in disc galaxies is probably not physically realistic, considering the range of colours observed within and among galaxies (e.g., de Jong 1996c, and references therein). Therefore, the predicted relationships should be treated with caution and only be used as general guidelines. \subsection{Colour gradients as diagnostics} Broad-band colours are relatively easy to obtain and are therefore the most widely used colour diagnostics to date. They immediately reveal the approximate nature of a galaxy, which is to first order determined by its dominant stellar population and dust content. Although for the detailed analysis of galaxy luminosity and colour profiles one needs to adopt {\it a priori} assumptions concerning the evolutionary stellar population synthesis, the initial mass function, the metallicity and the star formation history, as well as about the dust geometry and its characteristics, de Jong (1996c) shows that the colours formed from different broad-band combinations correlate strongly, which indicates that these colours are probably caused by the same physical process. Therefore, broad-band colours can be used as indicators of changes in the gross properties of galaxies (e.g., changes in metallicity and/or dust contamination). All systematic colour differences induced by stellar population changes and metallicity gradients are generally considerably smaller than the reddening due to dust, however. \subsubsection{Radial colour gradients in edge-on disc galaxies} \label{edgeongrad.sect} In contrast to the large number of studies of radial colour gradients in moderately inclined and face-on spiral galaxies (e.g., de Jong 1996c, and references therein), the colour behaviour of highly inclined and edge-on galaxies has not received much attention. In highly inclined galaxies, the study and interpretation of intrinsic colour gradients is severely hampered by the presence of dust in the galaxy planes, which causes the dust lane to appear as a red feature in vertical colour profiles (e.g., Hamabe et al. 1979; Hegyi \& Gerber 1979; van der Kruit \& Searle 1981b; Jensen \& Thuan 1982; de Grijs et al. 1997). In individual edge-on galaxies, it is generally found that the colours along the major axes, i.e., the locations of the dust lanes, remain nearly constant (e.g., Sasaki 1987; Wainscoat et al. 1990; Aoki et al. 1991; Peletier \& Balcells 1997), although in most cases the outermost disc regions tend to be slightly bluer on the major axis (e.g., Sasaki 1987), which may be explained in terms of an increasingly metal-poor population or a decreased amount of dust at larger galactocentric distances. Generally, as the height above the dust lane and its embedded young disc increases, the radial colour gradients become small or statistically insignificant (e.g., Hamabe et al. 1979, 1980; van der Kruit \& Searle 1982a,b; Jensen \& Thuan 1982; Peletier \& Balcells 1997). \subsubsection{Colour gradients from scale length ratios} \label{colgrads.sect} Since the dust influence varies as a function of passband, scale length ratios could be used as a diagnostic to estimate colour gradients and the dust content of a given galaxy. Evans (1994) studied the effects of dust on the stellar scale length as a function of wavelength, under the assumption that the resulting scale length differences are solely due to dust absorption. His models predict that these differences are small, at least for face-on galaxies, on the order of the observational uncertainties, and even smaller for galaxies with a prominent bulge component. According to his models, if the scale height ratio between dust and stars is $\sim 0.5$ (Peletier \& Willner 1992; Evans 1994), Evans' (1994) models exclude face-on galaxies with $h_B/h_H \approx 2$. On the other hand, larger ratios can be obtained if a galaxy is inclined with respect to the line of sight. The measurement of blue to red scale length ratios alone will not unambiguously reveal the dust content of a given galaxy, because any deviation from unity can equally well be explained by an intrinsic colour gradient, especially for face-on galaxies (Byun et al. 1994). | 98 | 4 | astro-ph9804337_arXiv.txt |
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9804 | astro-ph9804011_arXiv.txt | In this paper we demonstrate for the first time the connection between the spatial and temporal progression of star formation and the changing locations of the very dense regions in the gas of a massive disk galaxy (NGC~1144) in the aftermath of its collision with a massive elliptical (NGC~1143). These two galaxies form the combined object Arp 118, a collisional ring galaxy system. The results of 3D, time-dependent, numerical simulations of the behavior of the gas, stars, and dark matter of a disk galaxy and the stars and dark matter in an elliptical during a collision are compared with multiwavelength observations of Arp 118. The collision that took place approximately 22 Myr ago generated a strong, non-linear density wave in the stars and gas in the disk of NGC~1144, causing the gas to became clumped on a large scale. This wave produced a series of superstarclusters along arcs and rings that emanate from the central point of impact in the disk. The locations of these star forming regions match those of the regions of increased gas density predicted the time sequence of models. The models also predict the large velocity gradients observed across the disk of NGC 1144. These are due to the rapid radial outflow of gas coupled to large azimuthal velocities in the expanding ring, caused by the impact of the massive intruder. | Optically-identified, collision-produced 'ring galaxies' usually display vigorous star formation in an expanding ring or arc of high density gas (e.g., Joy \& Harvey, 1987; Joy \& Ghigo, 1988; Appleton \& Marston, 1997; see Appleton \& Struck-Marcell, 1996). This recent star formation often dominates the appearance of these galaxies, with H$\alpha$ emission tracing the giant HII regions (e.g., Hippelein, 1989a) and outshining the nucleus unless the latter is active. Near-infrared images of these systems show that the older stellar disk population has also been swept into an outwardly expanding, strong density wave. The density disturbance produced by a collision through (and roughly perpendicular to) a galaxy's disk propagates outward through a rotating disk of gas and stars, which is itself at first contracting and then expanding as the collision proceeds. This superposition of material motions and wave propagation produces a pattern of both closed loops and open-ended arcs of relatively high density gas. Shocks can occur in these regions because of the high relative velocities that are produced in the flows. The higher density features are well delineated in observations of these systems and in the combined N-body/hydrodynamic models that we and others have produced. The models can be exploited to further our understanding of the star formation that is triggered by a galaxy collision, a process that has likely occurred in many galaxies over their lifetimes and may have been very frequent earlier in the universe (see Lavery {\it et al.}, 1996). Here we compare our numerical models of these types of galaxy systems with observational data on Arp 118, one particular IR-luminous, gas-rich example of a collisionally produced ring galaxy. This system consists of a strongly distorted disk galaxy and an elliptical in close proximity. Hippelein (1989a) had difficulty in explaining the extreme velocity gradient and complex morphology in the Arp 118 system using a his simple picture of a collision between a gas-rich spiral and an elliptical. However, Gao {\it et al.} (1997) have observed that Arp 118 contains a large amount of molecular gas distributed exclusively along the ring formations (the first high resolution CO observations ever made of a ring galaxy) , and that the velocity structure in this gas is kinematically consistent with the simple collisional model. In this paper, we show that, using fully dynamical, 3D models, we can reproduce the morphology of the disturbed disk galaxy in the pair and the approximate relative positions of the two galaxies. The 'best-fit' model for the Arp 118 system was chosen from a grid of simulations produced by Gerber, Lamb, \& Balsara (1996). These simulations explore the results of face-on collisions (collisions parallel to the spin-axis of the disk) between an elliptical galaxy and a disk galaxy. This dynamical treatment confirms the correspondence between the models and the observed velocity structure. We use the chosen simulation to constrain the timescales for star formation in Arp 118 and to explore the history of the collision by comparing the predicted results of such a collision with current observations. | The simulation presented here is comprised of a sequence of models that follows the evolution of a pair of galaxies through and after a collision. One of the later models closely matches the observed morphology of the CO component in Arp 118 and its velocity field. Earlier models in the sequence provide a good match with the present radio continuum and H$\alpha$ emission, indicating that the sequence of post-collision star formation can be traced and timed by comparison to simulations of encounters between two massive galaxies with masses similar to those observed. Stars, particularly massive stars, are formed in the cores of giant molecular clouds in the highest density regions. Both multiwavelength observations and our numerical models of Arp 118 indicate that strong shocks in the gas together with large increases in the gas volume density are associated with star formation over volumes of 1 kpc$^3$. The observed morphology of the regions of dense gas and the clustering observed in the stars emitting H$\alpha$, which were formed in the gas since the collision, suggest that the observed clumping of the young stars results from a clumping of the densest gas on the same scale. The recent work by Marston \& Appleton (1995) and Appleton \& Marston (1997) also provides evidence that the clumping observed in the optical images of collisionally produced ring galaxies is not due to patchy dust obscuration, because the same clustering is also observed in the near infrared. Gas clumping on this same scale is found in the numerical simulations, suggesting that there is a global explanation for the observed morphology of the dense gas and the resulting giant stellar formations in these systems. The simulations show a relatively small perturbation (clumping) in the density of stars. We therefore predict that the distribution of the old stars in these systems will be observed to be smoother than that of the gas, although the models show that the stars are driven into a wide ring, and sometimes even a second, inner ring, by the collision (see Lamb {\it et al.} 1998). This study of the Arp 118 system demonstrates that a careful comparison between high resolution observations and detailed models can yield insight into the sequence of star formation that takes place in a gas-rich galaxy after a major collision. The intensity and location of the starburst at any particular epoch will depend upon the speed with which density waves are propagating through the expanding disk. Such quantities can now be predicted quite accurately from current models of colliding galaxies. Thus global star formation (on the scale of several hundred parsecs) as it occurs in these systems at the current epoch can be investigated more thoroughly than previously. The rate of galaxy collisions in the past was larger than it is today, due to the greater overall density, thus a considerable portion of the star formation that took place in young disk galaxies at earlier epochs was likely triggered by galaxy collisions. We expect, therefore, that studies like the one reported here will help in understanding this earlier star formation and its current consequences. | 98 | 4 | astro-ph9804011_arXiv.txt |
9804 | astro-ph9804227_arXiv.txt | A comparison is made between the properties of CAL 83, CAL 87, RX J0513.9$-$6951, 1E 0035.4$-$7230 (SMC 13), RX J0019.8+2156, and RX J0925.7$-$4758, all supersoft X-ray binaries. Spectra with the same resolution and wavelength coverage of these systems are compared and contrasted. Some new photometry is also presented. The equivalent widths of the principal emission lines of H and He II differ by more than an order of magnitude among these sources, although those of the highest ionization lines (e.g. O VI) are very similar. In individual systems, the velocity curves derived from various ions often differ in phasing and amplitude, but those whose phasing is consistent with the light curves (implying the lines are formed near the compact star) give masses of $\sim$1.2M$_{\odot}$ and $\sim$0.5M$_{\odot}$ for the degenerate and mass-losing stars, respectively. This finding is in conflict with currently prevailing theoretical models for supersoft binaries. The three highest luminosity sources show evidence of ``jet" outflows, with velocities of $\sim$1--4$\times$10$^3$ km s$^{-1}$. In CAL 83 the shape of the He II 4686\AA\ profile continues to show evidence that these jets may precess with a period of $\sim$69 days. | The close-binary ``supersoft sources" (SSS) are now recognized as a distinct class of very luminous (L$_{bol}\geq10^{38}$ erg s$^{-1}$) X-ray sources characterized by extremely soft X-ray spectra with little or no radiation above $\sim$0.5 keV (e.g. Tr\"umper et al.\ 1991). Several reviews of the observational properties of these sources have recently been published (e.g. Hasinger 1996, Greiner 1996, Kahabka \& van den Heuvel 1997). All SSS appear to have high mass-accretion rates and exhibit long-term X-ray and optical variability which are thought to reflect variations in the rate of mass transfer. In addition, some SSS show evidence of collimated outflows or ``jets" (Crampton et al.\ 1996, hereafter CHC96; Southwell, Livio \& Pringle 1997). Van den Heuvel et al.\ (1992) suggested that the X-ray properties of SSS are best explained by a model involving steady nuclear burning on the surface of a white dwarf accreting material at the Eddington rate. Many observations appear to support this model (Greiner 1996), although alternative interpretations (e.g. Kylafis \& Xilouris 1993) have not yet been ruled out. During a 1996 November CTIO observing run we obtained spectra and some photometry of six close-binary supersoft sources. One of these lies in the Small Magellanic Cloud, 1E 0035.4$-$7230 (hereafter SMC 13). CAL 83, CAL 87, and RX J0513.9$-$6951 (hereafter RX J0513) are all members of the Large Magellanic Cloud, while RX J0019.8+2156 (hereafter RX J0019) and RX J0925.7$-$4758 (hereafter RX J0925) are galactic systems. Since these objects were observed with the same spectrographic configuration, intercomparison of their spectra is very straightforward. In addition, we present previously unpublished photometry for CAL 83 and a few observations of RX J0513 and RX J0925. New data for CAL 87 is being published in a separate paper (Hutchings et al.\ 1998). Long-term spectroscopic and photometric monitoring of these sources is important since the SSS exhibit significant variations over timescales of months and years. A summary of the properties of these six supersoft binaries is given in Table 1, where they are listed in order of decreasing orbital period. The bolometric luminosity (L$_{bol}$) listed is the average value given in Greiner's catalog (1996); it depends strongly on the assumed model, adopted distance, and amount of absorption assumed. | 98 | 4 | astro-ph9804227_arXiv.txt |
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9804 | astro-ph9804157_arXiv.txt | The scattering diameters of \sgra\ and several nearby OH masers ($\approx 1\arcsec$ at 1~GHz) indicate that a region of enhanced scattering is along the line of sight to the Galactic center. We combine radio-wave scattering data and free-free emission and absorption measurements in a likelihood analysis that constrains the following parameters of the GC scattering region: The GC-scattering region separation, $\delgc$; the angular extent of the region, $\psi_\ell$ and $\psi_b$; the outer scale on which density fluctuations occur, $l_0$; and the gas temperature, $\te$. The maximum likelihood estimates of these parameters are $\delgc = 133_{-80}^{+200}$~pc, $0.5\arcdeg \le \psi_\ell \lesssim 1\arcdeg$, and $(l_0/1\,\mathrm{pc})^{2/3}\te^{-1/2} = 10^{-7 \pm 0.8}$. The parameter $\psi_b$ was not well constrained and we adopt $\psi_b = 0\fdg5$. The close correspondence between $\delgc$ and $\psi_\ell\dgc$ suggests that the scattering region encloses the \hbox{GC}. As host media for the scattering, we consider the photoionized surface layers of molecular clouds and the interfaces between molecular clouds and the $10^7$~K ambient gas. We are unable to make an unambiguous determination, but we favor the interface model in which the scattering medium is hot ($\te \sim 10^6$~K) and dense ($n_{\mathrm{e}} \sim 10$~cm${}^{-3}$). The GC scattering region produces a 1~GHz scattering diameter for an extragalactic source of 90\arcsec, if the region is a single screen, or 180\arcsec, if the region wraps around the GC, as appears probable. We modify the Taylor-Cordes model for the Galactic distribution of free electrons in order to include an explicit GC component. We predict that pulsars seen through this region will have a dispersion measure of approximately $2000$~pc~cm${}^{-3}$, of which approximately 1500~pc~cm${}^{-3}$ arises from the GC component itself. We stress the uniqueness of the GC scattering region, probably resulting from the high-pressure environment in the \hbox{GC}. | \label{sec:gc.intro} Davies, Walsh, \& Booth~(1976) established that the observed diameter of \sgra, the compact source in the Galactic center, scales as $\lambda^2$, as expected if interstellar scattering from microstructure in the electron density determines the observed diameter. The observed diameter of \sgra\ is now known to scale as $\lambda^2$ from 30~cm to 3~mm (\cite{rogersetal94}) and to be anisotropic at least over the wavelength range 21~cm to 7~mm (\cite{bzkrml93}; \cite{krichbaumetal93}; \cite{y-zcwmr94}). Maser spots in OH/IR stars within 25\arcmin\ of \sgra\ also show enhanced, anisotropic angular broadening (\cite{vfcd92}; \cite{fdcv94}). These observations indicate that a region of enhanced scattering with an angular extent of at least 25\arcmin\ in radius (60~pc at 8.5~kpc) is along the line of sight to \sgra. At 1~GHz the level of angular broadening produced by this scattering region is roughly 10 times greater than that predicted by a recent model for the distribution of free electrons in the Galaxy (\cite{tc93}, hereinafter TC93), even though this model includes a general enhancement of scattering toward the inner Galaxy. These observations do not constrain the \emph{radial} location of the scattering region for the following reason: All previous observations have been of sources in or near the Galactic center, and for such sources, a region of moderate scattering located far from the Galactic center can produce angular broadening equivalent to that from a region of intense scattering located close to the Galactic center. Previous estimates for the location of the scattering region have ranged from~10~pc to~3~kpc. Ozernoi \& Shisov~(1977) concluded that an ``unrealistic'' level of turbulence is implied unless the region is within 10~pc of the Galactic center. The level of turbulence they considered unrealistic, however, namely $\sqrt{\langle n_{\mathrm{e}}^2\rangle}/\langle n_{\mathrm{e}}\rangle \sim 1$, does appear to occur elsewhere in the interstellar medium (\cite{s91}). Further, van~Langevelde et al.~(1992) used the free-free absorption toward \sgra\ to constrain the region's distance from the Galatic center to the range 0.85--3~kpc, though suitable adjustment of free parameters (outer scale and electron temperature) can decrease the limit to 0.03~kpc. We shall refer to the case in which the region is a site of extreme scattering, $\lesssim 100$~pc from the Galactic center and presumably caused by processes occurring there, as the GC model. We shall refer to the case in which the region is far from the GC, $\gtrsim 1$~kpc and a site of enhanced but not extreme scattering, as the random superposition (RS) model. Although the GC model is attractive for phenomenological reasons, other sites of enhanced interstellar scattering are found throughout the Galaxy (e.g., NGC~6634, \cite{mrgb90}; Cyg~X-3, \cite{mmrj95}) and the mean free path for encountering such a region is approximately 8~kpc (\cite{cwfsr91}). Identifying the location of the scattering is important in establishing the origin of the scattering. Associating the scattering with a specific region may elucidate the mechanism for the generation of the density fluctuations responsible for the scattering. The currently favored mechanism is that velocity or magnetic field fluctuations---or both---generate the density fluctuations (\cite{h84}, 1986; Montgomery, Brown, \& Matthaeus~1987; \cite{s91}; \cite{sg94}; \cite{gs95}). Velocity or magnetic field fluctuations are also a natural means for inducing anisotropy in the density fluctuations and thereby in the scattering disks. If this mechanism is correct, the amplitude of the density fluctuations may provide a measure of the coupling between the density and velocity or magnetic field fluctuations or, more generally, provide information about the small-scale velocity or magnetic field in the scattering region. However, current observational constraints are uncertain by the ratio of the Galactic center-scattering region distance to the Galactic center-Sun distance. In the RS model, the ratio is a few while in the GC model the ratio could be as large as one hundred, so the location of the scattering region is a key free parameter. The location of the scattering region also has implications for pulsar searches toward the \hbox{GC}. Cordes \& Lazio~(1997) showed that even if the RS model is correct, pulsars seen through the scattering region will suffer pulse broadening of at least 5~s at 1~GHz (see also \cite{dwb76}; \cite{os77}). If the GC model is correct, only at frequencies above 10~GHz will pulsations be detectable (because of the $\nu^{-4}$ dependence of pulse broadening) and then only for pulsars with periods longer than 100~ms. In this paper we develop a likelihood analysis to quantify the most probable $\delgc$ for the scattering region. In \S\ref{sec:gc.model} we describe our model for the distribution of free electrons in the \hbox{GC}. In \S\ref{sec:likefunc} we assemble measurements from the literature relevant to radio-wave scattering and develop a likelihood method to constrain the properties of the scattering region, and in \S\ref{sec:gc.conclude} we discuss our results and present our conclusions. | \label{sec:gc.conclude} \subsection{Comparison with Previous Analyses} Isaacman~(1981) surveyed the central $2\arcdeg \times 4\arcdeg$ ($\ell \times b$) of the GC in a search for planetary nebulae. He finds an \emph{excess} number of sources as compared to that expected from extragalactic source counts, an excess he attributes to \ion{H}{2} regions and planetary nebulae. That he finds an excess at all is notable, though, since our analysis predicts that angular diameters of extragalactic sources seen through the GC scattering region will be at least 1\farcm5--3\arcmin. The resolution of his survey was $0\farcm4 \times 2\arcmin$, and he was able to detect sources with angular scales as large as 14\arcmin. Thus, we attribute his excess to the fact that, where his survey overlapped the GC scattering region, it was desensitized by the intense GC scattering to a considerably lesser degree than our survey. Anantharamaiah et al.~(1991) used their observations at~0.327~GHz and a 0.408~GHz $\log N$-$\log S$ relation to conclude that the number of observed extragalactic sources within a 4~deg${}^2$ area centered on \sgra\ is consistent with the number expected from high-latitude source counts. Outside of the inner 1~deg${}^2$, our source counts are also consistent with the expected number of extragalactic sources. Although scattering of other sources, such as B1739$-$298 and B1741$-$312 is heavy, the predicted diameter of these sources is less than 10\arcsec\ at~0.327~GHz, comparable to the size of their beam, so that these sources would not have been resolved out. Gray et al.~(1993) surveyed the Sgr~E region ($\ell = 358.7\arcdeg, b = 0\arcdeg$) at~0.843, 1.45, and~4.86~GHz. At 1.4~GHz the number of sources they find is consistent with that expected from the $\log N$-$\log S$ distribution. Figure~\ref{fig:sclike} shows that the likelihood function for our field 358.9$+$0.5 does not favor a large amount of scattering, i.e., the number of sources in this field is consistent with that expected. Further, two of the sources observed in our VLBI experiment (\cite{lc97}), B1739$-$298 and 1LC~358.439$-$0.211, are in this field. The former is heavily scattered, though not at a level sufficient for it to be seen through the \sgra\ scattering screen. We conclude that our source count results are in good agreement with previous source counts toward the \hbox{GC}. The only exception occurs over the 1~deg${}^2$ region centered on \sgra. This region has not been considered previously or has been subsumed into a much larger area. \subsection{Physical Conditions in the Scattering Region}\label{sec:gc.physical} Our global likelihood, Fig.~\ref{fig:global} and \S\ref{sec:gc.global}, attained a maximum for the following parameter values: $\delgc = 150$~pc, $0.5\arcdeg \le \psi_\ell \lesssim 1\arcdeg$, and $l_0^{2/3}\te^{-1/2} = 10^{-7}$. In this section we consider whether a medium exists within the GC for which such parameter values are plausible. There is a wealth of observational data available for the \hbox{GC}. We shall summarize those conclusions relevant to our study here; interested readers are referred to a number of recent reviews---Genzel, Hollenbach, \& Townes~(1994); Morris \& Serabyn~(1996); and Gredel~(1996)---and references within. Our criteria for the host medium of the density fluctuations are that the medium must have a sufficient density and that it must be capable of sustaining density fluctuations of the requisite magnitude. We establish our first criterion by estimating $\nbar$ from the scattering diameters of GC sources using equations~(\ref{eqn:measures}) and~(\ref{eqn:smweight}). The diameters of \sgra\ and the OH masers require a weighted scattering measure of $S \approx 10^2$~\smu, equation~(\ref{eqn:smweight}). Eliminating SM between equations~(\ref{eqn:measures}) and~(\ref{eqn:smweight}) and solving for $\nbar$ yields \begin{equation} \nbar \sim 10^3\,\mathrm{cm}^{-3}\frac{1}{\varepsilon\sqrt{f}}\left(\frac{l_0}{1\,\mathrm{pc}}\right)^{1/3}\left(\frac{\delgc}{150\,\mathrm{pc}}\right)^{-3/2}. \label{eqn:delne} \end{equation} Two factors could alter this estimate by about an order of magnitude. First, it is likely that $l_0 \ll 1$~pc, which would \emph{reduce} our estimate of $\nbar$. Second, as we noted in \S\ref{sec:gc.global}, the similarity between the values of $\delgc$ and $\psi_\ell\dgc$ suggests that the density fluctuations fill the region and $f \approx 1$. However, we might also associate $l_0$ with the characteristic size of a scattering cloudlet within the region. If the region contains few such cloudlets and $\delgc/l_0 \gg 1$, then $f \ll 1$, and our estimate above would be a considerable \emph{underestimate}. Yusef-Zadeh et al.~(1994) estimated that a typical line of sight might intersect only 10 or so scattering cloudlets. In any event, we conclude that the scattering medium must be dense, $n_{\mathrm{e}} \gtrsim 10^2$~cm${}^{-3}$. For comparison, Spangler~(1991) concludes that $n_{\mathrm{e}} \sim 1$~cm${}^{-3}$ for scattering regions in the Galactic disk. Our second criterion for the host medium is that it must be able to support density fluctuations of the required magnitude. This constraint has been lucidly reviewed by Spangler~(1991): The density fluctuations are presumed to arise from plasma turbulence. As this turbulence dissipates, it cannot heat the host medium at a rate that exceeds the medium's cooling capacity. This constraint is particularly acute in the situation we are proposing as the dissipation mechanisms considered by Spangler~(1991) scale as $l_0^{-a}$ with $a \approx 1$. Since we are considering $l_0 < 1$~pc, the heating rates could be excessive. The dominant damping mechanisms for $l_0 < 1$~pc are linear Landau damping, ion-neutral collisions, and a parametric decay instability. The first two mechanisms scale as $l_0^{-2/3}$ while the latter scales as $l_0^{-1}$. In addition to their dependence on $l_0$, the damping rates depend on the large scale magnetic field, $\Gamma \propto B^2$; the Alfv\'en wave speed, $\va$; and the amplitude of the magnetic fluctuations, $\Gamma \propto (\delta B/B)^2$ for linear Landau damping and ion-neutral collisions while $\Gamma \propto (\delta B/B)^3$ for the parametric decay instability. Linear Landau damping also depends upon the angle of propagation with respect to the direction of $\mathbf{B}$, $\chi$, and the plasma $\beta$. For values of these quantities appropriate for scattering regions in the Galactic disk, these damping mechanisms produce volumetric heating rates of $\Gamma \sim 10^{-23.5}$--$10^{-21.5}$~erg~s${}^{-1}$~cm${}^{-3}$. As Spangler~(1991) discussed, there are also a number of simplifications and additional assumptions which enter the calculation of these heating rates. Inferred magnetic field strengths in the GC are $B \sim 1$~mG, or $10^3$ that of the field strength in the disk. To estimate $\delta B/B$, we use (Cordes, Clegg, \& Simonetti~1990) \begin{equation} \frac{\delne}{n_{\mathrm{e}}} \sim \left(\frac{\delta B}{B}\right)^c \label{eqn:bB} \end{equation} with $c = 1$ for linear processes and $c = 2$ for non-linear processes like the parametric decay instability. Our first criterion for the host medium is that $\delne/n_{\mathrm{e}} \le 1$. For definiteness, and to provide the largest possible value of the heating, we take $\delta B/B \sim 1$. Finally, although $B$ is much larger in the GC than in the Galactic disk, $n$ is also larger. As a result $\va$ is larger than in the disk, but probably by no more than an order of magnitude. Thus, we expect $\Gamma$ in the GC to be about a factor of $10^7$ larger that in the Galactic disk. The heating rate from linear Landau damping in scattering regions in the Galactic disk (\cite{s91}) assumes the density fluctuations arise from obliquely propagating magnetosonic waves ($\chi \approx 6\arcdeg$). More aligned propagation results in less damping. It is not clear if the GC environment would favor highly aligned propagation or not. The large values of $B$ in the GC are inferred, in part, from the system of non-thermal filaments and threads seen throughout the \hbox{GC}. With only one exception, these filaments have no kinks or bends in them, even though they are observed to be interacting with molecular clouds having typical velocities of 10--100~km~s${}^{-1}$ (\cite{ms96}). This rigidity could be an indication that only highly aligned propagation is allowed. If $\chi$ is highly concentrated near 0\arcdeg, then the heating from linear Landau damping would be unimportant and the heating rates could be two orders of magnitude lower than those quoted above. Alternately, as Spangler~(1991) noted, the distribution of $\chi$ could be isotropic, but waves with large $\chi$ would then damp quickly and the heating rate will be unchanged or even larger than what we assume. The presence of small-scale ($\approx 0.1$~pc) magnetoionic cloudlets in the GC has already been inferred to explain large changes in the Faraday rotation measure of certain features (G~359.1$-$00.2, the ``Snake,'' \cite{gnec95}; G~359.54$+$0.18, the non-thermal filaments, Yusef-Zadeh, Wardle, \& Parastaran~1997), though the inferred density in these cloudlets, 0.3--10~cm${}^{-3}$, is less than our nominal estimate. In the Galactic disk a small body of observational evidence suggests that the magnetoionic medium responsible for Faraday rotation is also responsible for scattering and pulsar dispersion (\cite{sc86}; Lazio, Spangler, \& Cordes~1990; \cite{ahmsrs96}), and the same may be true in the \hbox{GC}. We now consider two models for the host medium. In both models, the scattering arises in thin layers on the surfaces of molecular clouds. Even if the filling factor, $f$, of these layers is not large, the \emph{covering factor}, i.e., the probability that a line of sight through the GC will intersect one of these layers, can still be close to unity. \subsubsection{Photoionized Surfaces of Molecular Clouds}\label{sec:gc.skins} Over the region $|\ell| \lesssim 1\fdg5$ and $|b| \lesssim 0\fdg5$, $n_{\mathrm{e}} \sim 10$~cm${}^{-3}$, as determined from single-dish recombination line and total intensity measurements (Matthews, Davies, \& Pedlar~1973; \cite{mp79}). Embedded within this large-scale region are smaller regions of much higher electron densities, $n_{\mathrm{e}} \sim 10^3$--$10^5$~cm${}^{-3}$, primarily within Sgr~A West and Sgr~B2 (e.g., \cite{mpgy-z93}). Some of these high density regions are the photoionized surfaces (size $\sim 10^{-4}$~pc) of molecular clouds ($n \gtrsim 10^4$~cm${}^{-3}$) irradiated by the ambient radiation field (effective temperature $\approx 35\,000$~K). Yusef-Zadeh et al.~(1994) identified these molecular skins as the source of the scattering and associated their thicknesses with the outer scale, $l_0$; Gray et al.~(1995) suggested that the magnetoionic medium responsible for the Faraday rotation toward G~359.1$-$00.2 also results from these molecular clouds. This model suffers from at least three potential difficulties. First, the molecular skins are photoionized by a radiation field having a temperature of $T_{\mathrm{eff}} \sim 10^4$~\hbox{K}. Our constraint on $l_0^{2/3}\te^{-1/2}$ therefore requires $l_0 \sim 10^{-7.1}$~pc. This value is considerably smaller than that derived by Yusef-Zadeh et al.~(1994) for the ionized molecular skins. However, in deriving their value for $l_0$, Yusef-Zadeh et al.~(1994) used a value of the ionizing flux appropriate to the inner few parsecs. The stellar density decreases as $r^{-2}$, so outside the inner few parsecs, the ionizing flux should be lower than that assumed by Yusef-Zadeh et al.~(1994). A lower ionizing flux would result in a smaller skin depth and bring their estimate and our estimate of $l_0$ into better agreement. A second potential difficulty with this model is that the medium would only barely be capable of cooling itself. If the outer scale is $l_0 \sim 10^{-7}$~pc, then the heating rate from the damping of the plasma turbulence is $\Gamma \sim 10^{-13}$--$10^{-12}$~erg~cm${}^{-3}$~s${}^{-1}$. The cooling capacity of gas near $\te \sim 10^4$~K depends sensitively upon the fractional ionization and temperature (\cite{dm72}). We estimate that a density of $n_{\mathrm{e}} \gtrsim 10^5$~cm${}^{-3}$ is required for these skins to be able to cool sufficiently in order to support the density fluctuations. This density is at the upper end of the range $10^3$--$10^5$~cm${}^{-3}$ inferred for the small-scale \ion{H}{2} regions. In determining the cooling function of the medium, we have used results that assume a solar abundance. The metallicity in the GC could be as much as twice solar, leading to an increased cooling efficiency. The third difficulty is that the required value for $l_0$ in this model, $l_0 \sim 3 \times 10^{11}$~cm, is considerably smaller than that in the Galactic disk. In the Galactic disk, a stringent lower limit on the outer scale is $10^{13}$~cm, and it may be as large as $10^{18}$~cm (\cite{ars95}). Although the physics for the generation and maintenance of small-scale density fluctuations is not well understood, we regard it as potentially troublesome that this model predicts such a small $l_0$. \subsubsection{``Warm'' Interfaces}\label{sec:gc.interface} X-ray observations have revealed a central, diffuse X-ray source with a (FWHM) size of $1\fdg8 \times 0\fdg9$ (\cite{ykkkt90}). Frail et al.~(1994) suggested that this X-ray emitting gas may be responsible for the scattering. Yusef-Zadeh et al.~(1997) suggested that this X-ray emitting gas is also the magnetoionic medium responsible for the Faraday rotation toward G~359.54$+$0.18. The density and temperature of this region are estimated at 0.05~cm${}^{-3}$ and $10^7$--$10^8$~\hbox{K}. This region cannot itself be the host of the density fluctuations because of its low density, cf.~equation~(\ref{eqn:delne}). However, this gas appears spatially coincident with the central zone of intense molecular emission and presumably abuts cooler gas in the clouds. We modify Frail et al.'s~(1994) proposal by identifying the interfaces where the GC molecular clouds are exposed to this ambient hot medium as the source of the scattering. We term these interfaces ``warm'' by analogy with McKee \& Ostriker's~(1977) model for the \hbox{ISM}. In that model cold clouds immersed in a hot ($10^6$~K) medium have $10^4$~K interfaces. In the GC densities and temperatures are 1--2 orders of magnitude higher, but we expect that clouds will still develop intermediate temperature interfaces. This model suffers from two of the same difficulties as the previous model. The X-ray emitting gas and molecular clouds appear to be in rough pressure equilibrium with $P \sim 5 \times 10^6$~K~cm${}^{-3}$ (\cite{bs91}). Even though there are supersonic motions \emph{within} the clouds, the clouds themselves ($v \sim 10$--100~km~s${}^{-1}$) are moving subsonically with respect to the hot medium ($c_{\mathrm{sound}} \sim 1000$~km~s${}^{-1}$). Taking pressure balance to extend throughout the interface region, we find a density $n_{\mathrm{e}} \sim 5$--50~cm${}^{-3}$ for $\te \sim 10^5$--$10^6$~\hbox{K}. From our likelihood results, the temperature within these interfaces implies an outer scale of $l_0 \sim 10^{-6.5}$--$10^{-6}$~pc, which, in turn, implies an rms density, equation~(\ref{eqn:delne}), of $\delne \sim 10$~cm${}^{-3}$. However, the cooling capacity of this medium is only $10^{-20}$~erg~cm${}^{-3}$~s${}^{-1}$. The predicted heating rate is $\Gamma \sim 10^{-13}$~erg~cm${}^{-3}$~s${}^{-1}$. The outer scale in this model remains troublesomely small. If the size of the interface region is set by thermal conduction, the portion of the interface with $\te < 10^6~K$ has a size $\lesssim 10^{-1}$~pc (\cite{mc77}). Clearly if not all of the interface contributes to the scattering better agreement would be obtained between $l_0$ and the interface size. Still, the outer scale remains an order of magnitude smaller than its lower limit in the Galactic disk. One point in favor of this model is the distribution of the X-ray emitting gas as compared to that of the molecular clouds. The size of the X-ray emitting region is similar to the extent of the scattering region, approximately 1\arcdeg. In contrast, the molecular cloud distribution extends over the range $-1\arcdeg \lesssim \ell \lesssim 2\arcdeg$. If the scattering traced massive stars within these molecular clouds, as the photoionized molecular cloud skins model suggests, the scattering should extend further in longitude than it does. In this respect, the lack of enhanced scattering for \hoh\ masers in Sgr~B is particularly problematic. We stress the importance, and probable uniqueness, of the high density in the \hbox{GC}. In the Galactic disk density fluctuations cannot be supported in media with $\te \sim 10^6$~K, a position with both theoretical and limited observational support (\cite{s91}; \cite{pc92}). Similarly, recent VLBI observations of 5 pulsars show no evidence for an enhanced level of turbulence at the boundary of the Local Bubble (\cite{cr87}; Britton, Gwinn, \& Ojeda~1996), potentially a local analog of an interface between a hot and cooler medium. However, the Local Bubble and ambient medium have densities a factor of $10^2$--$10^3$ smaller than that in the \hbox{GC}. In summary, we use our likelihood results, \S\ref{sec:gc.global}, to constrain host media for the scattering material. Potential media include the photoionized skins of molecular clouds or the interface regions between the clouds and the ambient X-ray emitting gas. There are difficulties with both models: Both models overpredict the outer scale and appear to have some trouble supporting the required level of density fluctuations. Although we have been unable to make an unambiguous identification of the scattering medium with either medium, we favor the interface model, in part, because it shows a better correspondence between the spatial distribution of scattering and proposed host medium. \subsection{Modification of the Taylor-Cordes Model}\label{sec:tcmodel} The TC93 model modelled the global distribution of free electrons in the Galaxy with four components: an extended component, an inner Galaxy component, a component confined to the spiral arms, and a component local to the Gum Nebula. We now extend the TC93 model to include a GC component\footnote{ As in TC93, the coordinate system has the $x$-axis directed parallel to $\ell = 90\arcdeg$, the $y$-axis toward $\ell = 180\arcdeg$, the $z$-axis toward $b = 90\arcdeg$, and $R = \sqrt{x^2 + y^2}$ is the Galactocentric radius.} (cf.\ eqn.~[11] of TC93): \begin{eqnarray} n_{\mathrm{e}}(x,y,z) & = & n_1(R,z) + n_2(R, z) + \sum_{j=1}^4 n_{\mathrm{arm},j}(x,y,z) + n_{\mathrm{Gum}}(x,y,z) \nonumber \\ & + & n_{\mathrm{GC}}g_{\mathrm{GC}}(R)h_{\mathrm{GC}}(z). \end{eqnarray} The first four components are discussed at length in TC93. We focus on only the last component, that toward the \hbox{GC}. Based on the estimate in equation~(\ref{eqn:delne}) and our estimates for $l_0$, we take $n_{\mathrm{GC}} = 10$~cm${}^{-3}$. Our estimate for $\delgc$ is $\delgc = 150$~pc. Heretofore, we have been treating the scattering region as a screen with sharp boundaries. It is more likely that the region has soft edges. We therefore adopt a radial dependence of \begin{equation} g_{\mathrm{GC}}(R) = \exp\left[-(R/0.150\,\mathrm{kpc})^2\right]. \end{equation} The latitude (or $z$) dependence of the screen is less well constrained. For definiteness we take \begin{equation} h_{\mathrm{GC}}(z) = \exp\left[-(z/0.075\,\mathrm{kpc})^2\right] \end{equation} corresponding to $\psi_b = 0\fdg5$. The resulting axial ratio for the electron density distribution is 0.5; the axial ratio for the X-ray distribution is also 0.5 and that of the molecular cloud distribution is 0.3 (\cite{ms96}). In the TC93 model the relationship between the free electron density and the scattering measure produced by a line of sight of length $ds$ through those electrons is $d\mathrm{SM} \propto F n_{\mathrm{e}}^2ds$. The parameter~$F$ is \begin{equation} F = \frac{\zeta\epsilon^2}{f}\left(\frac{l_0}{1\,\mathrm{pc}}\right)^2, \label{eqn:F} \end{equation} where $\zeta$ is the normalized variance of electron density fluctuations between cloudlets and their surroundings and $\epsilon$ and~$f$ are as in equation~(\ref{eqn:measures}). Taking $\zeta \sim \epsilon \sim 1$, our estimates for $l_0$ imply $F \gtrsim 10^4$. Both of the models we have considered here have $f < 1$. For definiteness, we take $f \sim 0.1$, recognizing that this may be an upper limit on~$f$. We therefore conclude $F \gtrsim 10^5$. For comparison, the parameter~$F$ has a value of~0.4 in the solar neighborhood, 6 in spiral arms, and 40 in the Galaxy's inner few kiloparsecs. A value of $F \sim 10^5$ produces an SM comparable to that suggested by the scattering diameters of \sgra\ and the OH masers. Their scattering diameters require a line-of-sight weighted scattering measure of $S \approx 10^2$~kpc~m${}^{-20/3}$. Our results suggest $\delgc \approx 150$~pc; correcting for the line-of-sight weighting, equation~(\ref{eqn:smweight}), implies that the GC has a scattering measure of $\mathrm{SM} \sim 10^{5.5}$~kpc~m${}^{-20/3}$. Integrating the TC93 expression for $d\mathrm{SM}$ through the GC, with $F \sim 10^5$, we find $\mathrm{SM} \sim 10^6$~kpc~m${}^{-20/3}$. Since the GC component is so localized, only for lines of sight through the GC are the results of TC93 altered. In this direction, however, the TC93 model underpredicts various quantities by a large amount. In the TC93 model, GC pulsars have $\mathrm{DM} \approx 600$--800~pc~cm${}^{-3}$; we predict that the DM will be somewhat larger, $\mathrm{DM} \approx 2000$~pc~cm${}^{-3}$, with approximately 1500~pc~cm${}^{-3}$ of that arising from the GC component itself. For comparison, the largest DM known is for PSR~B1758$-$23 with $\mathrm{DM} = 1074$~pc~cm${}^{-3}$ (Manchester, D'Amico, \& Tuohy~1985; \cite{klmjds93}). Further, from Cordes \& Lazio~(1997) the temporal broadening of pulses from pulsars seen through this region will be $350\,\nu_{\mathrm{GHz}}^{-4}$~\emph{seconds}, requiring high-frequency ($\nu \approx 10$~GHz) periodicity searches to detect pulsations. Although the DM we predict for GC pulsars is substantial, the dispersion smearing across a 1~GHz bandpass at 10~GHz ($\approx 5$~ms) is comparable to the pulse broadening, so that only a small number of filterbank channels, e.g., 16, would be necessary to combat the dispersion smearing. \subsection{Conclusions} We use a likelihood analysis to determine the following parameters of the GC scattering region: The GC-scattering region separation, $\delgc$; the angular extent of the region, $\psi_\ell$ and $\psi_b$; the outer scale on which density fluctuations occur, $l_0$; and the gas temperature, $\te$. \begin{itemize} \item From the literature we have assembled a list of all sources toward the GC for which angular broadening has been measured. A subset of these sources is OH/IR stars, for which the spatial distribution about the GC is known. We construct a likelihood function for the angular broadening of OH/IR stars, utilizing this distribution, \S\ref{sec:broaden}. Masers within approximately 1\arcdeg\ of \sgra\ have diameters consistent with $\delgc \approx 150$~pc (Fig.~\ref{fig:aggregate_as}). \item The likelihood analysis of our source counts, \S\ref{sec:counts}, indicates that a deficit of sources occurs within approximately 1\arcdeg\ of \sgra\ and is caused by a scattering region within 500~pc of \sgra\ (Fig.~\ref{fig:aggregate_sc}). The resulting scattering diameter, at least 20\arcsec\ at~1~GHz, causes extragalactic sources to be so broad as to be resolved out by our observations. \item \hoh\ masers in and an extragalactic source near Sgr~B and an extragalactic source near Sgr~C do not show the extreme scattering of sources closer to \sgra, indicating that the scattering region does not extend to more than 1\arcdeg\ in longitude. The latitude extent of the scattering region is poorly constrained, but is no more than 1\arcdeg. \item From the literature we have estimated the free-free emission and absorption toward five heavily scattered masers near \sgra. The likelihood function is dominated by the free-free emission. The relevant parameters, $\delgc$, $l_0$, and $\te$, are not independent for this likelihood and only the product $\delgc^{-2}l_0^{2/3}\te^{-1/2}$ can be constrained (Fig.~\ref{fig:aggregate_ff}). \end{itemize} The global likelihood, formed by multiplying the individual likelihoods, is shown in Fig.~\ref{fig:global}. The maximum likelihood estimates of the parameters are $\delgc = 150$~pc, $0.5\arcdeg \le \psi_\ell \lesssim 1\arcdeg$, and $l_0^{2/3}\te^{-1/2} = 10^{-7}$ with $l_0$ in pc and $\te$ in \hbox{K}. The parameter $\psi_b$ was not well constrained and we adopt $\psi_b = 0\fdg5$. The close correspondence between $\delgc$ and $\psi_\ell\dgc$ suggests that the scattering region encloses the \hbox{GC}. The GC scattering region produces a 1~GHz scattering diameter of 90\arcsec, if the region is a single screen, or 180\arcsec, if the region wraps around the GC, as appears probable. We modify the Taylor-Cordes model for the Galactic distribution of free electrons in order to include an explicit GC component. We predict that pulsars seen through this region will have a dispersion measure of approximately $2000$~pc~cm${}^{-3}$, of which approximately 1500~pc~cm${}^{-3}$ arises from the GC component itself, and suffer pulse broadening of $350\,\nu_{\mathrm{GHz}}^{-4}$~\emph{seconds}; pulsations will be detected only for frequencies above 10~GHz (\cite{cl97}). As host media for the scattering we consider the photoionized surface layers of molecular clouds and the interfaces between molecular clouds and the $10^7$~K ambient gas. We identify the host medium by requiring that it be sufficiently dense to support density fluctuations of the required magnitude. We are unable to make an unambiguous determination, but we favor the interface model which predicts that the scattering medium is hot ($\te \sim 10^6$~K) and dense ($n_{\mathrm{e}} \sim 10$~cm${}^{-3}$). The X-ray interface model also shows better spatial agreement, when compared to the photoionized skin model, with the region over which the scattering is observed. This model is summarized graphically in Fig.~\ref{fig:gc.summary}. The GC scattering region is likely to be unique in the Galaxy, probably because it is a high-pressure environment and can sustain densities and temperatures much higher than in the Galactic disk. | 98 | 4 | astro-ph9804157_arXiv.txt |
9804 | astro-ph9804294_arXiv.txt | { Lithium is one of the few primordially produced elements. The value of the primordial Li is taken to be that observed in metal--poor dwarfs, where it is not contaminated by stellar Li sources which act on longer time scales. The atmospheric abundance is currently derived from the LiI $\lambda\lambda 6707 \AA~$ resonance transition and the validity of the models employed has been questioned \markcite{k95} (Kurucz 1995). In this letter we report the first detection of the Li I $\lambda\lambda 6104 \AA~ 2^2P - 3^2D$ subordinate transition in the prototype population II star HD~140283. The same Li abundance of (Li/H) $=1.4\times 10^{-10}$ is found consistent with both the resonance and subordinate lines. The two lines form at different depths in the atmosphere implying that the 1-D homogeneous atmospheric models used in the abundance determination are essentially correct. When coupled with the standard big bang yields, the Li in the halo dwarfs provides two solutions for the baryon-to-photon ratio $\eta_{10}= n_{b}/n_{\gamma} \times 10^{10}$ and for the present baryon density $\Omega_b h_{70}^2=0.0748\eta_{10}$: a) a first solution at $\eta_{10}\approx 1.8$, consistent with the $\eta_{10}$ implied by the high deuterium values $D/H\approx 2\times 10^{-4}$ observed in some quasar absorption systems \markcite{webb} (Webb et al 1997) and b) a second solution at $\eta_{10}$ $\approx$ 4 which is consistent, within the errors, with the low deuterium D/H =$3.4\times 10^{-5}$ measured in other quasar absorption systems\markcite{burles} (Burles \& Tytler 1998). } | Lithium, together with D and $\rm ^{3,4}He$, is one of the few elements produced by nuclear reactions in the first minutes after big bang\markcite{wfh} (Wagoner, Fowler \& Hoyle 1967). The observations of these elements and their extrapolation to the primordial values are consistent with the predictions of the standard primordial nucleosynthesis providing, together with the relic radiation and the expansion of the Universe, a robust support to the big bang theory. Recently, additional support to the primordial nature of Li in halo dwarfs has come from the observations of Li in metal-poor stars of the thick disc\markcite{mbp97} (Molaro, Bonifacio \& Pasquini 1997). This population is chemically and kinematically distinct from the halo, but has the same Li abundance of the halo. Minniti et al (1997) \markcite{min97} claimed detection of Li, at the plateau level, in a metal-rich, but old star, belonging to the Galactic Bulge. Finally Li at the plateau level has also been detected in a star which was possibly born in an external galaxy and then accreted by the Milky Way\markcite{mol97} (Molaro 1997). So far the Li abundance has been always obtained only from the analysis of the Li I $\lambda\lambda$ 6707 \AA~ resonance doublet. This is not a very comfortable situation in the light of the importance of the determination of lithium abundances in stars for primordial nucleosynthesis, stellar structure and chemical evolution. Our ability to determine the Li abundance using simple plane-parallel homogeneous atmospheres, has been recently debated\markcite{k95,kis97,gp97} (Kurucz 1995; Kiselman 1997; Gadun \& Pavlenko 1997). The analysis of several lines, which sample different depths in the stellar atmosphere is crucial to test the correctness of the modelling. The one dimensional, homogeneous, static models which are currently employed may arise concern because they ignore the fine structure and hydrodynamic phenomena such as granulation which are seen on the Sun. The Li I $\lambda\lambda$ 6707 \AA ~resonance transition is the only one readily available to spectroscopic observation. The strongest subordinate line at 6104 \AA ~ is much fainter and blended with Fe I line and has been so far detected only in young T Tauri stars (Hartigan et al 1989)\markcite{hart} and Li-rich giants (Merchant 1967, Wallerstein \& Sneden 1982) \markcite{merchant,ws}, were Li is more than about 1 dex more abundant owing to the Galactic Li production. \par In this letter we report the detection of the Li I $\lambda\lambda$ 6104 \AA ~transition in the spectrum of the metal--poor star HD140283. Both this line and the resonance line are consistent with the computations made using a one dimensional, homogeneous model atmosphere, thus increasing our confidence that this model represents a satisfactory average of the complex fine structure expected in metal--poor stars. The use of Li observed in halo dwarfs as an indicator of primordial abundance rests on the absence of any Li depletion. Depletion is predicted by non-standard models which take into account rotational mixing \markcite{pin92} (Pinsonneault, Deliyannis \& Demarque 1992) or diffusion \markcite{vau95} (Vauclair \& Charbonnel 1995), but these models predict a downturn of the hot side of the Li plateau and considerable dispersion. It seems that neither the downturn nor the large dispersion is present in the observations, which suggests that diffusion or rotational mixing do not affect significantly the Li observed at the stellar surface of metal--poor dwarfs. However, the downturn can be very small ($\approx 0.2$ dex) for the purely diffusive case and a suitable choice of the mixing length parameter ($\alpha=1.5$, see Deliyannis et al 1990\markcite{deli90}) and the issue of intrinsic dispersion remains rather controversial with some positive claims. Ryan et al (1996)\markcite{ryan} identify a triplet of stars (G064-012, G064-037, CD -33$^\circ$ 1173) with similar colors, but different Li abundances by a factor of 2.5. Then there is the case of star BD+23 3912 which has a [Fe/H]$\approx -1.3$ to $-1.5$ and a Li abundance which is about 0.20-0.36 dex higher than the plateau (Rebolo et al 1988, King et al 1996\markcite{reb88,king96}). Moreover Boesgaard et al (1998)\markcite{boesgaard} find differences of up to $\approx 0.5$ dex among seven subgiants of M92 but the same objects show other chemical peculiarities, namely [Mg/Fe] is 0.55 dex lower and [Na/Fe] is 0.76 dex larger than in HD140283 (King et al 1998) \markcite{king98} . | The LiI line forming regions lie in the upper part of the atmospheric convective zone where Li is mostly ionized due to its low ionization potential. This is why the determination of precise Li abundances requires accurate observations, accurate stellar effective temperature and an appropriate modeling of the atmosphere of a metal poor star. The model-atmospheres employed are one dimensional (1-D), with plane parallel geometry and ignore any inhomogeneity effect, such as granulation. Qualitative computations, based on a two-stream model atmosphere, suggested that the abundance of Li in halo dwarfs could be underestimated by as much as a factor of 10\markcite{k95} (Kurucz 1995), but more recent calculations based on 2-D\markcite{gp97} (Gadun \& Pavlenko 1997) and 3-D\markcite{kis97} (Kiselman 1997) atmospheric models show that effects of granulation on the LiI lines are much less important. Granulation effects in the atmosphere have a depth dependence and this should produce different effects in the resonance and subordinate doublets. As can be seen from figures 1 and 2, the same Li abundance reproduces satisfactorily both the 6104 \AA ~and the 6707 \AA ~doublets. The two transitions form at different depths in the stellar atmosphere: unit optical depth at wavelength 6707.761 \AA~ is attained at $log (\tau_{Ross})\approx -0.57$, corresponding in our model to a local temperature of 5235 K, while at wavelength 6103.649 \AA~ it is already attained at $log (\tau_{Ross})\approx -0.09$, or T=5915 K. The resonance line receives contributions from a more extended region than the subordinate line. Unit optical depth at the wavelength at which the residual intensity is 0.999, is attained at $log (\tau_{Ross})\approx -0.11$ for the resonance line, but at $log (\tau_{Ross})\approx -0.08$ for the subordinate line. Thus the subordinate line samples deeper and hotter layers than the resonance line, as shown in Fig. 3. \par The lower level of the Li 6104 \AA~ transition is the upper level of the 6707 \AA $2^2S-2^2P$~ transition. Our synthetic spectra are computed under the LTE assumption and the consistency between the two lines implies a correct computation of the populations of the 2S, 2P and 3D levels. This is in agreement with the theoretical estimations which predict relatively small corrections for NLTE effects in the LI 6707 \AA ~line \markcite{carl94,pavmag} (Carlsson et al 1994; Pavlenko \& Magazz\`u 1996). Thus the detection of a subordinate LiI line, and its consistency with the abundance derived from the resonance 6707 \AA~ doublet, provides support to the correctness of this Li abundance. The consistency of the abundances based on the LiI 6707 \AA~ and 6104 \AA~ transitions observed in HD 140283 supports the Li abundances measured in the population II stars, using 1-D model atmospheres, in the last decades. The new generation of large telescopes will allow to measure the Li 6104 \AA ~Li I subordinate doublet in other much fainter population II stars, thus permitting to verify this consistency on the grounds of a statistically significant sample, and ultimately achieve a more accurate measurement of the primordial Li abundance. \par Among the light elements produced in the first minutes after the big bang, Li is the only one which shows a non monotonic behaviour with $\eta_{10}$, the so-called {\sl Li-valley}, which reflects the different nuclear reactions which synthesize Li at different baryonic densities. The most recent measurement of the Li primordial abundance is $\rm (Li/H)=1.73 \pm 0.05_{stat} \pm 0.2_{sys} \times 10^{-10}$\markcite{bm} (Bonifacio \& Molaro 1997), which is the mean value of 41 halo stars for which precise effective temperatures, determined by means of the infrared flux method \markcite{alonso} ({Alonso}, {Arribas}, \& {Martinez-Roger} 1996), were available. The systematic errors, which dominate the error budget, come from a possible offset of $\pm$ 75 K in the zero point of the temperature scale of cool stars. This Li abundance intercepts the primordial yields for two different values of $\eta_{10}$, which unfortunately do not help in resolving the deuterium and helium controversies. Each solution for $\eta_{10}$ obtained from Li is consistent with either the high-deuterium/low-helium\markcite{webb,olive} (Webb et al 1997; Olive, Steigman \& Skillman 1997) or the low-deuterium/high-helium \markcite{burles,izo97} (Burles \& Tytler 1998; Izotov, Thuan \& Lipovetsky 1997). The lower $\eta_{10}$ requires considerable D destruction to match the presently observed abundance in the local interstellar medium of the Galaxy. The higher $\eta_{10}$ value is also consistent with the low deuterium (D/H $=3.9 (\pm 1)\times 10^{-5}$) derived from the 92 cm hyperfine transition emission towards the unprocessed Galactic anti-center \markcite{chen97} (Chengalur, Braun \& Butler Burton 1997). | 98 | 4 | astro-ph9804294_arXiv.txt |
9804 | astro-ph9804249_arXiv.txt | Modifications in Friedmann-Lema\^itre-Robertson-Walker (FLRW) Hubble diagrams caused by mass density inhomogeneities are used to illustrate possible effects on a determination of the mass parameter $\OM$ and the cosmological constant $\Lambda$. The values of these parameters inferred from a given set of observations depend on the fractional amount of matter in inhomogeneities and can differ from those obtained by using standard FLRW Hubble diagrams by as much as a factor of two. | The pressure-free FLRW models of GR have, for a very long time, been assumed to adequately describe the large scale geometry of our universe but only recently have observational techniques become available which promise to determine values for all three of their requisite parameters $\{H_0,\ \OM$, and $\OL\equiv {\Lambda c^2/(3H_0^2)}\}$ and implicitly test this assumption. As an example, corrected magnitudes and redshifts ($m$-$z$) for Type Ia supernovae (SNe Ia) are measured, plotted, and compared with theoretical $m(H_0,\ \OM,\OL; z)$ curves computed for the FLRW models \cite{PS1,PS2,GP}. Because these models are isotropic and homogeneous, and our universe appears quite inhomogeneous, modifications in the FLRW predictions have long been proposed and estimated \cite{YZ,KR}. When a wide angle measurement of the CMB is made the average mass density of the universe (the FLRW value for $\rho_0$) likely exists within the radio beam collected by the antenna. However, when small objects such as SNe Ia are observed at $z< 1$, a mass density significantly less than the average is `likely' to be in the observing optical path. In particular, if the underlying mass density approximately follows luminous matter (\ie associated with bounded galaxies) then effects of a diminished mass density in the observing beams on relations like $m(\OM,\OL;z)$ are important. The majority of currently observed SNe Ia are not being seen through foreground galaxies and whether or not this is due to selection (rather than statistics) is not important. If the objects observed do not have the average FLRW mass density $\rho_0$ in their foregrounds then the FLRW \mz\ relation does not apply to them. Ultimately some SNe Ia should exist behind foreground galaxies and for these, \mz\ should be computed using the lensing formulas. These formulas \cite{BR,CJ} contain source-observer, deflector-observer, and source-deflector distances, respectively $D_s, D_d$, and $D_{ds}$, all of which depend on the mass density in the observing beam, {\bf excluding} the deflector. These distances will not be given by the standard FLRW result if the observing beam contains less than the average FLRW mass density, but instead given by the `intergalactic' distance discussed here. | Numerous arguments have been made since the '60s against the existence of any effect on apparent magnitudes such as given here (see \cite{KR2} for current refs. to dissenting opinions). These are almost all weak-lensing arguments, valid as long as density perturbations don't produce multiple images (or absorb photons), \ie the FLRW result will coincide with the theoretical mean of weak lensing observations. Even if weak lensing arguments can be extended to strong lensing perturbations, \eg by not resolving separate images, the theoretical mean of a distribution of magnitudes may be of little use in determining $\OM$ or $\OL$. If mass is as inhomogeneous as luminosity, the distribution of magnitudes is expected to be so skewed as to make the mean statistically insignificant. The `most likely' value should be far more useful in a determination of the cosmic parameters. However, to determine the most likely \mz not only requires knowledge of the average mass density $\rho_0$, it requires the modeling of galaxies masses etc. What can be determined by fixing a single additional parameter $\nu$ (which proportions total FLRW mass density into an intergalactic component and a galactic component) is the `intergalactic' Hubble curve. This \mz can be much more useful in determining $\OM$ or $\OL$ than the mean Hubble curve, \eg if galaxies are not larger than they appear optically, the distribution of magnitudes (at a given redshift) is expected to peak much closer to the intergalactic value than to the mean value (see the numerical work \cite{HD}). Additionally the effects of galaxy lensing can be controlled by simple selection. All that is required is that the observed SNe Ia are separated into those with foreground galaxies and those without. Those without should be fit to the intergalactic Hubble curve and those with (when any are found) could be included by correcting for lensing and/or by averaging in with the others to see if the weak-lensing FLRW value can be obtained. Only with enough unbiased and absorption corrected data will the FLRW Hubble curve be useful. The intergalactic Hubble curve contains the additional parameter $\nu$; however, if luminous matter is essentially the whole story most matter is in galaxies and one can put $\nu=2$ as a good approximation. Use of the intergalactic Hubble curve is then, in principle, no more involved than use of the standard FLRW Hubble curve. The division of $\rho_0$ into galactic and intergalactic parts for gravitational-optics purposes seems simplistic but it is certainly less simplistic than ignoring optical effects of inhomogeneities altogether as is done by using FLRW. | 98 | 4 | astro-ph9804249_arXiv.txt |
9804 | astro-ph9804280_arXiv.txt | The history of the transition from a neutral intergalactic medium to one that is almost fully ionized can reveal the character of cosmological ionizing sources. In this talk I will discuss the implications for rival reionization scenarios of the rapid decline observed in the space density of quasars and star-forming galaxies at redshifts $z\gta 3$. The hydrogen component in a highly inhomogeneous universe is completely reionized when the number of ionizing photons emitted in one recombination time equals the mean number of hydrogen atoms. At $z\sim 5$, the local character of the UV metagalactic flux allows one to define a {\it critical} emission rate of hydrogen-ionizing photons per unit comoving volume, ${\dot{\cal N}}_{\rm ion}=10^{51.5\pm 0.3}\ndotunits$. Models based on photoionization by bright QSOs and/or young galaxies with star formation rates in excess of $0.3-1\sfr$ appear to fail to provide the required number of hydrogen-ionizing photons at these redshifts by large factors. If stellar sources are responsible for keeping the universe ionized at $z\approx 5$, the rate of star formation per unit comoving volume at this epoch must be comparable or greater than observed at $z\approx 3$. | The existence of a filamentary, low-density intergalactic medium (IGM) which contains the bulk of the hydrogen and helium in the universe is predicted as a product of primordial nucleosynthesis and of hierarchical models of gravitational instability with ``cold dark matter'' (CDM) (Cen \etal 1994; Zhang \etal 1995; Hernquist \etal 1996). The application of the Gunn-Peterson constraint on the amount of smoothly distributed neutral material along the line of sight to distant objects requires the hydrogen component of the diffuse IGM to have been highly ionized by $z\approx 5$ (Schneider \etal 1991), and the helium component by $z\approx 2.5$ (Davidsen \etal 1996). It thus appears that substantial sources of ultraviolet photons were present at $z\gta 5$, perhaps low-luminosity quasars or a first generation of stars in virialized dark matter halos with $T_{\rm vir}\gta 10^4\,$K (Couchman \& Rees 1986; Ostriker \& Gnedin 1996; Haiman \& Loeb 1997; Miralda-Escud\`e \& Rees 1997). Early star formation provides a possible explanation for the widespread existence of heavy elements in the IGM (Cowie \etal 1995), while reionization by QSOs may produce a detectable signal in the radio extragalactic background at meter wavelengths (Madau \etal 1997). Establishing the character of cosmological ionizing sources is an efficient way to constrain competing models for structure formation in the universe, and to study the collapse and cooling of small mass objects at early epochs. While the nature, spectrum, and intensity of the background UV flux which is responsible for maintaining the intergalactic gas and the Ly$\alpha$ clouds in a highly ionized state at $z\lta 3$ has been the subject of much debate in the last decade, it is only in the past few years that new observations have provided reliable information on the presence and physical properties of the sources and sinks (due to continuum opacities) of UV radiation in the interval $3\lta z\lta 5$. In this talk I will focus on the candidate sources of photoionization at early times and on the time-dependent reionization problem, i.e. on the history of the transition from a neutral IGM to one that is almost fully ionized. The starting point of this study can be found perhaps in the simple realization that the {\it breakthrough epoch} (when all radiation sources can see each other in the Lyman continuum) occurs much later in the universe than the {\it overlap epoch} (when individual ionized zones become simply connected and every point in space is exposed to ionizing radiation), and that at high redshifts the ionization equilibrium is actually determined by the {\it instantaneous} UV production rate. In the following I will adopt an Einstein-de Sitter universe ($q_0=0.5$) with $H_0= 50h_{50}\,\kmsmpc$. | 98 | 4 | astro-ph9804280_arXiv.txt |
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9804 | astro-ph9804233_arXiv.txt | \noindent Nonlinear time-dependent calculations are being carried out in order to study the evolution of vertically-integrated models of non-selfgravitating, transonic accretion discs around black holes. In this paper we present results from a new calculation for a high-$\alpha$ model similar to one studied previously by Honma, Matsumoto \& Kato who found evidence for limit-cycle behaviour connected with thermal instability. Our results are in substantial agreement with theirs but, in our calculation, the disc material does not always remain completely optically thick and we include a suitable treatment for this. We followed the evolution for several cycles and determined the period of the cycle as being about 780 seconds. Advective cooling is dominant in the region just behind the outward-moving peak of surface density. The behaviour of this model is significantly different from what we saw earlier for low-$\alpha$ models (which we discussed in a previous paper) and we contrast and compare the two situations. | 98 | 4 | astro-ph9804233_arXiv.txt |
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9804 | astro-ph9804005_arXiv.txt | We have obtained intermediate resolution spectra of eleven candidate brown dwarf members of the Pleiades open cluster using the Keck II telescope and LRIS spectrograph. Our primary goal was to determine the location of the ``lithium depletion edge" in the Pleiades and hence to derive a precise age for the cluster. All but one of our 11 program objects have radial velocities appropriate for Pleiades members, have moderately strong \Ha\ emission, and have spectral types M6 to M8.5 as expected from their (R-I)$_c$\ colors. We have constructed a color-magnitude diagram for the faint end of the Pleiades main sequence, including only stars for which high S/N spectra in the region of the lithium $\lambda$6708$\AA$\ absorption line have been obtained. These data allow us to accurately determine the Pleiades single-star lithium depletion edge at I$_{c0}$\ = 17.80, (R-I)$_{c0}$\ = 2.20, spectral type = M6.5. By reference to theoretical evolutionary models, this converts fairly directly into an age for the Pleiades of $\tau$\ = 125 Myr. This is significantly older than the age that is normally quoted, but does agree with some other recent estimates. | The Pleiades was recognized in the 1980's as the best open cluster to attempt to identify brown dwarfs (Stauffer \etal 1989; Jameson \& Skillen 1989) because of its fortuitous combination of proximity, youth and richness. A number of brown dwarf candidates were identified in those and subsequent papers, usually on the basis of photometry obtained from deep imaging surveys in two or more colors. A means to establish that a candidate brown dwarf was at least near, if not necessarily below, the hydrogen burning mass limit was proposed by Rebolo, Mart\'{\i}n and Magazz\`u (1992; hereafter RMM). Below about 0.065 \MSUN, brown dwarfs should never develop core temperatures sufficient for lithium ignition, and thus for objects less massive than this we should find a lithium abundance the same as for the interstellar medium from which the object formed, independent of the object's age. For slightly higher masses, lithium acts as an age scale because the length of time it takes for the core to reach 2.5x10$^6$\ K is a sensitive function of mass. Because these stars are fully convective, once the core temperature exceeds the necessary limit, the entire lithium content of the star should be exhausted rapidly and thus be reflected in an observable change in the photospheric lithium abundance. The current generation of theoretical models make specific predictions about the time evolution of this lithium depletion boundary. For example, D'Antona \& Mazzitelli (1997) predict that at ages 30, 70 and 140 Myr, the lithium depletion edge should occur at 0.17 \MSUN, 0.09 \MSUN\ and 0.07 \MSUN, respectively. Other recent models by Baraffe \etal (1998) and Burrows \etal (1997) make nearly identical predictions of the variation of this lithium depletion boundary with age. Indeed, Bildsten \etal (1997) and others have argued that the age for an open cluster derived in this manner should be better than by any other method. The initial attempts to detect lithium in very low mass open cluster members were not successful (RMM; Marcy, Basri \& Graham 1994; hereafter MBG94). However, Basri, Marcy \& Graham (1996; hereafter BMG96) eventually detected lithium in PPL 15, a Pleiades brown dwarf candidate originally identified by Stauffer, Hamilton \& Probst (1994). Rebolo et al. (1996) later showed that two other Pleiades members about 1 magnitude fainter than PPL 15 also have strong lithium absorption. BMG96 derived an age for the Pleiades based on the presence of lithium in PPL 15 but the absence of lithium in another only slightly brighter Pleiades member (HHJ3). However, Basri \& Mart\'{\i}n (1998) have recently provided evidence that suggests that PPL 15 is a nearly equal-mass spectroscopic binary and thus its individual components would be $\sim$0.75 mag fainter. This allows the location of the lithium boundary to be considerably fainter than had been assumed. The published data therefore no longer constrain the age of the Pleiades nearly as well as one would like because too few stars have been measured spectroscopically in the magnitude range of interest. In this paper, we report on the results of a program to determine lithium abundances for a number of candidate Pleiades members with apparent magnitudes chosen to bracket the possible magnitude range within which the lithium depletion boundary might be located. | The goal of this project was to precisely determine the location of the lithium depletion edge in the Pleiades, and hence to determine an accurate age for the cluster. Figure 2 shows a color-magnitude diagram for very low mass Pleiades members for which lithium data are available. The lithium data are from Oppenheimer \etal (1997 = Opp97), MBG94, BMG96, Rebolo \etal (1996 = Reb96) and from this paper. The (R-I)$_c$\ colors for the Opp97 stars are from Stauffer \etal (1995), where we have converted these spectroscopic (V-I)$_c$\ colors to (R-I)$_c$\ via a relation derived from the Gliese catalog M dwarfs in Leggett (1992). The (R-I)$_c$\ colors for some stars in Reb96 and BMG96 are from PC2 indices provided by Mart\'{\i}n \etal (1996) and our calibration of PC2 vs. (R-I)$_c$. Two points for CFHT-PL-15 are shown: (R-I)$_c$\ = 2.24 as derived from its PC2 index, and (R-I)$_c$ = 2.41 as derived from its VO index. Using either color, CFHT-PL-15 lies below the main sequence defined by the other stars, possibly indicating that it is a non-member. Because it has detected lithium and a radial velocity compatible with Pleiades membership, we prefer to believe it is a member and that the inferred color is anomalous (either intrinsically or due to measurement error). Our primary conclusions are unaffected by how we interpret this star. Finally, the dashed line in Figure 2 is a field star ZAMS derived from the M dwarf photometry provided by Leggett (1992). The location of the stars in Figure 2 show a well-defined correlation with the measured lithium equivalent widths. Bluer than (R-I)$_c$\ = 2.2, none of the stars have detected lithium while redder than that color all of the measured stars have lithium. Similarly, fainter than I$_{c0}$\ = 17.8 all the stars have detected lithium, whereas brighter than that limit all but one of the stars have no lithium. Finally, there is a trend for lithium equivalent width to increase going to lower inferred mass (i.e. fainter and redder). We further believe that the dispersion in I magnitude at a given color is to a large extent real and is primarily an indication that some of the observed stars are photometric binaries. In particular, we suggest that HHJ6 (I$_{c0}$\ = 16.93, R-I$_c$\ = 2.18), Lick-PL1, PPL1, and CFHT-PL-12 are good candidates to be photometric binaries. CFHT-PL-12 also has considerably stronger \Ha\ emission than the other Pleiades stars observed - possibly indicating it is a short-period binary or that we have observed it during a flare. The attribution of PPL1 as a nearly equal mass binary would explain why it has strong lithium absorption despite an I$_c$\ magnitude equal to or brighter than two other cluster members with no detected lithium. An alternative explanation would be that there is a significant age spread in the Pleiades and these four over-luminous stars are the youngest in our sample. By comparison to theoretical models (e.g. D'Antona \& Mazzitelli 1997), their displacement about 0.5 mag above other stars of the same color would require them to be more than 50 Myr younger than the other stars, which we believe is unlikely (see, for example, discussions of this issue by Soderblom \etal 1993 and Steele \& Jameson 1995). Based on the above interpretation of Figure 2, we determine that the single star lithium depletion edge in the Pleiades is at I$_{c0}$\ = 17.8 $\pm$ 0.1 or (R-I)$_{c0}$\ = 2.20 $\pm$ 0.05. The uncertainty estimates are not rigorous, and arise mostly from the uncertainties in the absolute calibration of the photometry in Bouvier \etal (1998). For theoretical evolutionary models which incorporate realistic model atmospheres as their outer boundary condition, and hence which can predict observational colors and magnitudes for brown dwarfs, it is possible to convert directly the empirical lithium depletion edge to an age estimate for the Pleaides. In Figure 3, we plot the absolute I$_c$\ magnitude of the lithium depletion boundary (defined here as the point where lithium has been depleted by a factor of 100) for the most recent models of Baraffe \etal (1998) as a function of age. To place our empirically measured point into this diagram, we assume that the cluster has (m-M)$_o$\ = 5.60 (r $\sim$ 130 pc), and A$_I$\ = 0.06 (c.f. Pinsonneault \etal 1998), leading to M(I$_c$) = 12.2 $\pm$\ 0.15 for the lithium depletion boundary, where we have assumed plausible but again unrigorous 1$\sigma$\ uncertainties of 0.1 mag for the distance modulus and 0.03 mag for the extinction and that the uncertainties add in quadrature. The age derived in this way is then 125 $\pm$\ 8 Myr, where the uncertainty estimate only comes from propagating the 0.15 mag uncertainty of the boundary through the model shown in Figure 3. In order to try to assess the model dependence of this estimate, we have also made a similar calculation for theoretical evolutionary models by Burrows \etal (1997) and D'Antona \& Mazzitelli (1997). Those models do not provide R and I magnitudes, so instead we have used the (R-I)$_c$ color to estimate an I$_c$-band bolometric correction (we adopted the Monet \etal (1992) BC$_I$ vs. (V-I)$_c$\ relation, and a conversion from (V-I)$_c$\ to (R-I)$_c$\ based on data in Leggett 1992). In that manner, we estimate that the lithium depletion boundary in the Pleiades is at M(Bol) = 11.99. Comparing this number to the predictions of the two theoretical models, we get an age for the Pleiades of 130 Myr based on the D'Antona \& Mazzitelli (1997) calculations and 125 Myr for the Burrows \etal (1997) models. The most commonly quoted age for the Pleiades is of order 70-80 Myr (Patenaude 1978; Mermilliod 1981). However, models with a relatively large amount of convective core overshoot can yield much larger ages, as was originally shown by the models of the Padova group (c.f. Mazzei \& Pigatto 1989, who derived an age for the Pleiades of 150 Myr). Other recent models give ages intermediate between these values (e.g. 100 Myr for Meynet, Mermilliod \& Maeder 1993 and $\geq$\ 120 Myr for Ventura \etal 1998). Given the disagreement over the age derived from the upper main sequence turn-off, it is particularly useful to have an independent means to derive the age from low mass stars. Using the lithium detection in PPL15 and the non-detection of lithium in HHJ3, BMG96 estimated the age of the Pleiades to be 115 Myr. Our new result pushes this age even slightly older, but more importantly does so using many more stars and thus provides a much better defined age from the lithium data. Finally, we note that we have chosen to use the ``traditional" distance scale to the Pleiades (r $\sim$\ 130 pc). That distance conflicts with the new Hipparcos distance to the Pleiades of about 116 pc (van Leeuwen \& Ruiz 1997; Mermilliod \etal 1997). We have done this because we believe that the Hipparcos distance for the Pleiades is not correct, as has been discussed in Pinsonneault \etal (1998) and Soderblom \etal (1998). The zeroth order effect of simply using the Hipparcos distance and keeping everything else the same would lead to an even older Pleiades age of about 140 Myr. | 98 | 4 | astro-ph9804005_arXiv.txt |
9804 | astro-ph9804097_arXiv.txt | Multiwavelength observations of high energy flare in 1996 from 3C 279 seems to favour the so called mirror model between different inverse Compton scattering models proposed as a possible explanation of gamma-ray emission in blazars. We performed kinematic analysis of the relativistic blob - mirror system and found that only part of the mirror located very close to the jet axis (very likely inside the jet cone) can re-emit soft photons which serve as a target for production of $\gamma$-rays by relativistic electrons in the blob. Since the presence of well localized scattering mirror inside the jet is problematic, this makes problems for the mirror model. The time scale and the shape of the $\gamma$-ray flare should reflect, in terms of the mirror model, the blob dimensions and the longitudinal distribution of relativistic electrons inside the blob. For the $\gamma$-ray light curve of the type observed in 1996 from 3C 279, i.e. the rising time of the flare during a few days with a sharp cut-off towards the end of the flare, the density of electrons inside the blob should increase exponentially starting from the front of the blob and reach maximum towards the end of the blob. Such distribution of electrons is difficult to explain in a model of a relativistic shock moving along the jet, which would rather inject electrons more efficiently at the front of the blob with a trail of particles on its downstream side. | About 50 blazars have been detected by the Compton Gamma Ray Observatory in the MeV - GeV energy range (Fichtel et al. 1994, von Montigny et al. 1995, Thompson et al. 1995, Mukherjee et al. 1997), and 3 blazars, of the BL Lac type, are discovered in the TeV $\gamma$-rays by the Whipple Observatory (Punch et al. 1992, Quinn et al. 1996, Catanese et al. 1997). These blazars can reach very high $\gamma$-ray luminosities which are variable on time scales as short as a part of a day, in the case of optically violent variable quasars, or even several minutes, in the case of BL Lacs. These observations strongly suggest that $\gamma$-ray emission from blazars is collimated towards the observer within a small angle as a result of relativistic motion of plasma in the jet or directional acceleration of particles. High energy processes occurring in blazars are popularly explained in terms of the inverse Compton scattering (ICS) model in which $\gamma$-rays are produced in ICS of soft photons by electrons in a blob moving relativistically along the jet. Different modifications of this general model mainly concern the origin of soft photons, i.e. whether they come internally from the blob in the jet (synchrotron self-Compton (SSC) model, e.g. Maraschi et al. 1992, Bloom \& Marscher 1993), directly from the disk (e.g. Dermer et al. 1992, Bednarek et al. 1996a,b), are produced in the disk but reprocessed by the matter surrounding the disk (external comptonization (EC) model, e.g. Sikora et al. 1994, Blandford \& Levinson 1995), or produced in the jet but reprocessed by the matter surrounding the jet (the so-called mirror model, Ghisellini \& Madau~1996, henceforth GM). In this last paper it is mentioned that SSC model and external comptonization of photons produced by the broad line region clouds (BLR) illuminated by the disk (EC model) may also contribute to the $\gamma$-ray emission producing a first $\gamma$-ray pre-flare. For the SSC model the amplitude of the $\gamma$-ray variation is expected to be proportional to the square of the variation observed in IR-optical-UV energy range. For the EC model the $\gamma$-ray emission should vary linearly with the low energy synchrotron emission. Such behaviour is not observed in the case of the 1996 flare from 3C 279 in which the $\gamma$-ray variation is more than the square of the synchrotron variation. Moreover, in the $\gamma$-ray light curve of this flare (see Fig.~1 in Wehrle et al. 1997), there is no clear evidence for a double peak structure which could eventually correspond to the first $\gamma$-ray flare produced in terms of SSC or EC models and the second $\gamma$-ray flare produced in terms of the mirror model. Therefore, although the SSC model can not be completely rule out, Wehrle et al. (1997) concludes that the mirror model is favourite by the multiwavelength observations of a strong flare in February 1996 from 3C 279 since it predicts $\gamma$-ray flare with observed features. In this paper we test the mirror model by comparing predictions of the kinematic analysis with the observational results. The possible contributions from SSC and EC models to the $\gamma$-ray production during this flare are neglected since, as we mentioned above, there is no observational support for their importance. Simultaneous analysis of all these models will require an introduction of additional free parameters (density of electrons in the blob, the perpendicular extend of the blob, definition of the disk radiation) which are not all well constrained by the observations. | We discuss details of the mirror model proposed by Ghisellini \& Madau. This model seems to be favourite by the multiwavelength observations of the $\gamma$-ray flare in 1996 from 3C 279 (Wehrle et al. 1997). Based on the analysis of the kinematics of the emission region (a blob moving relativistically along the jet) we come to the conclusion that only relatively small part of the mirror is able to re-emit soft photons which serve as a target for production of $\gamma$-rays. For the parameters of the $\gamma$-ray flare observed in 1996 from 3C 279, the radius of this part of the mirror should be comparable to the longitudinal extend of the blob. It has to be of the order of $2 \times 10^{16}$ cm in order to be consistent with the rising time of the flare. This part of the mirror should lay inside the jet cone provided that its opening angle is of the order of $\sim 1/\gamma$. As mentioned in Ghisellini \& Madau (GM), the physical processes in the jet may prevent the presence of the well localized mirror inside the jet. The calculations of density of photons re-emitted by the mirror are done by Ghisellini \& Madau (see Fig.~2 in GM) in a time independent picture which do not take into account the dynamics of the blob. As a consequence they integrate over the parts of the mirror at distances from the jet axis which are much larger than the maximum distance $h_{\rm u}$ (Eq.~(\ref{eq17})), found in our dynamical (time dependent) analysis. The photon densities seen by the blob can not be directly compared with that ones obtained by us in a time dependent version of the mirror model. Ghisellini \& Madau results are only correct for the continuous (time independent) flow of relativistic plasma along the jet axis but overestimates the density of soft photons seen by the relativistic electrons in the blob with limited longitudinal extend. The relativistic blobs in blazars has to be confined to the part of the jet in order to produce the $\gamma$-ray flares with the observed rising time scale. We computed the $\gamma$-ray light curves expected in the dynamical version of the mirror model for different distribution of relativistic electrons inside the blob and assuming that the density of electrons in the blob changes during propagation along the jet. Slowly rising $\gamma$-ray flux with sudden cut-off towards the end of the flare, as observed in 3C 279, is obtained in the case of inhomogeneous blob with electron densities exponentially rising towards the end of the blob. Such electron distribution is difficult to understand in the popular scenario for $\gamma$-ray production in which relativistic shock moves along the jet. It seems that such shock should rather inject relativistic electrons with high efficiencies close to the front of the blob, with the trail of electrons on its downstream side (Kirk, Rieger \& Mastichiadis~1998). However the $\gamma$-ray light curve expected in this case is different than observed during the flares in the blazar 3C 279. Since $\gamma$-rays are produced in a region which is close to the mirror, therefore the shape of the $\gamma$-ray light curve is not very sensitive on the variations of the density of electrons during the time of propagation of the blob between the base of the jet and the mirror. Of course the absolute $\gamma$-ray fluxes produced by the blobs with different evolutions of electron densities in time may differ significantly. The $\gamma$-ray light curves presented in Figs.~\ref{fig2} show very sharp cut-offs towards the end of the flare due to our assumption on the negligible thickness of the mirror. In fact, the observed width of the peak in the $\gamma$-ray light curve of 3C 279, of the order of $t_m\sim 1$ day (see Fig.~1 in Wehrle et al. 1997), may be related to the time in which relativistic blob is moving though the mirror with the finite thickness. If this interpretation is correct then the thickness of the mirror has to be limited to $\rho_{\rm m}\approx c t_{\rm m} \beta (1+\beta)\gamma^2 \approx 4\times 10^{17}$ cm which is comparable to the distance of the mirror from the base of the jet. In this analysis we do not consider production of $\gamma$-rays in terms of the SSC and EC models simultaneously with the mirror model since there is no clear evidence of their importance in the $\gamma$-ray light curve and the multiwavelength spectrum observed in 1996 from 3C 279 (Wehrle et al. 1997). The $\gamma$-ray light curves reported in Figs.~2 show only relative change of the $\gamma$-ray flux with time. They are not straightforwardly dependent on the parameters of the blob (the magnetic field, electron density, blob perpendicular extend, disk radiation) which are not uniquely constrained by the observations. The SSC and EC models will require to fix these parameters in order to guarantee reliable comparisons. | 98 | 4 | astro-ph9804097_arXiv.txt |
9804 | astro-ph9804268_arXiv.txt | Observations of galactic black hole candidates made by the instruments aboard the Compton GRO in the hard X-ray and $\gamma$-ray bands have significantly enhanced our knowledge of the emission properties of these objects. Understanding these observations presents a formidable challenge to theoretical models of the accretion flow onto the compact object and of the physical mechanisms that generate high-energy radiation. Here we summarize the current state of observations and theoretical interpretation of the emission from black hole candidates above 20 keV. The all-sky monitoring capability of BATSE allows, for the first time, nearly continuous studies of the high-energy emission from more than a dozen black hole candidates. These long-term datasets are particularly well-suited to multiwavelength comparison studies, from the radio upward in frequency (Zhang et al. 1997a, these proceedings). Energy spectral evolution and/or spectral state transitions have been observed from many of the black hole candidates. Moderately deep searches of the galactic plane suggest a deficit of weak $\gamma$-ray transients. Such population studies have implications for the origin of black hole binaries and the nature of accretion events. Observations above 50 keV from OSSE demonstrate that in the $\gamma$-ray band there exist two spectral states that appear to be the extensions of the X-ray low (hard) and high (soft), or perhaps very high, states. The former state, the ``breaking'' state, cuts off with e-folding energy $\sim$100 keV and has its peak luminosity near this energy; thus substantial corrections need to be made to historical estimates of the bolometric luminosity of black holes in the ``low'' state. In contrast, in the X-ray high (soft) state, the luminosity peaks in the soft X-rays and the spectrum extends with an unbroken power law, even up to energies above 500 keV in some cases. COMPTEL has detected emission above 750 keV from Cyg X-1 and the transient GRO~J0422+32. In both cases the data suggest that an additional weak, hard spectral component is required beyond that observed by OSSE at lower energies, although the precise spectral form is yet to be determined. The breaking $\gamma$-ray spectrum can be well modeled by Comptonization of soft photons from the accretion disk in a hot thermal plasma. However, recent studies of the combined X-ray and $\gamma$-ray spectrum of Cyg~X-1 and GX339--4 cast severe doubts on the simple geometry of a hot corona overlying a thermal accretion disk. Furthermore, timing studies of the former source are inconsistent with spectral formation by Compton scattering in a uniform, compact hot cloud, suggesting instead a decline in electron density with increasing radius. The power-law $\gamma$-ray spectral state creates more significant theoretical challenges, particularly in explaining the lack of a break at energies exceeding the electron rest mass. It has been suggested that in the X-ray high (soft) state, the high-energy emission arises from bulk-motion Comptonization in the convergent accretion flow from the inner edge of the accretion disk. Such a process can conceivably generate the $\gamma$ ray spectrum extending without a cutoff, if the accretion rate approaches that of Eddington. | The most reliable evidence for the presence of a black hole in a binary system comes from determination of a mass function through optical measurements of the radial velocity of the companion star. If the resulting lower limit on the mass of the compact object exceeds 3$\Msun$, the upper limit for the mass of a stable neutron star based on current theory, then one can reasonably assume that the compact object is a black hole. There are at least nine X-ray binary systems with minimum mass estimates exceeding 3$\Msun$, of which three (Cyg~X-1, GRO~J0422+32, and GRO~J1655--40) have been clearly detected by GRO instruments. Other objects are identified as BHCs based on the similarity of their high-energy spectra and rapid time variability to those of Cyg~X-1. Such classification is, of course, somewhat tenuous. Before neutron stars and black holes can be reliably distinguished based on their X-ray and $\gamma$-ray spectra, the full range of spectral forms from both classes must be observed and characterized. Extensive knowledge of the X-ray emission of these objects has accumulated in the literature, but the broad nature of the $\gamma$-ray emission is only now coming to light, with the high sensitivity of current-generation instruments. The instruments of the Compton GRO have made extensive observations in the hard X-ray and $\gamma$-ray bands of galactic black hole candidates (BHCs). With its all-sky capability, BATSE has monitored emission on a nearly continuous basis from at least three persistent sources (Cyg~X-1, 1E1740.7--2942, GRS~1758--258) and eight transients (GRO~J0422+32, GX339--4, N Mus 1991, GRS~1716--249, GRS~1009--45, 4U~1543--47, GRO~J1655--40, and GRS~1915+105). Lightcurves are presented below in Fig. \ref{lightcurve}. OSSE has made higher-sensitivity, pointed observations of all of these sources, spectra of which appear below in Fig. \ref{two_state}. COMPTEL has detected emission above 750 keV from Cyg~X-1 and GRO~J0422+32. To date, there have been no reported detections of galactic BHCs by EGRET. The French coded-aperture telescope Sigma on the Russian Granat spacecraft has imaged at least a dozen BHCs, including most of those in the list above, but with the addition of TrA~X-1, GRS~1730--312, and GRS~1739--278. The latter two objects were weak transients discovered during a multi-year survey of the galactic center region and have been classified as BHCs by their outburst lightcurves and the hardness of their spectra (Vargas et al. 1997). In this survey, Sigma regularly detected the persistent, variable sources 1E1740.7--2942 and GRS~1758--258, both of which are classified as BHCs on spectral grounds. The most striking result from Sigma observations of BHCs is the detection of broad spectral features below 500 keV from 1E1740.7--2942 (Sunyaev et al. 1991, Bouchet et al. 1991, Churazov et al. 1993, Cordier et al. 1993) and N Mus 1991 (Goldwurm et al. 1992, Sunyaev et al. 1992). These features have been interpreted as thermally broadened and red-shifted annihilation radiation from the vicinity of the compact object. | 98 | 4 | astro-ph9804268_arXiv.txt |
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9804 | astro-ph9804118_arXiv.txt | We present results of \ASCA\ deep exposure observations of the hardest X-ray source discovered in the \ASCA\ Large Sky Survey (LSS) project, designated as AX~J131501$+$3141. We extract its accurate X-ray spectrum, taking account of the contamination from a nearby soft source (AX~J131502$+$3142), separated only by 1$'$. AX~J131501$+$3141 exhibits a large absorption of $N_{\rm H} = (6^{+4}_{-2})\times 10^{22}$ \NHUNIT\ with a photon index $\Gamma = 1.5^{+0.7}_{-0.6}$. The 2--10 keV flux was about $5\times 10^{-13}$ \FLUXUNIT\ and was time variable by a factor of 30\% in 0.5 year. From the highly absorbed X-ray spectrum and the time variability, as well as the results of the optical follow-up observations (\cite{Akiyama98}), we conclude that AX~J131501$+$3141 is a type 2 Seyfert galaxy. Discovery of such a low flux and highly absorbed X-ray source could have a significant impact on the origin of the cosmic X-ray background. | Since the discovery of the Cosmic X-Ray Background (CXB) more than 30 years ago (\cite{Giacconi62}), its origin has been a long standing puzzle. With the \ROSAT, $\sim$70\% of the CXB below 2 keV has been resolved into point sources, more than 60\% of which are type 1 active galactic nuclei (AGNs) (\cite{Vikhlinin95}; \cite{Hasinger96}; \cite{McHardy98}; \cite{Hasinger98}). Origin of the CXB above the 2 keV band, however, is less clear due to the absence of the imaging instrument in this energy band. One problem in the hard X-ray band, often referred as the spectral paradox (e.g. \cite{Fabian92}), is that the X-ray spectrum of the CXB in the 2--10 keV band is harder than that of the typical type 1 AGN, which is presumably the main contributor to the CXB below 2 keV. The 2--10 keV X-ray spectra of type 1 Seyfert galaxies (most of the bright AGNs) are approximated by a power-law with a mean photon index of 1.7 (\cite{Mushotzky93}), which is significantly steeper than that of the 2--10 keV CXB of about 1.4 (\cite{Gendreau95}). This fact implies that the origin of the CXB above 2 keV differs, at least in part, from that below 2 keV. In addition, 20\% of the total energy of the CXB is contained in the 2--10 keV band, whereas only a few percent of the CXB is contained below 2 keV (see review by Fabian \& Barcons (1992), Hasinger (1996)). Thus, the 2--10 keV band would be an essential energy range to solve the origin of the CXB. \ASCA\ is the first satellite with the capability of the hard X-ray (up to 10keV) imaging and spectroscopy, hence is presently the best satellite to investigate the CXB in the 2--10 keV band. In the \ASCA\ Large Sky Survey project (LSS: \cite{Inoue96}; \cite{Ueda98a}), a continuous field of 7 deg$^2$ near the north galactic pole was surveyed with a sensitivity higher than any previous surveys in this energy band. Ueda \etal\ (1998a) resolved a significant fraction of the CXB, about 30\% of the CXB, into discrete sources at a sensitivity limit of $F_{\rm X}\sim 10^{-13}$ \FLUXUNIT\ (2--10 keV). The mean photon index in the 2--10 keV band for these resolved sources ($F_{\rm X}=$(1--4)$\times 10^{-13}$ \FLUXUNIT\ in 2--10 keV) was found to be $\Gamma = 1.5\pm 0.2$. This result is consistent with the idea that the photon index approaches to that of the CXB, $\Gamma\sim 1.4$, as the source flux decreases. However, due to limited photon statistics, the spectral information of the resolved sources was too poor to address the nature of individual X-ray sources. The hardest source in the LSS (hereafter we refer it as the ``LSS hardest source'') was found to show a photon index of $\Gamma\sim -0.2$ with no correction of an absorption (\cite{Ueda96}). However, it is unclear whether the apparent hard spectrum is due to a large absorption or due to flatness of the intrinsic spectrum. The LSS hardest source, which was found in an unbiased survey, would provide us a good opportunity to investigate the nature of faint and hard sources which could significantly contribute to the CXB above 2 keV. Hence, we have performed follow-up \ASCA\ and optical observations on the LSS hardest source. This paper reports results of the \ASCA\ deep exposure observations, while those of the optical observations are given by Akiyama \etal\ (1998). | We extracted the accurate spectrum of the LSS hardest source, AX~J131501$+$3141, taking account of the contamination from the nearby soft source, AX~J131502$+$3142, from which no significant X-ray emission was found in the LSS (Ueda 1996). We found that AX~J131501$+$3141 exhibits a large absorption of $N_H = (6^{+4}_{-2})\times 10^{22}$ \NHUNIT\ with a photon index $\Gamma = 1.5^{+0.7}_{-0.6}$. It showed a long-term time variability between two observations separated by 0.5 year. While the photon index of AX~J131501$+$3141 is consistent with the canonical value of type 1 AGNs (e.g., Mushotzky 1993), its absorption column density is larger than that of typical type 1 AGN by more than an order of magnitude, although a part of type 1 AGNs, about 10\% of them (\cite{Schartel97}), shows a column density larger than $5\times 10^{22}$ \NHUNIT. It is important that this source is selected fully unbiasedly. Hence, its X-ray properties should provide a key to understand the general nature of the missing hard X-ray populations which constitute the CXB above 2 keV. Two major possibilities have been proposed to account for the apparent hard spectrum of the CXB: one is to introduce large absorptions of sources (e.g., \cite{Awaki91}), and the other is to consider populations of sources with intrinsically flat spectra (e.g., \cite{Morisawa90}; \cite{Matteo97}). Our results of the LSS hardest source strongly suggest that highly absorbed sources play an important role in considering the origin of the hard X-ray background. The large absorption of $6\times 10^{22}$ \NHUNIT, the photon index of $\Gamma\sim 1.5$, and the time variability are common properties seen in type 2 Seyfert galaxies. In fact, systematic studies of type 2 Seyfert galaxies by Awaki \etal\ (1991) and Ueno (1996) revealed that they commonly show large absorptions of $\sim 10^{23}$ \NHUNIT\ and photon indices of 1.5--1.7. Akiyama \etal\ (1998) found one bright optical galaxy with $B=17.25$ mag near the center of the X-ray error circle of 0.5$'$ radius in the optical follow-up observations. No other optical source with the flux larger than $B=22.4$ mag is found in the error circle. Akiyama \etal\ (1998) performed spectroscopic observations of the bright galaxy and found that ratios of emission lines are similar to those found in type 2 Seyfert galaxies. The redshift of this galaxy was determined to be 0.07. From the redshift, the observed flux in the 2--10 keV band, $5\times 10^{-13}$ \FLUXUNIT, can be converted to the absorption corrected luminosity of $L_{\rm X}\sim 2\times 10^{43}\ {\rm erg~s}^{-1}$. This luminosity is consistent with those of Seyfert galaxies. Thus, we identify AX~J131501$+$3141 found in the unbiased X-ray survey as a type 2 Seyfert galaxy. Using the LogN-LogS relation in Hasinger \etal\ (1998), we estimate that the chance coincidence between AX J131501$+$3141 and AX J131502$+$3142 is $\sim$ 3\%. However, these two sources have probably no physical correlation, because AX J131502$+$3142 is likely to be a QSO\footnote{ In the optical imaging observations of R- and B-band by Akiyama \etal\ (1998), we found two point-like optical sources located at about 10$''$ north from the center of the error circle with 30$''$ radius for the soft source (AX J131502$+$3142). They are very close to each other and their magnitudes are comparable. Total magnitudes of the two sources are B=20.8 mag and R=19.6 mag; B-R color is 1.21. Since both the optical color and the optical to soft X-ray flux is consistent with that for type-1 AGNs, it is possible that one of them is a quasar which is also responsible for AX J131502$+$3142. However there are fainter optical sources (R$\gtrsim$22) in the error circle and the optical identification for AX J131502$+$3142 is not clear yet. } which is more distant than the new type 2 Seyfert AX J131501$+$3141. Awaki (1991), Madau, Ghisellini \& Fabian (1994) and Comastri \etal\ (1995) predicted that the combination of type 1 and type 2 AGNs can reproduce the CXB spectrum, based on the unified AGN scheme (e.g., \cite{Antonucci93}). In the scheme, type 1 and type 2 AGNs are essentially the same objects, observed from different viewing angle. These type 2 AGNs, which exhibit apparently fainter and harder X-ray spectra than those of type 1 AGNs, should become detectable as the detector sensitivity increases. Although we have examined only one sample from the LSS at this moment, the result is encouraging not only for the unified AGN scheme, but also for solving the origin of the CXB. | 98 | 4 | astro-ph9804118_arXiv.txt |
9804 | astro-ph9804162_arXiv.txt | We present observations of NGC 4038/39 in the [\ion{C}{2}] 158 \micron \/ fine structure line taken with the MPE/UCB Far-infrared Imaging Fabry-Perot Interferometer (FIFI) on the KAO. A fully sampled map of the galaxy pair (without the tidal tails) at 55\arcsec \ resolution has been obtained. The [\ion{C}{2}] emission line is detected from the entire galaxy pair and peaks at the interaction zone. The total [\ion{C}{2}] luminosity of the Antennae is $L_{\rm [C II]} = 3.7 \times 10^{8} L_{\sun}$, which is about 1\% of the far-infrared luminosity observed with IRAS. The main part of the [\ion{C}{2}] emission probably arises from photodissociation regions (PDRs), and a minor fraction may be emitted from \ion{H}{2} regions. A small part of the [\ion{C}{2}] emission comes from standard cold neutral medium (CNM); however, for high temperature ($T \sim 100$~K) and high density ($n_{\rm H} \sim 200$~cm$^{-3}$) about one third of the observed [\ion{C}{2}] emission may originate from CNM. From PDR models we derive densities of the order of $\sim 10^{5}$~cm$^{-3}$ and far-UV (FUV) intensities of $460\chi_{\circ}$, $500\chi_{\circ}$, and $240\chi_{\circ}$ for the PDRs in the interaction zone, NGC~4038, and NGC~4039, respectively. However, PDRs with densities of the order of $\sim 10^{2}$~cm$^{-3}$ and FUV intensities of the order of $\sim 100\chi_{\circ}$ could also explain the observed [\ion{C}{2}] emission. The minimum masses in the [\ion{C}{2}] emitting regions in the interaction zone and the nuclei are a few $\times 10^{7}~M_{\odot}$. A comparison with single dish CO observations of the Antennae shows a [\ion{C}{2}] to CO intensity ratio at the interaction zone a factor of 2.6 lower than usually observed in starburst galaxies, but still a factor of about 1.3 to 1.4 higher than that at the nuclei of NGC 4038/39. Therefore, no global starburst is taking place in the Antennae. [\ion{C}{2}] emission arising partly from confined starburst regions and partly from surrounding quiescent clouds could explain the observed [\ion{C}{2}] radiation at the interaction zone and the nuclei, though the star formation activity toward the nuclei is lower. Accordingly there are small confined regions with high star formation activity in the interaction zone and with a lower star formation activity in the nuclei. This supports the high density and high FUV intensity for the PDRs in the interaction zone and the nuclei. | The galaxy pair NGC 4038/39 (Arp 244) is an interacting system in an early stage of merging at a distance of about 21 Mpc from our own galaxy. On long-exposed images in the optical (e.g. Arp 1966) the interaction is clearly visible because of the tails (``Antennae'') emerging from two uniformly luminous, partly overlapping ovals and because of the dwarf galaxy that appears to have formed at the tip of the southern tail through the interaction (Zwicky 1956, Schweizer 1978, Mirabel, Dottori, \& Lutz 1992). The tails contain about 70~\% of the total amount of \ion{H}{1} in the system (van der Hulst 1979). Short-exposed images (Laustsen, Madsen, \& West 1987) reveal hints of the interaction on somewhat smaller scales; the distorted arrangement of H$\alpha$ knots and the velocity distribution of the individual knots lead Rubin \etal \ (1970) to conclude there is an interaction of two rotating galaxies. Computer simulations carried out in the classical paper of Toomre \& Toomre (1972) and later by Barnes (1988) can account for the present morphological appearance of the system quite well by assuming an interaction of two rotating spiral galaxies. Spectra of the nuclei of both galaxies taken in the optical range do not look like pure starburst spectra but consist of a composition of early-type stars and late giants (Keel \etal \ 1985). The detection of bright near-infrared peaks at the nuclei lead Bushouse \& Werner (1990) to the same result. Their images of the Antennae in J- and R-band and in H$\alpha$ show also the same pattern of bright knots in the surroundings of NGC 4038 and in the bridge connecting both galaxies. The Antennae system as a whole shows a relatively low star forming efficiency according to the ratio $L_{\rm IR} /M({\rm H}_2) \approx 8.55$ measured by Young \etal \ (1986); it is only a factor of 3 higher than in the Milky Way. The ratio determined by Sanders \& Mirabel (1985) is almost twice as high. However, they observed a smaller region in CO and therefore probably underestimated the molecular mass. Comparison of the Antennae with the sample of interacting and isolated galaxies of Young \etal \ (1986) shows NGC 4038/39 to have characteristics more like isolated galaxies. Measurements in the radio continuum at 1.5~GHz and 4.9~GHz of NGC~4038/39 (Hummel \& van der Hulst 1986) reveal a number of discrete knots which coincide in general with H$\alpha$ knots, and an underlying diffuse component. This diffuse component has a steep spectral index on average which indicates non-thermal emission, and the peak of the diffuse radio emission is at the dust patch near the overlapping region. The discrete radio knots account for roughly 35~\% of the total radio emission and have a spectral index of $\alpha \approx -0.5$ on average, probably due to a thermal contribution (Hummel \& van der Hulst 1986). Interferometric observations of the Antennae in CO ($1 \to 0$) by Stanford \etal \ (1990) show three main concentrations of CO emission. Two are associated with the nuclei and the third with the interaction zone. The overlap region is the strongest CO source and contains $\approx 10^{9}~M_{\sun}$ of gas, roughly as much H$_{2}$ as both nuclei together. Based on 10$\mu$m and H$\alpha$ data, the authors have calculated a star forming rate of $5 M_{\sun}$~yr$^{-1}$, and consequently the life time of the molecular gas is $2 \times 10^{8}$~yr. Single dish observations in CO were made at the interaction zone and both nuclei of the Antennae by Aalto \etal \ (1995). The ratio of the emission lines of $^{12}$CO and $^{13}$CO measured in the nuclei and the overlapping region of the Antennae is similar to that found in the central regions of ``normal'' starburst galaxies (Aalto \etal \ 1995). An excellent tracer of star formation activity in galaxies is the strong [\ion{C}{2}] 158~$\mu$m $^{2}$P$_{3/2}\to \, ^{2}$P$_{1/2}$ fine structure line which arises mainly from photodissociation regions (PDRs) created by far-ultraviolet photons from hot young stars impinging on nearby dense interstellar clouds (Crawford \etal \ 1985, Stacey \etal \ 1991). In combination with CO and FIR observations, the [\ion{C}{2}] emission can be used with PDR models (Tielens \& Hollenbach 1985, Wolfire, Hollenbach, \& Tielens 1989, Wolfire, Tielens, \& Hollenbach 1990) to derive densities and far-UV intensities and estimates of the star formation activity. Extragalactic surveys of [\ion{C}{2}] emission (Crawford \etal \ 1985, Stacey \etal 1991) concentrated mainly on nuclei while more recent observations have imaged individual galaxies to study the distribution of [\ion{C}{2}] emission on large scales. [\ion{C}{2}] images of M83 (Geis \etal \ 1998) and NGC~6946 (Madden \etal \ 1993) demonstrate that the emission is extended at least over the full optical extent and often follows the distribution of the FIR and CO within the disk of the galaxies. In the case of NGC~6946 [\ion{C}{2}] emission beyond the optical extent of the galaxy has been found. This [\ion{C}{2}] emission has been attributed to diffuse gas and not to PDRs. Therefore we also estimate the possible contribution of [\ion{C}{2}] emission from neutral atomic gas and from ionized gas in NGC~4038/39. Because of its relative proximity, NGC~4038/39 is a unique source for carrying out spatially resolved measurements of the [\ion{C}{2}] line in an interacting system. We present the results of our imaging spectroscopy study of the [\ion{C}{2}] line in NGC~4038/39 and compare them with observations obtained with ISO. | We present a map of NGC 4038/39 in the [\ion{C}{2}] 158 \micron \ fine structure line. [\ion{C}{2}] emission is detected over the optical extent of the system of galaxies and peaks at the interaction zone. The total luminosity of the [\ion{C}{2}] line is $3.7 \times 10^{8} L_{\sun}$ which is about 1\% of the FIR luminosity of the Antennae. Only a negligible fraction of the observed [\ion{C}{2}] emission can originate in the WNM, if conditions are similar to Galactic atomic clouds. Under normal conditions the [\ion{C}{2}] emission from standard CNM makes only a small contribution to the total [\ion{C}{2}] emission, however it may rise to $\slantfrac{1}{3}$ of the total [\ion{C}{2}] emission for individual positions. Only minor fractions of the [\ion{C}{2}] emission at the interaction zone and the nuclei can arise from \ion{H}{2} regions. PDRs are the dominant source for the [\ion{C}{2}] emission. We estimate minimum hydrogen masses associated with the [\ion{C}{2}] emitting region of $1.9 \times 10^{8} M_{\sun}$ for the entire merging system and $6.8 \times 10^{7} M_{\sun}$, $3.2 \times 10^{7} M_{\sun}$, and $3.7 \times 10^{7} M_{\sun}$ within one beam centered at the interaction zone, NGC~4038, and NGC~4039, respectively. Assuming a single emission component in the beam we derive a density of the [\ion{C}{2}] emitting gas in PDRs of $1 \times 10^{5}$ cm$^{-3}$ for the interaction zone and of $2 \times 10^{5}$ cm$^{-3}$ and $1 \times 10^{5}$ cm$^{-3}$ for NGC~4038 and NGC 4039, respectively, and a far-UV intensity of $450 \chi_{\circ}$ for the interaction zone, $500 \chi_{\circ}$ for NGC~4038, and $250 \chi_{\circ}$ for NGC~4039. The derived beam filling factor of the emission from the PDRs is 20\% for the interaction zone and 10--15\% for the nuclei. However, the PDR model also allows a solution with low density ($\sim 10^{2}$~cm$^{-3}$), high beam filling factor ($\sim 50$\%), and low FUV intensity ($\sim 100 \chi_{\circ}$). The low, beam averaged [\ion{C}{2}]/CO ratio of 2350 toward the interaction zone and the even lower ratios at the nuclei indicate that no global starburst is going on either in the area surrounding the interaction zone or in the nuclear region. The high-excitation lines observed with ISO SWS which trace the starburst must therefore arise from a small, confined region in the interaction zone. This result is also supported from the observations with ISOCAM. Therefore on the scale of our [\ion{C}{2}] beam a single emission component for the PDR is only a crude approximation. Using interferometric and single dish CO observations and an expected [\ion{C}{2}]/CO ratio for starburst regions and for quiescent clouds, we constructed a two-component model consisting of a confined starburst region and of molecular clouds enveloping the starburst. From this model we find that the [\ion{C}{2}] emission at the interaction zone originates partly from confined starburst regions and partly from surrounding quiescent clouds. If we apply this model to the nuclei we get also enhanced star formation activity in NGC~4039 with a low [\ion{C}{2}]/CO ratio for the quiescent clouds but only moderate star forming activity in NGC~4038. This two-component model supports the high-density solution for the PDRs. Future investigation of the Antennae in the [\ion{N}{2}] 205 \micron \ fine structure line would be very helpful to further disentangle the origins of the [\ion{C}{2}] line. Also observations at higher spatial resolution in the FIR regime (e.g. with FIRST and SOFIA) would be a great step forward to investigate this and other spatially very complex objects in more detail. | 98 | 4 | astro-ph9804162_arXiv.txt |
9804 | astro-ph9804024_arXiv.txt | Understanding the properties of interstellar turbulence is a great intellectual challenge and the urge to solve this problem is partially motivated by a necessity to explain the star formation mystery. This review deals with a recently suggested inversion technique as applied to atomic hydrogen. This technique allows to determine 3D turbulence statistics through the variations of 21~cm intensity. We claim that a radio interferometer is an ideal tool for such a study as its visibility function is directly related to the statistics of galactic HI. Next, we show how galactic rotation curve can be used to study the turbulence slice by slice and relate the statistics given in galactic coordinates and in the velocity space. The application of the technique to HI data reveals a shallow spectrum of the underlying HI density that is not compatible with a naive Kolmogorov picture. We show that the random density corresponding to the found spectrum tends to form low contrast filaments that are elongated towards the observer. | The properties of the interstellar medium strongly suggest that it is turbulent. Here turbulence is understood as unpredictable temporal behavior of nonlinear systems as preached by J. Scalo (1985, 1987). The importance of turbulence in molecular clouds and its relation to star formation has long been appreciated (Dickman 1985). Recent progress in numerical simulations of molecular cloud dynamics (see Ostriker, this volume) indicates the intrinsic connection between the turbulence in different phases of the interstellar medium (McKee \& Ostriker 1977). In what follows we shall mostly discuss the turbulence in atomic hydrogen (HI), although the formalism presented here is applicable to other spectral lines. Statistical description is a nearly indispensable strategy when dealing with turbulence and a big advantage of statistical techniques is that they extract underlying regularities of the flow and reject incidental details. Kolmogorov notion of energy cascade from large to small scales has been proved an extremely valuable concept and Kolmogorov spectrum of turbulence has been measured in various media. At the same time, astrophysical turbulence, unlike that in incompressible fluids, is a much more complicated phenomenon, and therefore one cannot {\it a priori} hope that Kolmogorov's (1941) description is adequate (cf. Armstrong et al. 1995). Energy injection at small scales, shocks, compressibility may make interstellar turbulence spectrum much more informative, and we should expect to see deviations from the boring $-11/3$ slope. The advantage of using 21~cm emission data is that a continuum of separations between data points is available. This property is shared by diffuse emission in other spectral lines, but 21~cm measurements allow to disregard dust adsorption. As our review deals with HI studies within the galactic plane we do not discuss in detail interesting results obtained for HI in Large Magelanic Clouds (Spicker \& Feitzinger 1988a,b). To avoid possible misunderstanding we should stress that Spicker \& Feitzinger (1988a,b) deal with velocity fluctuations, while only intensity fluctuations are available when one studies HI turbulence in galactic plane. Statistics of random velocity and density fields may be different and therefore a direct comparison of the results obtained for these fields may be misleading. Being limited in space we refer the interested reader to the earlier reviews on interstellar turbulence, among which the one by Dickman (1985) can serve as an excellent introduction to the basic statistical techniques. A more advanced reader will enjoy a thoughtful analysis of problems associated with the statistical analysis of observational data given in Houlahan \& Scalo (1990). Important aspects of the statistical analysis are discussed, for instance, by Dickman \& Kleiner (1985), Roy \& Joncas (1985) Perault et al. (1986), O'Dell \& Casta\~{n}eda (1987), Rickett (1988), Van Langevelde et al. (1992) Kitamura et al. (1993), Meisch \& Bally (1994), Armstrong, Rickett, \& Spangler (1995), Wallin, Watson, \& Wyld (1998) and by the contributors to the present volume. A brief discussion of the very early studies of interstellar statistics can be found in Lazarian (1992). Studies of interstellar turbulence frequently deal with samples which are not statistically homogeneous (see Miesch \& Bally 1994). Indeed, whenever individual molecular complexes are studied, the statistics (especially at large separations) may be dominated by regular gradients rather than the random component. To eliminate the inhomogeneous component, various types of spatial filtering are used (see Zurfleh 1967). These problems are alleviated for HI studies, since molecular clouds tend to be localized objects in sharp contrast to more pervasive distribution of atomic hydrogen. Further on we shall deal with the two point structure functions and power spectra. Naturally, one cannot place pickup devices at different points of HI. Instead, only the 21-cm intensity fluctuations with pointwise emissivity integrated along the lines of sight are available. It is obvious, that given a statistical description of the transparent emitting astrophysical medium, it is possible to predict statistical properties of the observable diffuse emission (Kaplan \& Pikelner 1970), which would correspond to the solution of the {\it forward problem}. However, more important is to solve the {\it inverse problem}, namely, to deduce the 3D statistics of HI from observations. These issues are dealt with in sections 2 and 3. In section 4 we discuss the application of the technique to interferometric data. The galactic rotation curve allows one to study turbulence slice by slice. This slicing, however, is far from trivial (see section 4). Indeed, topologically disconnected blobs of HI can overlap in the velocity space if their velocities are the same. We show that interferometric study can potentially provide the information about both random density and velocity fields. Addressing the issue of HI topology we show that HI with the measured spectrum of density fluctuations forms low contrast filaments (section 5) and these filaments are elongated towards the observer due to the presence of velocity fluctuations. | The following are the principal conclusions of this review: \begin{itemize} \item The application of a newly developed inversion technique reveals a shallow spectrum of HI density in galactic plane. The spectrum is different from the Kolmogorov one, which may be indicative of non-trivial physics involved. \item The spectrum of HI density is anisotropic in velocity space with velocity fluctuations altering the statistics along the line of sight. The degree of anisotropy and the spectrum in velocity space depend on the scale under study. The spectrum become isotropic for scales larger than 200-300~pc. \item Interferometers are useful tools for studying HI turbulence in galactic plane and may provide the spectrum of 3D random density field and useful information on random velocity field. \item Atomic hydrogen with Gaussian density and the shallow spectrum corresponding to observations forms low contrast filaments. These filaments are anisotropic in the velocity space and directed towards an observer. \end{itemize} | 98 | 4 | astro-ph9804024_arXiv.txt |
9804 | astro-ph9804206_arXiv.txt | As part of a continuing study of the effect of cluster environment on the star formation properties of galaxies, we have undertaken an H$\alpha$ objective prism survey of the nearby cluster, Abell 1060. We detect 33 galaxies in emission, 24 of which are cluster members. We present new radial velocity measurements and H$\alpha$+[\ion{N}{ii}] equivalent widths and fluxes for a number of these galaxies. We distinguish between galaxies with diffuse and compact emission, the latter having been associated in previous work with a disturbed morphology of the galaxy and most likely resulting from tidally-induced star formation from galaxy--galaxy or cluster--galaxy interactions. The fraction of cluster spirals in Abell 1060 detected with compact emission agrees with the expected fraction for a cluster of its richness, as derived from results of a previous survey of 8 clusters. Some of the detected cluster early-type spirals exhibit anomalously high global H$\alpha$ equivalent widths, as compared to galaxies of similar type in the field. | \label{resint} \noindent The effect of cluster environment on the star formation properties of galaxies has long been a matter of debate. While some studies have suggested a lower star formation rate for cluster spirals as compared to the field (e.g. \citeNP{gisler,dress85}), other work has suggested a similar or enhanced rate, particularly for early-type spirals(e.g. \citeNP{kenn84,gav91,moss93,enacs}). With the discovery that in distant rich clusters there is a high fraction of blue, star-forming galaxies, often with unusual morphology suggestive of mergers and tidal interactions (e.g. \citeNP{lavhen86,tom88,couch94}), there is renewed interest in tidally-induced star formation by mergers and interactions in nearby clusters. In order to address these issues, Moss, Whittle and co-authors have completed an objective-prism survey of eight nearby clusters of galaxies to detect global \mbox{H$\alpha$+[\ion{N}{ii}]} emission as an indicator of the current rate of massive star formation. The survey technique is described by \citeN{mwi}, hereafter MWI, and initial results have been discussed by \citeN{moss93} (see also \citeNP{moss90,moss95,moss96,moss97}). We have extended this survey to a ninth cluster, the Hydra I cluster, Abell 1060. Abell 1060 is one of the nearest of the Abell clusters, at a redshift of $z\sim$0.01, and is the nearest large cluster beyond the Virgo and Fornax clusters. It has a high spiral fraction (e.g. \citeNP{solanes92}) and is a relatively poor cluster, with a low density intracluster medium \cite{lowmush} and low X-ray luminosity \cite{edge91}. Since it is the nearest of the clusters surveyed by us so far, it can be surveyed to a fainter limit in absolute magnitude. However, its proximity means that it has a large projected diameter on the sky, with one Abell radius, \mbox{$1\, r_{A}=2\fdg3=1.5\, h^{-1}$ \rm{Mpc}} (where $h$ is defined in terms of the Hubble constant $H_{0}=100h\ \rm{km\ s}^{-1}\rm{Mpc}^{-1}$). Whereas other clusters were surveyed in a region of radius 1.5 $r_{A}$, the photographic plate size restricted survey of Abell 1060 to a region of radius somewhat less than one Abell radius (see \S \ref{platty}). \citeN{r89}, hereafter R89, presents a catalogue of 581 galaxies in the cluster area, which contains a sample which is complete to the magnitude limit $V_{25}=16.65$, within 2\degr \/ of the cluster centre. This is a convenient complete sample for the present H$\alpha$ survey, extending to a fainter apparent magnitude than the Zwicky Catalogue used to define samples for other clusters. Cluster properties are summarised in Table \ref{Hydratab}. The (B1950.0) position of the central cluster galaxy NGC 3311 is given as the cluster centre in columns 2 and 3. The mean heliocentric radial velocity $\langle v_{\odot}\rangle$ and velocity dispersion $\sigma$ determined using $N_{gal}$ individual galaxy redshifts are given in columns 4, 5, and 6 \cite{bird94}. The Abell richness class \cite{aco} is given in column 7, the Bautz-Morgan and Rood-Shastry type classes are given in columns 8 and 9 respectively \cite{bm,strood}. \begin{table*} \centering \caption{\label{Hydratab} Cluster properties} \begin{tabular}{ll@{\hspace{1.0ex}}l@{\hspace{1.0ex}}ll@{\hspace{1.0ex}}l@{\hspace{1.0ex}}lcccccc} \\ \hline \\ Name & \multicolumn{6}{c}{R.A.~~(1950.0)~~Dec.} & $\langle v_{\odot}\rangle$ & \(\sigma\) & $N_{gal}$ &Richness&\multicolumn{2}{c}{Type Class} \\ \cline{12-13} &\multicolumn{3}{c}{l}&\multicolumn{3}{c}{b}&km\,s$^{-1}$&km\,s$^{-1}$&&&B-M&R-S\\ \\ (1)&\multicolumn{3}{c}{(2)}&\multicolumn{3}{c}{(3)}&(4)&(5)&(6)&(7)&(8)&(9)\\ \\ \hline \\ Abell 1060 &10$^{{\rm h}\hspace{-0.85ex}}$ &34$^{{\rm m}\hspace{-1.05ex}}$ &21\fs 6&-27$^{\circ\hspace{-0.85ex}}$ &16$^{\prime\hspace{-0.65ex}}$ &05$^{\prime\prime\hspace{-0.85ex}}$&3697 & 630 & 132 & 1 & III & C\\ (Hydra I)&\multicolumn{3}{c}{269\fdg6}&\multicolumn{3}{c}{26\fdg49}\\ \\ \hline \\ \end{tabular} \end{table*} The observations and data reduction are described in \S \ref{obses}. Details of the observational technique are given in \S \ref{platty}, and of the process of identifying the emission-line galaxies in \S \ref{ident}, where a table of the detected emission-line galaxies (ELGs) is given. Measurements of radial velocities for the detected emission-line galaxies are presented in \S \ref{twoplate}, and those of \mbox{H$\alpha$+N[II]} equivalent widths and fluxes in \S \ref{ews}, where measured \mbox{H$\alpha$+N[II]} fluxes are also converted into effective star formation rates. A comparison of detected cluster emission-line galaxies in Abell 1060 with field galaxies and detected emission-line galaxies in other clusters is given in \S \ref{frac}. Notes on individual galaxies are given in \S \ref{indiv}. Finally, we present a brief discussion of our results in \S \ref{discus}. | 98 | 4 | astro-ph9804206_arXiv.txt |
|
9804 | astro-ph9804030_arXiv.txt | We present the first radio observations of a sample of 13 optically and IR bright southern hemisphere classical Be stars made from the Australian Telescope Compact Array at 3.5cm and 6.3cm simultaneously. One star, $\delta$ Cen was detected at 3.5cm and a second, $\mu$ Cen was also thought to have been detected; further observations of this source are required to confirm this detection. No sources were detected at 6.3cm, although $\delta$ Cen was previously detected at this wavelength by other observers at a higher flux than our detection limit. The radio observations show that the spectral energy distribution undergoes a turnover between the far IR and radio wavelengths, as was seen in previous studies. Likewise we find no simple correlation between far IR and radio flux. Lower limits to the outer disc radius were found to be of the order of few hundred solar radii; of the order of those found previously by Taylor et al. | Classical Be stars are defined as non supergiant B stars that have, or have had Balmer lines in emission. They are further characterised by the presence of a continuum excess, arising from free free and bound free emission from a stellar wind. Comparison of optical and UV spectra show that two different wind regimes must coexist. Consequently a high velocity component responsible for the high excitation lines visible in the UV, and a denser component that produces the near IR continuum excess and the optical emission lines was proposed to explain the observations. That the denser component was concentrated in the equatorial plane has long been suspected; recent interferometric data shows that the the envelopes around Be stars are non spherical (Dougherty \& Taylor 1992; Quirrenbach et al 1994; Stee et al 1995). However, only a few of the brightest systems are ammeanable to an interferometric approach. Consequently, other approaches to the study of Be star circumstellar discs have been attempted. One such approach was long wavelength (mm-radio) flux measurement. When combined with near IR photometry modeling of the spectral energy distribution (SED) leads to a profile of the base density and density gradient (and hence a radial velocity law) within the circumstellar disc. However, observations of Be stars at mm and radio wavelengths showed that they were much fainter than implied by a simple extrapolation of their near IR and IRAS fluxes (e.g. Taylor et al 1990; henceforth Ta90), indicating a change in the ion density gradient within the disc. Several explanations were advanced to explain this, the most favourable being a change in disc opening angle or re-acceleration of material at large radii. Other possibilities exist, such as recombination at large radii or a truncation of the disc; see Waters et al. (1991) and Ta90 for a thorough discussion of all the scenarios. Given the paucity of detections in the northern hemisphere (only 6 Be stars have been detected), and the lack of multiple observations it is difficult to quantify the behaviour of the continuum excess at long wavelengths. Because of this shortfall we made observations of a sample of 13 bright southern hemisphere Be stars from the Australian Telescope Compact Array (ATCA) in 1997 April/May in an attempt to increase the sample size. | As a result of observations of a sample of 13 IR bright southern Be stars we have identified $\delta$ Cen as a radio source, which may also be variable. A second star, $\mu$ Cen, is also though to be a radio emitter, although further observations are needed to confirm this. This brings the total number of Be stars detected at radio wavelengths to eight (inclusive of $\mu$ Cen). We confirm earlier results of Ta90 that demonstrate a turnover in the spectral energy distribution between the far IR and radio wavelengths. We find no compelling evidence for a direct correlation between stellar or far IR luminosity and radio flux. Clearly further observations of Be stars at higher sensitivities are required before such a correlation can be confirmed or rejected. Measurement of the minimum outer disc radius reveals that the circumstellar discs extend to several hundred solar radii, again reproducing the results of Ta90. | 98 | 4 | astro-ph9804030_arXiv.txt |
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