T. A. Oosterloo. and M. E. Putman 2

Transcription

1 The Astronomical Journal, 123: , 2002 April # The American Astronomical Society. All rights reserved. Printed in U.S.A. H i FINE STRUCTURE IN MAGELLANIC TIDAL DEBRIS 1 B. P. Wakker Department of Astronomy, University of Wisconsin, 475 North Charter Street, Madison, WI 53706; wakker@astro.wisc.edu T. A. Oosterloo Netherlands Foundation for Research in Astronomy, Postbus 2, NL-7990 AA Dwingeloo, Netherlands; oosterloo@nfra.nl and M. E. Putman 2 Center for Astrophysics and Space Astronomy, University of Colorado at Boulder, Campus Box 389, Boulder, CO 80309; mputman@casa.colorado.edu Received 2001 January 22; accepted 2002 January 3 ABSTRACT We report on H i observations with the Australia Telescope and the Parkes telescope of one of the cores of the high-velocity cloud (HVC) WW 187 (also known as HVC and HVC ). With our H i map, this becomes the first HVC for which single-dish and interferometer data have been combined. The analysis of the H i fine structure shows that the smallest concentrations are almost resolved and may be as small as 5 15 pc. The structure of the cloud is hierarchical: smaller, denser cores are embedded in a smoother envelope. At 1 0 resolution, the peak brightness temperatures of the cores lie in the range K, while the FWHMs of the individual spectra range from 3 to 9 km s 1, with a modal value of 5kms 1. This limits the kinetic temperature of the gas to lie between 35 and 500 K. The peak column densities range from 1: to 3: cm 2. The distribution of column densities in the field is an exponential: count / exp ½ NðH iþš. The volume densities in the cores could be as high as 20 cm 3, which is sufficiently high to provide a long-time source for the molecular hydrogen in the HVC that is seen in the direction of NGC The original motivation for this observation came from the detection of interstellar UV absorption lines in the spectrum of the extragalactic background source NGC 3783 from which metal abundances for WW 187 were derived that are similar to those in the SMC. We describe in detail the derivation of the value for N(H i) used to derive these abundances. In combination with model calculations, these abundances support the contention that WW 187 is part of the leading arm of the Magellanic Stream and that the Magellanic Stream is a tidal feature, rather than being formed by ram pressure. Key words: Galaxy: evolution Galaxy: halo ISM: clouds Magellanic Clouds radio lines: ISM 1. INTRODUCTION High-velocity clouds (HVCs) are conglomerates of neutral hydrogen moving at velocities incompatible with differential galactic rotation (see review by Wakker & van Woerden 1997). These objects seem to have several different origins. Specific examples have now been found for each of these. HVCs with heights above the Galactic plane up to several kiloparsecs could be a manifestation of a Galactic fountain, in which on timescales of hundreds of Myr hot disk gas rises, cools, condenses into neutral H i clouds, and falls back (Bregman 1980, 2001). This model predicts nearsolar metallicity for such HVCs. Other HVCs (including complex C) have low metallicity (0.1 solar; Wakker et al. 1999; Richter et al. 2001; Wakker 2001). Such clouds may represent gas falling in toward the Milky Way for the first time, or even independent Local Group clouds. Blitz et al. (1999) proposed that most HVCs are Local Group clouds. However, as originally presented, this model is not compatible with the H i mass function (Zwaan & Briggs 2000), but a population of some clouds at distances up to 300 kpc from the Milky Way is not excluded. The origin of these clouds could lie in ancient tidal interactions between dwarf 1 Based on observations with the Australia Telescope. 2 Hubble Fellow galaxies and the Milky Way, or they could have formed in the early universe in ways similar to the building blocks that formed the Milky Way. Among the most prominent HVCs is a tidal stream near the Milky Way the Magellanic Stream. Originally interpreted as tidal debris from the Magellanic Clouds (Lin & Lynden-Bell 1977; Davies & Wright 1977; Murai & Fujimoto 1980), discrepancies between the models and the data led some authors to propose that the gas in the Stream was pushed out of the Magellanic Clouds by a drag force exerted by an extended tenuous corona of the Milky Way (Mathewson, Schwarz, & Murray 1977; Heller & Rohlfs 1994; Moore & Davis 1994). However, these models cannot produce a leading stream of debris. In the tidal model the combined tidal force exerted on the Small Magellanic Cloud (SMC) by the Large Magellanic Cloud (LMC) and the Milky Way peaked about 2 Gyr ago and was sufficiently strong to make the gas in the outer parts of the SMC unbound. One full orbit (2 Gyr) later, the decelerated gas has formed a 90 long trailing structure, which is the Magellanic Stream. The gas that was accelerated had to pass close by the LMC, and a large fraction was absorbed by that galaxy. The remainder of the gas had its orbits further disturbed by the LMC and formed an irregular leading arm up to 90 ahead. The clouds classified as population EP (extreme-positive velocity clouds) by Wakker & van Woerden (1991) are the best candidates to be this leading debris, as was already

2 1954 WAKKER, OOSTERLOO, & PUTMAN Vol. 123 Fig. 1. Map of the HVCs in the region l ¼ , b ¼ 35 to 50, with v LSR > 140 km s 1, based on the list of Morras et al. (2000), which covers the sky south of declination 25 on a 0=5 0=5 grid with 1 km s 1 velocity resolution and 0.07 K rms. Colors show the galactic standard of rest (GSR) velocity. The color scale is chosen such that bluish colors correspond to v GSR < 0 (gas moving toward the Milky Way) and greenish to reddish colors correspond to v GSR > 0 (gas moving away from the Milky Way). Contours are for brightness temperatures of 0.05, 0.5, 1, 3, 6, and 9 K. The LMC lies near l ¼ 280, b ¼ 32. The position of NGC 3783 is shown, as are the numbers in the catalog of Wakker & van Woerden (1991) for the most prominent clouds. The thick curved line shows the orbit of the LMC in the model of Gardiner & Noguchi (1996), while the thin curved line shows the orbit of the SMC. The dots show the present-day positions of the test particles in that model. proposed by Mathewson, Cleary, & Murray (1974) in the paper announcing the discovery of the Magellanic Stream. Figure 1 shows the velocity field and distribution of brightness temperature for these clouds based on the high-velocity component list of Morras et al. (2000), with the positions of the Gardiner & Noguchi (1996) tidal model particles overlaid. Population EP is defined as the gas in the region l ¼ , b ¼ 30 to 30 with v LSR > 150 km s 1. West et al. (1985) used the Seyfert galaxy NGC 3783 to detect Ca ii absorption associated with the HVC cataloged as WW 187 by Wakker & van Woerden (1991), which is one of the more prominent cores in population EP. Lu, Savage, & Sembach (1994) used the Goddard High Resolution Spectrograph (GHRS) on the Hubble Space Telescope (HST ) to observe the column density of singly ionized sulphur. Sulphur is a good element for a metallicity measurement as it is found to be undepleted onto dust grains in low-velocity gas (Savage & Sembach 1996) and in the neutral interstellar medium (ISM) its dominant ionization stage is S +. Lu et al. (1994) derived an S ii abundance of 0.15 times solar, using the same value for N(H i) as West et al. (1985). This result was deemed consistent with a tidal origin for WW 187 as well as with it being an independent intergalactic cloud.

3 No. 4, 2002 H i FINE STRUCTURE IN MAGELLANIC TIDAL DEBRIS 1955 However, the column density of sulphur is obtained along the (almost) infinitesimal pencil beam toward the background probe, while the H i column density is based on 21 cm emission data obtained with a large radio beam (34 0 for the Villa Elisa telescope used by West et al. 1985). Wakker & Schwarz (1991) showed that within HVC cores there can be variations in the H i column density of up to a factor of 5 on scales of arcminutes. This implies that without a highresolution H i map the derived abundance may still be uncertain up to a factor of 5. Thus, to obtain a reliable determination of the intrinsic metallicity of cloud WW 187, we obtained time with the Australia Telescope Compact Array (ATCA) to observe the H i in the field around NGC 3783 at 1 0 resolution. This is the first high-resolution H i observation of tidal debris and the first of an HVC with high positive velocities. The ATCA data were combined with an observation made using the Green Bank 140 ft telescope to derive an improved H i column density in the direction of NGC This comes out to be ð8:3 2:0Þ10 19 cm 2, 30% lower than the earlier value, leading to an HVC sulphur abundance of S=H ¼ 0:25 0:07 times solar, very close to that of the SMC (Lu et al. 1998). We also combine the ATCA total column density map with single-dish observations obtained with the Parkes telescope (17 0 beam). Since the latter only has 26 km s 1 velocity resolution, we cannot combine individual channel maps. Nevertheless, this provides an opportunity to study the characteristics of the H i fine structure. Previous studies of high-resolution maps of HVCs (Schwarz & Oort 1981; Wakker & Schwarz 1991; Schwarz, Wakker, & van Woerden 1995; Wakker et al. 1996) were for HVCs with negative velocities relative to the LSR, and for these analyses the interferometer data were not combined with single-dish data. The current data allow a detailed study of the structure of an isolated H i cloud on scales of ,or pc if the cloud is at the distance of kpc predicted by the tidal model. Our understanding of H i small-scale structure is still incomplete. How this varies with the origin and environment of the HVCs remains unclear as a result of the low spatial resolution of many HVC surveys, the low number statistics, and the lack of maps of large-scale structure to complement high-resolution maps. Interaction with different regions of the Galactic halo, as well as the energetics from which the clouds originated, may lead to distinctly different H i properties of the clouds. For instance, Wolfire et al. (1995) predicted that embedding the HVCs in a hot Galactic corona (10 6 K) will result in a two-phase H i structure for clouds within 20 kpc of the Galactic disk, while no cold cores will be found in clouds beyond this radius (e.g., the Magellanic Stream). We measured some of the properties of the H i maps and connect these with the Wolfire et al. (1995) model. The paper is organized as follows. We first discuss the observational parameters and mapping method in detail (x 2). Then we describe the data, first presenting channel maps and position-velocity maps (x 3.1). Next we discuss in detail how the single-dish and interferometer data were combined (x 3.2). A short description of the velocity field (x 3.3) and the continuum sources in the field (x 3.4) are also given. We describe the derivation of the H i column density toward NGC 3783 in x 4. In x 5 we analyze the H i smallscale structure and its implications in the context of an association with the Magellanic System and location in the Galactic halo. 2. OBSERVATIONS 2.1. Observational Setup The important observational parameters for the ATCA observations are listed in Table 1. Two fields were observed, one centered on NGC 3783 and one adjacent field shifted toward the brightest part of the HVC as seen in a single-dish map (see Fig. 2). Observations were done at four different dates, using four different configurations of ATCA. All tracks cover the full 12 hr. The coverage of the UV plane is mostly regular between spacings of 30.6 and m, in intervals of 15.3 m; only 12 of the 50 baselines are missing, most of which lie between 260 and 380 m. The longest used baseline (B max ) corresponds to structure on scales of 0.72/ B max or (41=cos ¼ on the declination axis), while the shortest baseline corresponds to spatial scales of Data for baselines between 3000 and 5000 m are also available (corresponding to scales of ), but there is insufficient signal at these long baselines and so they have been discarded. The correlator was set up to have a starting frequency of GHz and to give 1024 channels of 36 khz (0.824 km s 1 ) width. On-line Doppler tracking was unavailable at ATCA, thus each configuration leads to a slightly different range in velocity space. Typically, velocities range from approximately 150 to 690 km s 1. From single-dish spectra we know that in this direction H i emission only occurs at velocities between 50 and 270 km s 1. Thus, from each configuration 461 channels between 78 and 300 km s 1 were extracted, and the rest were discarded. Discarding the other channels as well as the baselines longer than 3000 m led to a reduction in the size of the data set from 3.0 to 0.94 Gbyte. Calibration was done using the software in the MIRIAD package (Sault, Teuben, & Wright 1995). Either PKS or PKS (see Table 1) was used as the phase calibrator, while the amplitude and bandpass calibration were based on PKS Interferometer Mapping After the calibration step, we had a calibrated UV data set for each of the two fields and four dates. Plotting the calibrated data showed that just 10 UV points needed to be flagged. There is emission in three different velocity intervals: 20 to 10 km s 1, km s 1, and km s 1. The latter corresponds to the HVC. In order to be able to subtract the continuum sources in the field, maps were made for all channels between 190 and 280 km s 1. The velocity scales of the four different UV data sets for each of the two fields were aligned by MIRIAD s mapping program. Because the channels of the individual tracks were not aligned, we chose to rebin the velocity channels to 1 km s 1, rather than use the original km s 1. Uniform weighting was used in order to obtain the maximum possible resolution, giving a synthesized beam with FWHM (R:A: decl:). The sampling was made to be 2.2 pixels beam 1, while the mapped fields cover 2 or 3 times the primary beam FWHM in right ascension or declination, respectively. After the mapping was completed, two more steps had to be taken to create the final maps.

4 1956 WAKKER, OOSTERLOO, & PUTMAN Vol. 123 TABLE 1 Observational Parameters A. Parameter 1994 Jun Mar Jun Jun 10 Integration time (s) Shortest spacing (m) Phase calibrator... PKS PKS PKS PKS Velocity range to to to to Extracted channels Extracted velocities to to to to Parameter Field 1 Field 2 Center (J2000.0) , , (l, b) , , Grid spacing (arcsec) , 26.1 (1.35 s) 16.0, 26.1 (1.35 s) Mapped field (arcmin) Primary beam (arcmin) Shortest spacing (m) (144) 30.6 (144) Longest spacing (m) (3625) (3625) Synthesized beam (arcsec) TB/S (K Jy 1 ) rms (mjy beam 1 ) (1.72 K) 6.61 (1.77 K) Velocity range (km s 1 ) B. Note. Units of right ascension are hours, minutes, and seconds, and units of declination are degrees, arcminutes, and arcseconds. First, a continuum map was created by averaging the channels between 190 and 213 km s 1 and between 265 and 279 km s 1. This was subtracted from each channel map. A total of 15 continuum sources can be discerned, which are further discussed in x 3.4. Next, the channel maps had to be deconvolved to correct for beam artifacts. There are three major artifacts. First, the H i emission is distributed over an area that is larger than the primary beam of the telescope. As the two fields we observed cover most of the bright core of the HVC (see Fig. 2), our data reduce this problem. Second, the 15.3 m baseline increment yields a grating ring, at a radius of about Deconvolution (see below) gets rid of most of this, but because the cloud is more extended than the mapped field, the possibility of erroneous deconvolution remains. Finally, since the UV coverage is regular and fairly complete, there is a bowl due to the unobserved 15 m baseline. Short of obtaining single-dish data to fill in this baseline, the best method to correct for the bowl is the Multi-Resolution CLEAN (MRC; Wakker & Schwarz 1988). This method preserves the simplicity of the basic CLEAN but is better adapted for deconvolving extended sources. MRC works by first smoothing the map and beam pattern, in order to emphasize the large-scale structure and enhance the signalto-noise ratio (S/N). We increased the beam diameter by a factor of 2. Concurrently, a difference map and beam are created, which are the differences between the original and smoothed versions. Both maps are separately cleaned in a standard manner and then properly recombined. Even after this step there will still be missing flux in the interferometer map, as the deconvolution step can only reconstruct part of the flux on the 15 and 0 m baselines. For example, the observed UV plane of a source with a Gaussian shape will also look like a Gaussian, but with a hole in the center. The deconvolution can fit this Gaussian and restore the hole, but only if the width of the Gaussian was sufficiently large that the S/N was large enough at the observed baselines. Thus, the deconvolution cannot restore the very extended structure. In particular, for the synthesized beam of the observations discussed in this paper, any structure on scales larger than 30 0 will be completely filtered out, while structures on scales of about 20 0 are suppressed by a factor of 2. The mapping and deconvolving were done separately for each of the two fields. The resulting maps were then corrected for primary beam attenuation and linearly combined. To suppress the noise near the edges of the primary beam, a maximum correction of a factor of 6 was allowed. 3. RESULTS 3.1. Channel and Velocity-Declination Maps Figure 3 shows the final channel maps that were created, for velocities ranging from 220 to 259 km s 1.Mapsat2 0 resolution are shown, since the S/N of the full-resolution maps is so low that the contours would encircle many very small areas and the plot becomes impossible to read. The two circles in each panel indicate the FWHM of the primary beam of each pointing. The star shows the position of NGC 3783 (11 h 39 m 0198, ). There is structure on scales down to the size of the beam. Several subconcentrations can be discerned, most of which are fairly round; there is also a linear feature crossing the northern field from upper left to bottom right. Note that in the area where the two fields overlap there seems to be a gap in the emission at velocities of km s 1 and km s 1. This is probably due to the fact that in these channels the signal is weak enough that primary

5 No. 4, 2002 H i FINE STRUCTURE IN MAGELLANIC TIDAL DEBRIS 1957 Fig. 2. Column density map of cloud WW 187 as extracted from the HVC-reduced (minmed5 ) HIPASS (Putman et al. 2001), for velocities between 204 and 270 km s 1. Contours are at levels of 0.10, 0.50, 0.90, and 1: cm 2. The star near 11 h 39 m, shows the position of NGC The circles show the area included in the final maps (primary beam correction less than a factor of 6). beam attenuation makes it disappear below the noise level. It is thus not recovered, leading (probably) to an artificial gap. In the intermediate channels ( km s 1 ), the emission in the overlap region can be seen in the original maps all the way to the edge of the primary beam, and in the mosaicked map there is thus a smooth connection. To show the velocity structure of the HVC better, a set of velocity-declination maps was created from the mosaicked channel maps. These are shown in Figure 4. A cut was made through every fourth R.A. pixel, i.e., cuts are (about two beams) apart. The R.A. of each cut is indicated on each of the panels of Figure 4. From this figure it becomes apparent that most subconcentrations show a small or no velocity gradient. The profile on NGC 3783 is found from the cut indicated by the thick line in the rightmost panel of the second row Combined Interferometer and Single-Dish H i Map All channel maps were smoothed from the original resolution of to resolutions of 2 0,4 0, and 8 0, with the purpose of studying the influence of resolution on parameters such as peak brightness temperature, column density, concentration size, and line width. The smoothed maps were also used in the construction of a total column density map: a cutoff of 3 in the smoothed map at one step worse resolution was used to create a mask in order to select those pixels in the map at the original resolution where the S/N is suffi-

6 1958 WAKKER, OOSTERLOO, & PUTMAN Vol. 123 Fig. 3. Tile plot of primary beam corrected mosaicked channel maps, at 2 0 resolution. The velocity of each channel is indicated in the upper right corner of each panel. The maximum primary beam correction is a factor of 6. The gray scale ranges from 0 to 130 mjy beam 1 (20 ). The contour level at 39 mjy beam 1 (6 at the center) is shown to emphasize the areas with significant signal (except near the edges where the primary beam correction increases the noise level). The two circles indicate the FWHM of the primary beam (33<7). The star shows the position of NGC 3783, and the smaller beam around it shows the size of the Green Bank beam. ciently high. The masked maps were then corrected for primary beam attenuation, converted to brightness temperature (T B ), and integrated. The resulting total column density maps (based on interferometer data only) are shown in Figure 5. These are of comparable quality to the similar interferometer-only total H i maps for other HVC cores shown by Wakker & Schwarz (1991). The maps in Figure 5 are for the fine structure only, that is, structures with size scales smaller than about The more extended background has been filtered out. This was found from the HVC-reduced (Putman et al. 2001) H i Parkes All Sky Survey (HIPASS; Barnes et al. 2001). The standard HIPASS reduction method filters out emission that extends over more than 2 in declination, while the HVC method (referred to as minmed5) recovers the emission unless it completely fills an 8 scan. The Parkes telescope has a beamwidth of 14 0, but the gridding process that combines the individual scans decreases the spatial resolution to 15< The resulting Parkes column density map (for velocities between 204 and 270 km s 1 ) is shown in Figure 2. The primary beams of the ATCA observations are shown, as is the position of NGC The two interferometer fields cover most of one of the main concentrations. The availability of this low-resolution map of the HVC allowed us to construct a total column density map containing both the extended and the fine structure. This was done in the manner described by Schwarz & Wakker (1991). This method is based on the realization that the single-dish map contains (1) the emission filtered out by the interferometer and (2) the emission seen by the inter-

7 No. 4, 2002 H i FINE STRUCTURE IN MAGELLANIC TIDAL DEBRIS 1959 Fig. 4. Tile plot of primary beam corrected velocity-declination maps, with the right ascension of the cut indicated in the upper left corner. The gray scales and contour level are the same as for Fig. 3. The position of NGC 3783 is indicated by the thick line in the rightmost panel of the second row. ferometer. By smoothing the primary beam corrected interferometer column density map to a 16 0 beam, component (2) can be isolated. Subtracting this smoothed interferometer map from the single-dish map leaves the emission that was filtered out. The final map is then found by adding the interferometer map back in. The steps and result of this process are shown in Figure 6. Figure 6a gives the interferometer map, Figure 6b the interferometer map smoothed to 16 0 resolution. Figure 6f shows the original single-dish map. Figure 6e gives the residual single-dish map (i.e., single dish minus smoothed interferometer); this is the emission that was filtered out by the interferometer. Figure 6d shows the final map: residual single dish plus interferometer. Figures 6g, 6h, and 6i show an enlargement of the interferometer and combined map in the region around NGC The combination process detailed above was repeated for the interferometer maps at resolutions of 2 0,4 0, and 8 0. The results are shown in Figure 7. This figure also shows the single-dish (Parkes) map as well as the Parkes map smoothed to a beam of Since the shortest baseline in the interferometer map is 30.5 m, there is some sensitivity to structures up to 24 0 in size, but only within the 35 0 size of the primary beam. Thus, in the inner 10 0 of the two interferometer fields, the combination with the 16 0 Parkes data provides a sample of structures on all scales. Outside this region the sensitivity of the interferometer for structures on larger scales becomes progressively worse, until it is about zero for structures more than 30 0 from the field center. Thus, the combined maps are most reliable within about 20 0 from the field centers. Farther out the combined map is dominated by the largescale structure. It is instructive to compare the combined map of Figure 7 to the interferometer-only map of Figure 5. This shows that the interferometer picks up most of the structure and that the visual impression one gets of the details does not change much. However, adding the single-dish map makes the contrast between cores and off-core regions smaller. Comparing Figures 7a and 7g shows that the linear feature running across the northern half of the interferometer map hardly shows up in the 16 0 single-dish map: at full reso-

8 1960 WAKKER, OOSTERLOO, & PUTMAN Vol. 123 Fig. 5. H i column density as observed with the ATCA interferometer, at four different resolutions. Top left: ; top right: ; bottom left: ; bottom right: The darkest gray scale corresponds to a column density of 1: cm 2. The size of the beam is indicated by the filled circle in the bottom left corner, the position of NGC 3783 by the star. lution, the peak column density in the densest concentration (at R:A: ¼ 11 h 39 m 30 s,decl: ¼ ) is about 1: cm 2, of which 1: cm 2 is in the small-scale structure. Thus, the interferometer recovers two-thirds of the flux in this concentration. On the other hand, the main concentration seen in the lower half of the single-dish map has a peak column density of 3: cm 2, of which about 1: cm 2, or 55%, is recovered in the interferometer map. The smallest structures in this part of the map seem to be slightly offset to the east from the brightest part of the more extended structure. The bulge in the cm 2 column density contour seen in the upper right part of the single-dish map is completely due to the concentration seen at 11 h 38 m 20 s, On the other hand, between the concentrations most of the flux is in extended structure. Thus, the cloud appears like a set of dense cores embedded in a smooth extended envelope. Comparing Figures 7a, 7b, 7d, and 7e shows that the cores seen at 8 0 resolution show internal structure at 4 0 resolution. At 2 0 resolution these structures break up into smaller cores, which show more internal structure and break up at 1 0 resolution. The amount of small-scale structure outside the cores is limited, however. Thus, the structure is hierarchical, i.e., smaller cores are embedded in larger cores Velocity Field Figures 7c and 7f show the velocity field in the core of the HVC. This was obtained by fitting one or two Gaussians to the spectrum at each pixel. The figure shows the result for fits made to the maps at 2 0 resolution. Fits were also made to the maps at 1 0,4 0, and 8 0 resolution. However, at 1 0 resolution, many pixels cannot be fitted reliably because of a low S/N. The results at 4 0 and 8 0 are not as detailed. The pattern of velocities that results can be interpreted as showing that the cores have more or less constant central velocity but that different cores have different velocity. In some cases there is a regular velocity gradient for the gas between cores, but often the velocity jumps by 3 10 km s 1 over a space of a few pixels. These characteristics are very similar to those seen in other HVC cores observed at high resolution (Wakker & Schwarz 1991). These also showed that the profile widths of spectra observed with beams larger than 10 0 are mostly due to the velocity differences between cores.

9 No. 4, 2002 H i FINE STRUCTURE IN MAGELLANIC TIDAL DEBRIS 1961 Fig. 6. Intermediate steps in constructing the combined map of H i column density. Original data are in panels (a) and ( f ); then (a) is smoothed to (b), which is subtracted from ( f ) to create (e); next (a) and (e) are added to create (d ). The position of NGC 3783 is shown by the star. For panels (e) and ( f ) column density contours are at ð0:70 1:30Þ10 20 cm 2 in steps of cm 2, with a thick line for the contour at cm 2. Contour levels in panel (b) are 1.25, 2.5, 3.75, and 5: cm 2 ; in panel (g)4: and 9: cm 2 ; and in panel (h)0: and 1: cm 2. The gray-scale values are indicated by the wedge on the right. The size of the interferometer beam is indicated in the lower left corner of panels (g), (h), and (i). The primary beams of the interferometer fields are shown in panels (a) and ( f ). The size of the Parkes beam, centered on NGC 3783, is shown in panels (a), (d ), and ( f ). The box in panel (a) gives the area of detail in panels (g), (h), and (i). A histogram of the fitted FWHM line widths (Fig. 8) shows that 90% of the profiles have a width between 3 and 9 km s 1, with the modus of the distribution at a width of 5 km s 1. There are also a few components with larger widths, but these are mostly located in regions where the S/N is low, making them suspect. A width of 5kms 1 sets an upper limit on the kinetic temperature of the gas of about 500 K Search for Absorption by Continuum Sources Several weak continuum sources are present in the field. This gives the possibility of determining a spin temperature for the HVC by comparing the emission with the H i absorption toward the continuum source. A total of 15 continuum sources can be found. To allow possible further investigations into determining the H i spin temperature of WW 187,

10 1962 WAKKER, OOSTERLOO, & PUTMAN Vol. 123 Fig. 7. Maps of the velocity field and the H i column density. Panels (a), (b), (d ), and (e) show the combined ATCA and Parkes data, at resolutions of ,2 0,4 0, and 8 0. Panel (g) shows the Parkes data (16 0 beam), panel (h) the Parkes data smoothed to The size of the beam is indicated by the filled circle in the bottom left corner of each panel. The column density scale is shown by the left wedge in the lower right panel. Panels (c) and ( f ) show the velocity of the Gaussian components fitted to the spectra at each pixel, for the map at 2 0 resolution. The scale is indicated by the right wedge in the lower right panel. Two panels are shown as there are many pixels where two velocity components can be fitted. we list in Table 2 the positions and fluxes of these continuum sources, as well as the parameters pertaining to the measurement of spin temperature. Object 2 is NGC Only three others can be identified with a known source in the SIM- BAD database: the radio source PMN J (object 3), the X-ray source 2E (object 4), and the IR source IRAS (object 11). The method we followed to determine a spin temperature was fully described by Wakker, Vijfschaft, & Schwarz (1991). However, just six of the continuum sources yield a spectrum with an S/N larger than 5. No obvious absorption is visible toward any source, implying that only lower limits can be set on the spin temperature. Assuming a detection limit of 3 for absorption, this yields a lower limit on the spin temperature of about 10 K, which is unfortunately not very informative. 4. N(H i) IN THE DIRECTION OF NGC 3783 The main motivation for observing the H i in WW 187 at high angular resolution was to obtain an improved value for N(H i) in the direction of NGC 3783, in order to derive more reliable values for the abundances of heavy elements seen in absorption. However, the Parkes data we originally had available have a velocity resolution of only 26 km s 1,soit

11 No. 4, 2002 H i FINE STRUCTURE IN MAGELLANIC TIDAL DEBRIS 1963 Fig. 8. Distribution of line widths fitted to the spectra at 2 0 resolution. is not possible to obtain a combined spectrum (we note that data with higher velocity resolution will be available in the near future). Instead, we used the same 1 km s 1 resolution spectrum as Lu et al. (1994), which was obtained with the Green Bank 140 ft telescope (21 0 ). We again applied the method described in x 3.2 to combine interferometer and single-dish data, but now just for the direction toward NGC 3783 and using the 21 0 Green Bank spectrum as the singledish spectrum. Figure 9 presents (a) the interferometer spectrum, (b) the spectrum of the emission recovered by the interferometer within a 21 0 beam, (c) the original single-dish spectrum, and (d ) the final spectrum in the direction of NGC Gaussian fits were made to each spectrum, the parameters of which are given in Table 3. The single-dish profile gives a column density toward NGC 3783 of ð1:12 0:03Þ10 20 cm 2. The error in this number is determined just by the residual of the fit; it does not include a possible baseline offset. Of the single-dish flux, the interferometer recovers ð3:4 0:2Þ10 19 cm 2, where again the error is just the statistical error. Since the cloud is larger than the size of the mapped fields, there may be unrecognized structure that is due to grating rings of cores outside the field rather than to real signal. Assuming that there is a core of 1: cm 2, which is suppressed by a factor of 6 (i.e., at the edge of our field) and of which 25% is recovered by the interferometer, we estimate a possible systematic error of cm 2. Thus, the extended background contains ð7:8 0:6Þ10 19 cm 2. Before fitting the combined spectrum, the interferometer data were Hanning smoothed to reduce the noise to about 1.2 K. A single-component fit then yields a final column density in the direction of NGC 3783 of ð8:3 2:0Þ10 19 cm 2. Looking at the total column density map near NGC 3783 (Fig. 6g), it is clear that NGC 3783 sits within a low point in the fine structure. From Figure 9 it can be seen that the recovered flux has a profile shape that is similar to that of the single-dish profile. Both show a narrower central component on top of a broader underlying component. The recovered spectrum also appears to have a small hump near v LSR ¼ 253 km s 1 with height 0.27 K, but this has an FWHM of only 7 km s 1 and a column density of only cm 2. This wiggle can also be explained as due to a small deficit in the recovered flux near 250 km s 1, which could be caused by problems with the deconvolution, related to the fact that most of the signal at this velocity occurs near the edges of the northern field, so that some of it has an insufficiently high S/N to be discerned. Referring to Figures 4 and 7, one can see that the line wings of the recovered flux profile correspond to concentrations with narrow width but at velocities away from the central peak, not to concentrations with intrinsically broad TABLE 2 Continuum Sources Number (1) R.A. (2000.0) (2) Decl. (2000.0) (3) Flux (Jy) (4) PBC (5) Flux (Jy) (6) T B,c (K) (7) max (8) S/N c (9) T em (K) (10) T ext (K) (11) T s (K) (12) < > < > < > > > < > < > < > > < > < > < > < > > > < >3 Note. Col. (4): continuum flux of source in the original data. Col. (5): correction factor for primary beam attenuation. Col. (6): corrected continuum flux. Col. (7): brightness temperature of continuum source. Col. (8): optical depth of a 3 absorption. Col. (9): S/N in spectrum of continuum source. Col. (10): brightness temperature of emission with small-scale structure near continuum source. Col. (11): brightness temperature of extended structure near continuum source. Col. (12): lower limit set on spin temperature.

12 1964 WAKKER, OOSTERLOO, & PUTMAN Vol. 123 Fig. 9. Spectra in the direction of NGC The filled circles in panel (a) show the ATCA spectrum (1 0 beam) after Hanning smoothing, for which the rms is 1.2 K. The solid line is obtained by smoothing the spectrum to 5 km s 1 velocity resolution. Panel (b) shows the flux recovered by the interferometer and the three-component Gaussian fit. Panel (c) shows the Green Bank spectrum (21 0 beam) and the two-component Gaussian fit. Panel (d ) shows the final spectrum toward NGC 3783: single dish minus recovered plus interferometer; the smooth curve shows the best-fitting single Gaussian. profiles. Whether or not this is also true for the extended structure is entirely possible but not immediately clear. To find this out, the higher velocity resolution single-dish data should be used to combine the extended and fine structure in the manner described in x 3.2, separately for each velocity channel. Velocity (km s 1 ) TABLE 3 Gaussian Fits for Spectra toward NGC 3783 T B (K) FWHM (km s 1 ) Green Bank 140 ft Spectrum N(H i) (cm 2 ) Flux Recovered by ATCA in Green Bank Beam Combined Spectrum on NGC DISCUSSION OF H i FINE STRUCTURE AND CLOUD PHYSICS 5.1. Characteristics of H i Fine Structure The combined ATCA and Parkes map allows a study of the characteristics of the small-scale structure of the HVC. However, this paper is not the place for a full-fledged analysis. Only a few obvious points will be made, which relate to the interpretation of brightness temperature, column density, and volume density measurements. Figure 10 shows the peak brightness temperature and peak column density measured within each of the 10 concentrations and six off-core areas outlined in Figure 11. In the left panels the H i interferometer spectrum is shown at the position of the highest brightness temperature. This clearly shows that within these cores the apparent spectrum is strongly influenced by the resolution. Velocity gradients appear to be small enough that the velocity of the peak does not shift.

13 No. 4, 2002 H i FINE STRUCTURE IN MAGELLANIC TIDAL DEBRIS 1965 Fig. 10. Plots showing the effect of resolution of measured parameters for the concentrations outlined in Fig. 11. Left panels: Spectra at 1 0,2 0,4 0, and 8 0 resolution, at the position where the brightness temperature peaks. Middle panels: Filled circles connected by solid line show the peak brightness temperature in the concentration as function of resolution, while the dotted lines show curves expected if the structure is very extended (horizontal line), a Gaussian with FWHM ¼ 20 0,4 0,2 0, and 1 0, or a point source. Right panels: Column density measured as a function of offset to the position of the maximum column density in the concentration, in the pattern shown in Fig. 11. The middle set of panels shows the run of peak brightness temperature with resolution. This is compared to the expectation if the column density crosscut through the core were Gaussian: the dotted lines give the expected run of T B with resolution for Gaussian cores with FWHM ¼ 20 0,4 0,2 0,and 1 0, as well as for a point source (bottom line) and infinitely extended structure (horizontal line). This reveals that, although not actually having a Gaussian shape, crosscuts through most of the cores are most closely matched by those of a Gaussian with FWHM between 1 0 and 2 0. When the beam becomes much larger than the typical distance between cores ( ), they no longer stand out and T B varies much more slowly. We note, however, that the intrinsic sizes of the cores are probably less than their apparent sizes, since they are convolved with the interferometer beam (see also x 5.2). One of the larger areas where signal is present in the interferometer map, but where no clear core can be defined

14 1966 WAKKER, OOSTERLOO, & PUTMAN Vol. 123 Fig. 11. Full-resolution map with several overlays. The contour is at 1: cm 2. Filled stars with a number next to them give the positions of the continuum sources listed in Table 2. The boxes delineate the concentrations for which the peak brightness temperature as function of resolution and the column density contrasts are studied (x 5.1). The star pattern in the upper right corner shows the pattern of positions at which column densities are plotted in the right column of Fig. 10. (labeled Fld2 mid in Fig. 10), behaves more like Gaussians with an FWHM of , while outside the core (panels labeled OffCore top/bot/right ) T B shows very little variation with resolution, as expected. The fact that the curve giving T B as a function of resolution flattens above 10 0 indicates that the cores are embedded in an extended envelope that is fairly smooth on scales larger than about This conclusion is consistent with the impression one gets from the maps in Figure 7. The right panels in Figure 10 show how the column density drops away from the center of each core. The maximum column density is given by the point at radius 0 and can reach up to 3: cm 2. In almost all cases the column density has dropped considerably 1 0 away, but it drops more slowly farther out. This is another indication that the cores are very small. Also noteworthy is that the cores are not symmetrical: in almost all cases N(H i) can differ by up to a factor of 2 depending on which side of the core one looks. Some, but not all, of this variation is due to noise. Of particular interest is the pattern for the cores labeled Fld1 mid2 and Fld1 right. NGC 3783 lies almost exactly between these two (see Fig. 11), some 4 0 from their centers. No great variations are seen in N(H i) at4 0 radii for these cores, suggesting that the value of ð8:3 2:0Þ10 19 cm 2 found for N(H i) is reliable. The peak column density in these two cores is 1: and 1: cm 2, respectively. From the right set of panels in Figure 10, it is clear that near the centers of cores the column density contrast can be a factor of 2 over 1 0 scales and a factor of 3 on scales of several arcminutes. Away from the core centers the contrast is much less. A final characteristic of the small-scale structure is displayed in Figure 12. This gives the distribution of N(H i) in the mosaicked field for the combination of interferometer and single-dish data, in the region where the primary beam correction is below a factor of 4. This shows that an observation of a random direction within the cloud is most likely

15 No. 4, 2002 H i FINE STRUCTURE IN MAGELLANIC TIDAL DEBRIS 1967 Fig. 12. Number of pixels in cm 2 intervals of column density. N(H i) is given on a linear scale on the right and on a logarithmic scale on the left. The different curves show the counts at different resolutions, as indicated by the labels. A power law would result in a straight line in the plots on the right, and an exponential gives a straight line on the left. to yield a column density of about cm 2, with 50% of the values falling between 0: and 1: cm 2. In only 4% of the directions is the column density greater than cm 2. For column densities above 1: cm 2 the slope of the distribution is not a power law [in which case the curves in the upper right panel showing logðcountþ vs. log NðH iþ would be straight lines] but an exponential: count / exp½ NðH iþš. It is clear from Figure 12 that at lower resolution the highest column densities are missed, but that the distribution does not change for the lower column densities. This shows up as a turndown in the distribution, which happens at 1: cm 2 for 8 0, 1: cm 2 for 4 0, 2: cm 2 for 2 0, and ð2:8 3:0Þ10 20 cm 2 for 0< Internal Structure The maps at the highest possible resolution of show structure down to the scale of the beam (Fig. 7c). There are about 10 such cores, with peak column densities ranging from 1: to 3: cm 2. These cores probably sit deep within the cloud. They are surrounded by the cloud envelope, which has a column density ranging from to cm 2. To get the column density associated with the central core itself, we took the difference between the peak column density and the mean column density in a region right next to it. This mean varies between 0: and 1: cm 2 and is typically 1: cm 2. The column densities associated with the cores themselves are then in the range ð1:5 2:6Þ10 20 cm 2. These cores show relative motions of up to 20 km s 1 (Fig. 7c). This is similar to the results for other HVC cores obtained by Wakker & Schwarz (1991). Gaussian fits to the line profiles at each pixel yield a modal width of 5 km s 1, with a range of 3 9 km s 1 (Fig. 8). This sets an upper limit to the kinetic temperature of the gas of K for the cores visible in the map. The highest brightness temperatures that are observed in the cores are K, setting a lower limit to the kinetic temperature. A study of the 15 continuum sources in the field unfortunately failed to yield interesting limits on the H i spin temperature. Recently, Sembach et al. (2001) observed absorption lines from N i,nii,siii,andfeii, as well as H 2 (making this only the second HVC in which molecular hydrogen has been detected). They derive a molecular fraction of 1: and H 2 rotational temperatures of T 01 ¼ 130 K and T 23 ¼ 240 K.

16 1968 WAKKER, OOSTERLOO, & PUTMAN Vol. 123 Thus, the temperature of the cool component of the gas in the cores of WW 187 is probably 150 K. Elsewhere in the ISM kinetic temperatures for H i tend to be in the range K (Dickey & Lockman 1990). At kpc distance 1 0 corresponds to pc. Since at 1 0 resolution the cores have excess column densities on the order of ð1:5 2:6Þ10 20 cm 2, the average density within this volume is cm 3. At 1 0 resolution, the distribution of the number of pixels with a given column density (Fig. 12) shows a slight turndown at N(H i) of about ð2:8 3:0Þ10 20 cm 2, similar to what happens at lower column densities after smoothing. Thus, the cores might show higher column densities if they were observed at higher resolution. There now are 1 2 pixels where N(H i) reaches 3: cm 2 (in the Center and the Fld2 bot cores. If our beam area were 4 times smaller, and if the distribution were to keep the same form, we would expect 4 8 pixels at NðH iþ ¼3: cm 2 and 1 2 pixels at NðH iþ ¼4: cm 2, and we would derive densities that are about 4.5 times higher, i.e., 7 25 cm 3. For most of the reliable cases (cores near the beam centers, named Fld1 mid1, Fld1 mid2, Fld1 right, and Center ), the run of brightness temperature with resolution (Fig. 10, middle panels) seems to indicate a size between 1 0 and 2 0. However, this may be slightly misleading, since a core with an intrinsic size of about 1 0 would be smeared with the beam to appear as a core with a size of about 1<5. Thus, the cores appear to be almost resolved but may not quite be completely resolved. Thus, considering (1) that the distribution of column densities shows a turndown even at 1 0 resolution and (2) that, not taking into account beam smearing, the cores appear to be between 1 0 and 2 0 in size, we argue that the cores may be a factor of 2 4 (or more) smaller in area (or about a factor of smaller in diameter). This would imply that the cores may be as small as 5 15 pc and that densities as high as 20 cm 3 may be present. Sembach et al. (2001) assumed a volume density of 10 cm 3 and estimated the formation time for the cm 2 of H 2 that they detected. This depends on the formation rate, which is usually taken to be ð1 3Þ10 17 cm 3 s 1 (van Dishoeck & Black 1986). Sembach et al. (2001) thus found a formation time (the time to completely convert the H i to H 2, assuming that there is no or little destruction) of Myr. 3 Since this is a significant fraction of the orbital time of the Magellanic Clouds (2 Gyr), they favored the idea that the molecular hydrogen has survived the destructive radiation since the gas was stripped. Considering that the line widths and relative motions of the cores are about 5 km s 1, with sizes of about 10 pc, it takes only about 2 Myr for the core to shift over its own diameter. A typical distance between two cores is about 10 0, or pc, so the cores can mingle on timescales of tens of megayears. This is a small fraction (<1%) of the orbital time. It thus seems unlikely that the cores have survived since the gas was extracted from the SMC. However, with core densities as high as 20 cm 3, the formation time may be as low as Myr. Suppose that the 3 Note that in their eq. (3) Sembach et al. (2001) use a density of 3 cm 3, which should give a formation time of 1 Gyr for R ¼ cm 3 s 1, rather than the stated 100 Myr. core lifetimes are only about 1 Myr. In that time only about 1% of the H i can be converted to H 2 before the core dissolves. The H 2 will disappear slowly since the dissociating radiation is comparatively weak at distances of kpc from the Milky Way. Such spreading of H 2 can easily yield the observed molecular fraction of 1: toward NGC Since we see several cores, and since they should quickly dissolve, cores must constantly form and dissolve, providing a source for the molecular hydrogen. Finally, we can estimate the mass of the smallest, densest cores by combining the column density and diameter: M ¼ m H i nð4r 3 =3Þ ¼ 330ðN=10 20 ÞðR 2 =10Þ M. It is interesting to compare this to the Jeans mass of a core, M J 2: n 0:5, where is the velocity dispersion in km s 1 and n the particle density in cm 3. Thus, for cores with the highest column density ( cm 2 ), highest volume density (20 cm 3 ), and lowest velocity dispersion (5 kms 1 ), the core mass (10 3 M ) still is far below the Jeans mass (10 6 M ), suggesting that star formation is unlikely to take place WW 187 in the Context of the Multiphase ISM Model The fact that small, cold, dense cores are present in the WW 187 HVC can be used to derive some limits on the physical conditions of its environment, as well as to set limits on the structure of the Galactic halo. Specifically, a comparison with the model of Wolfire et al. (1995) is of interest. These authors apply the basic physics of heating and cooling in neutral gas to derive the relation between density and pressure for different combinations of metallicity and dust content, for a range of different external pressures (expressed as height above the plane). To tie the internal cloud physics to a model of the Milky Way, they analyzed the parameters for which a two-phase medium can exist, assuming that the Milky Way is surrounded by hot halo gas that provides the external pressure. This halo is constructed as consisting of isothermal hot gas at a temperature of ð1 2Þ10 6 K, which fills the potential of the Galaxy (with fairly standard parameters). Given a certain metallicity and dust content (which determine the cooling and heating of the gas), a two-phase medium occurs when the external pressure provided by the hot halo balances the internal pressure at a value such that both a dense, cold phase (n 100 cm 3, T 25 K) and a tenuous, hot phase (n 1cm 3, T 2500 K) are thermally stable. Wolfire et al. (1995) then study at what range of distances this two-phase medium occurs. The appropriate model to compare to WW 187 is that for gas stripped from the Magellanic Clouds, with a metallicity about 0.3 times solar and a dust content about 0.3 times that in the local ISM. Given the model for the hot halo, the presence of a twophase medium is predicted to be possible for a z range of 1 20 kpc if T halo ¼ 10 6 K, but for all z if T halo ¼ K. Wolfire et al. (1995) prefer the T ¼ 10 6 K value, and since previous lower resolution observations of the Magellanic Stream had not revealed cold, dense cores, they concluded that the absence of cold cores in the Stream was consistent with their model. However, HVC WW 187 clearly shows a multiphase medium. Assuming a distance of kpc, typical core densities are 2 10 cm 3, and a typical core temperature is 150 K (see x 5.2), corresponding to pressures on the order