L75 UNUSUAL DEPLETIONS TOWARD THE SMC STAR Sk 155 ...

The Astrophysical Journal, 554:L75–L79, 2001 June 10 䉷 2001. The American Astronomical Society. All rights reserved. Printed in U.S.A.

UNUSUAL DEPLETIONS TOWARD THE SMC STAR Sk 155—DIFFERENCES IN DUST COMPOSITION IN THE SMC INTERSTELLAR MEDIUM?1 Daniel E. Welty,2 J. T. Lauroesch,3 J. Chris Blades,4 Lewis M. Hobbs,5 and Donald G. York2,6 Received 2001 April 19; accepted 2001 May 8; published 2001 June 6

ABSTRACT We report initial results from an analysis of Hubble Space Telescope/Space Telescope Imaging Spectrograph echelle spectra of interstellar absorption toward Sk 155, located in the “wing” region of the SMC. There are at least 12 Milky Way (disk and halo) and at least 25 SMC components discernible in the profiles of various neutral and singly ionized species. Fits to the line profiles yield column densities and relative elemental abundances for the individual components. In the SMC components, the depletions of Fe and Ni range from mild ([Fe, Ni/Zn] ∼ ⫺0.3 to ⫺0.8 dex) to severe ([Fe, Ni/Zn] ∼ ⫺1.7 dex); Mg and Si, however, appear to be essentially undepleted ([Mg, Si/Zn] ∼ ⫺0.1 to ⫹0.2 dex) throughout. The combination of severe depletion for Cr, Mn, Fe, and Ni and negligible depletion for Mg and Si has not been seen in any Galactic sight line. If Si is generally undepleted in the SMC, then current models of the SMC dust—which rely heavily on silicates—will have to be reconsidered; oxides and/ or metallic grains may dominate. We briefly discuss possible the implications of these SMC depletions for understanding the elemental abundances in QSO absorption-line systems. Subject headings: dust, extinction — galaxies: abundances — galaxies: ISM — ISM: abundances — Magellanic Clouds — stars: individual (Sk 155) The relative abundances in the main components toward one star in the Magellanic Bridge (where the overall metallicities may be even lower) also are comparable to those in the halo clouds (Lehner et al. 2001). These similarities in the relative gas-phase abundances in the Milky Way and MCs suggested that the interstellar depletion patterns might also be similar in the three galaxies—despite the differences in metallicity and dust-to-gas ratio. In this Letter, we report the initial results from an analysis of Hubble Space Telescope/Space Telescope Imaging Spectrograph (STIS) echelle spectra of interstellar absorption toward the SMC star Sk 155. The main SMC interstellar components toward Sk 155 exhibit more severe depletions (of some elements) than those inferred toward Sk 108. The overall pattern of relative abundances, however, does not resemble any of those found (to this point) in the Galactic ISM. In the following sections, we briefly describe the new STIS spectra, the derived interstellar abundances, and the implications of those abundances for models of SMC dust and for interpreting the gas-phase abundances seen in QSO absorption-line systems. A subsequent paper will provide a more comprehensive analysis of all the interstellar components—both Galactic and SMC—along this line of sight (D. E. Welty et al. 2001, in preparation).

1. INTRODUCTION

Studies of the interstellar medium (ISM) in the Magellanic Clouds (MCs) explore somewhat different environmental conditions than those typically probed in our own Galactic ISM. The MCs are characterized by lower overall metallicities (⫺0.3 dex for the LMC, ⫺0.6 to ⫺0.7 dex for the SMC), lower dust-to-gas ratios, generally stronger ambient radiation fields, and significant differences in UV extinction (especially in the SMC). Determination of elemental abundances, depletions, and physical conditions in the MCs’ ISM therefore provides interesting tests for theoretical models of both interstellar clouds and interstellar dust grains. Understanding the ISM of the MCs should also yield insights into the nature of the absorption-line systems seen toward distant QSOs, which exhibit some similar properties. Several recent studies using high-resolution optical and/or UV spectra have provided accurate gas-phase abundances for the predominantly neutral clouds along several MC lines of sight. Toward SN 1987A, the main LMC component groups have relative abundances [X/Zn]7 similar to those found for warm, diffuse Galactic disk clouds; the “intermediate velocity” groups have relative abundances more to those in Galactic halo clouds (Welty et al. 1999a). Toward Sk 108, the relative abundances in most of the SMC components resemble those in the Galactic halo clouds (Welty et al. 1997; Mallouris et al. 2001).

2. STIS SPECTRA OF Sk 155

The O9 Ib star Sk 155 (AzV 479; V p 12.42; B⫺V p ⫺0.16; Ardeberg & Maurice 1977; Lennon 1997), located in the “wing” region of the SMC at a distance of about 65 kpc (Wayte 1990), lies behind an SMC hydrogen column density N(H i) ∼ 3.5 # 10 21 cm⫺2 (Fitzpatrick 1985). Sk 155 was observed with the STIS during five visits in 2000 April and October, under GO program 8145. Four wavelength settings were obtained using the high-resolution E140H and E230H echelle gratings, at a resolution of about 2.7 km s⫺1 (FWHM). The procedure developed by Howk & Sembach (2000) was used to improve the background (scattered-light) correction in the extracted spectra. Spectra from multiple exposures and adjacent orders were coadded; the summed spectra were then normalized via polynomial fits to the continuum regions adjacent to various interstellar ab-

1

Based on observations with the NASA/ESA Hubble Space Telescope, obtained at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. 2 University of Chicago, Astronomy and Astrophysics Center, 5640 South Ellis Avenue, Chicago, IL 60637; [email protected], don@oddjob .uchicago.edu. 3 Northwestern University, Department of Physics and Astronomy, 2131 Sheridan Road, Evanston, IL 60208-2900; [email protected]. 4 Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218; [email protected]. 5 University of Chicago, Yerkes Observatory, 373 West Geneva Street, Williams Bay, WI 53191; [email protected]. 6 Also at the Enrico Fermi Institute. 7 [X/Y] p log [N(X)/N(Y)] ⫺ log [A(X),/A(Y),]. The solar reference abundances A, are the meteoritic values listed by Anders & Grevesse (1989).

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Fig. 1.—Profiles of selected UV lines toward Sk 155, observed with STIS at resolutions of about 2.7 km s⫺1. Only the SMC absorption, at 60 km s⫺1 ⱗ v ⱗ 240 km s⫺1, is shown; tick marks indicate the components found in fits to the line profiles. Component groups A, B, and C are noted above several of the profiles. Note the striking differences between the profiles of the mildly depleted species Mg ii, Si ii, S ii, and Zn ii (left) vs. those of the more heavily depleted species Cr ii, Mn ii, Fe ii, and Ni ii (right)—indicative of significant componentto-component differences in relative gas-phase abundances. The stronger Fe ii l2586 line reveals additional lower column density components. For Mg ii l1239, most of the absorption between 110 and 130 km s⫺1 is due to Milky Way Mg ii l1240; additional absorption from SMC Ti ii l1910.6 (with Ti ii l1910.9) and from SMC Mg i l2026 (with Zn ii l2026) is noted.

sorption lines of interest. The empirical signal-to-noise ratios determined from the continuum fits range from about 15 to 40 per half-resolution element. The STIS spectra reveal very complex interstellar component structure in this line of sight, with at least 12 Milky Way and 25 SMC components discernible over a total velocity range of nearly 300 km s⫺1. Interstellar absorption is evident from trace neutral species, from many species dominant in neutral (H i) gas, from more highly ionized species, and from several excited fine-structure levels. Because we focus in this Letter on the abundances in the main neutral components in the SMC, Figure 1 shows lines from dominant ions (most of moderate strength) over the heliocentric velocity range from ⫹50 to ⫹250 km s⫺1; Milky Way disk and halo components are seen from about ⫺60 to about ⫹50 km s⫺1. The stronger Fe ii l2586 line reveals weaker SMC components located astride the main SMC absorption. We used the method of profile fitting to estimate column densities, line widths, and velocities for each of the individual components discernible in the ensemble of absorption-line profiles.

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In the fits, we assumed a Gaussian instrumental profile, with FWHM p 2.7 km s⫺1, and we used the rest wavelengths and f-values tabulated by Morton (1991; as updated in Welty et al. 1999b). We performed a simultaneous fit to all the profiles shown in Figure 1 (except for the weak Ti ii l1910 and strong Fe ii l2586 lines)—using the same number of components, with the same relative velocities, for all species but allowing for differences in overall zero point. The individual component column densities were all allowed to vary in the fits (except for the strongest components in Si ii); the component line widths were fixed at b p 2.5–3.5 km s⫺1. The 11 SMC components found between 110 and 180 km s⫺1 are shown by tick marks in Figure 1; in the next sections, we discuss the abundances found for the components at 134, 140, and 146 km s⫺1 (group A), at 158 km s⫺1 (group B), and at 165 and 170 km s⫺1 (group C).

3. COLUMN DENSITIES AND RELATIVE ABUNDANCES

Inspection of Figure 1 reveals striking differences between the profiles of Mg ii, Si ii, S ii, and Zn ii (left) and those of Cr ii, Mn ii, Fe ii, and Ni ii (right). For the species on the left, the component groups A and B are of comparable strength; group C is somewhat stronger. For the species on the right, however, groups B and C are comparable, and group A is stronger. In Tables 1 and 2, we list the column densities derived from the profile fits and the corresponding abundance ratios (relative to solar values) [X/Zn] and [Zn/H], respectively. Comparisons of the overall SMC N(S ii) and N(Zn ii) (which should be at most mildly depleted) with the corresponding total N(H) and N(O i) (from the weak l1355 line) yield [Zn/H] p ⫺1.05 dex, [S/H] p ⫺0.89 dex, [O/H] p ⫺1.01 dex, and [O/Zn] p 0.04 dex. These relative abundances suggest that the strongest SMC components are primarily neutral, with slight depletions (∼⫺0.2 dex) of Zn and O and perhaps a slightly lower metallicity (∼⫺0.85 dex) than is typical for the SMC [although the uncertainties in N(H) would permit the typical ⫺0.6 to ⫺0.7 dex]. The various X ii listed in Table 1 therefore should be the dominant ions in SMC groups A, B, and C. The abundance ratios [X/Zn] listed in Table 2 for each of those three component groups are plotted in Figure 2 (left), together with representative abundance patterns seen in the Galactic ISM. Tables 1 and 2 and Figure 2 (right) also show integrated abundances for the SMC gas toward Sk 155, Sk 78 (HD 5980), and Sk 108 (all observed at relatively high resolution).

TABLE 1 SMC Column Densitiesa Components and Starsb A (134, 140, 146) . . . . . . . B (158) . . . . . . . . . . . . . . . . . . . C (165, 170) . . . . . . . . . . . . . Sk 155 (all) . . . . . . . . . . . . . . Sk 78 (all)c . . . . . . . . . . . . . . . Sk 108 (all)d . . . . . . . . . . . . . Solar abundancee . . . . . . a

N(Mg ii) (cm⫺2)

N(Si ii) (cm⫺2)

15.43⫹0.05 ⫺0.05 15.35⫹0.04 ⫺0.05 15.65⫹0.04 ⫺0.04 16.01⫹0.02 ⫺0.03 15.64⫹0.09 ⫺0.11

15.53⫹0.05 ⫺0.06 15.60⫹0.14 ⫺0.20 15.86⫹0.15 ⫺0.23 16.22⫹0.07 ⫺0.09 15.69⫹0.02 ⫺0.02 15.52⫹0.03 ⫺0.03 7.55

… 7.58

N(S ii) (cm⫺2) 15.37⫹0.07 ⫺0.09 15.19⫹0.11 ⫺0.13 15.54⫹0.09 ⫺0.12 15.92⫹0.05 ⫺0.06 15.51⫹0.01 ⫺0.01 ⲏ15.29 7.27

N(Ti ii) (cm⫺2)

N(Cr ii) (cm⫺2)

N(Mn ii) (cm⫺2)

N(Fe ii) (cm⫺2)

N(Ni ii) (cm⫺2)

N(Zn ii) (cm⫺2)

12.67⫹0.07 ⫺0.09 12.09⫹0.14 ⫺0.21 12.27⫹0.15 ⫺0.22 13.04⫹0.08 ⫺0.09 12.38⫹0.04 ⫺0.04

13.15⫹0.03 ⫺0.04 12.30⫹0.08 ⫺0.10 12.48⫹0.08 ⫺0.10 13.38⫹0.02 ⫺0.03 13.26⫹0.15 ⫺0.26 13.08⫹0.01 ⫺0.02 5.68

12.67⫹0.04 ⫺0.03 12.04⫹0.06 ⫺0.08 12.15⫹0.06 ⫺0.07 12.98⫹0.03 ⫺0.03 12.90⫹0.10 ⫺0.12 12.83⫹0.03 ⫺0.04 5.53

14.84⫹0.03 ⫺0.03 14.07⫹0.06 ⫺0.07 13.98⫹0.09 ⫺0.10 15.08⫹0.03 ⫺0.02 14.94⫹0.02 ⫺0.02 14.84⫹0.03 ⫺0.03 7.51

13.39⫹0.03 ⫺0.04 12.59⫹0.10 ⫺0.13 12.57⫹0.11 ⫺0.16 13.63⫹0.03 ⫺0.03 13.56⫹0.02 ⫺0.02 13.34⫹0.09 ⫺0.11 6.25

12.56⫹0.03 ⫺0.03 12.47⫹0.04 ⫺0.04 12.74⫹0.03 ⫺0.04 13.14⫹0.02 ⫺0.02 12.84⫹0.02 ⫺0.03 12.50⫹0.07 ⫺0.09 4.65

… 4.93

Uncertainties are 1 j (photon noise and continuum uncertainty); values are logarithmic. Numbers in parentheses indicate components included (in units of kilometers per second). Values for Mg, S, Fe, and Ni from new fits to STIS spectra described by Koenigsberger et al. 2001; values for Si, Cr, Mn, and Zn from fits to IUE spectra using component structures obtained from STIS spectra; and value for Ti from fit to equivalent widths reported by Roth, Blades, & Albert 1995. d Welty et al. 1997 (updated for newer f-values and meteoritic reference abundances); Mallouris et al. 2001 (S only). e Anders & Grevesse 1989; logarithmic, with H p 12.0. b c

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TABLE 2 SMC Relative Abundancesa Components and Starsb A (134, 140, 146) . . . . . . B (158) . . . . . . . . . . . . . . . . . C (165, 170) . . . . . . . . . . . . Sk 155 (all) . . . . . . . . . . . . . Sk 78 (all)d . . . . . . . . . . . . . Sk 108 (all)e . . . . . . . . . . . .

[Mg/Zn] ⫺0.06 ⫺0.05 ⫺0.02 ⫺0.06 ⫺0.13

⫹0.05 ⫺0.06 ⫹0.06 ⫺0.07 ⫹0.05 ⫺0.05 ⫹0.02 ⫺0.02 ⫹0.10 ⫺0.11

…

[Si/Zn] ⫹0.07 ⫹0.22 ⫹0.22 ⫹0.18 ⫺0.05 ⫹0.12

⫹0.06 ⫺0.07 ⫹0.14 ⫺0.21 ⫹0.16 ⫺0.23 ⫹0.07 ⫺0.08 ⫹0.03 ⫺0.03 ⫹0.07 ⫺0.09

[S/Zn] ⫹0.19 ⫹0.10 ⫹0.18 ⫹0.16 ⫹0.05

⫹0.07 ⫺0.10 ⫹0.11 ⫺0.14 ⫹0.09 ⫺0.13 ⫹0.06 ⫺0.06 ⫹0.02 ⫺0.03

[Ti/Zn] ⫺0.17 ⫺0.66 ⫺0.75 ⫺0.38 ⫺0.74

ⲏ0.17

…

⫹0.08 ⫺0.10 ⫹0.14 ⫺0.22 ⫹0.15 ⫺0.23 ⫹0.07 ⫺0.09 ⫹0.05 ⫺0.05

[Cr/Zn] ⫺0.45 ⫺1.20 ⫺1.29 ⫺0.79 ⫺0.61 ⫺0.46

⫹0.04 ⫺0.05 ⫹0.09 ⫺0.11 ⫹0.08 ⫺0.10 ⫹0.03 ⫺0.03 ⫹0.15 ⫺0.26 ⫹0.07 ⫺0.08

[Mn/Zn] ⫺0.77 ⫺1.31 ⫺1.47 ⫺1.04 ⫺0.82 ⫺0.55

⫹0.05 ⫺0.05 ⫹0.07 ⫺0.09 ⫹0.07 ⫺0.07 ⫹0.03 ⫺0.04 ⫹0.10 ⫺0.13 ⫹0.07 ⫺0.09

[Fe/Zn] ⫺0.58 ⫺1.26 ⫺1.64 ⫺0.92 ⫺0.76 ⫺0.52

⫹0.04 ⫺0.05 ⫹0.07 ⫺0.08 ⫹0.09 ⫺0.11 ⫹0.03 ⫺0.04 ⫹0.03 ⫺0.03 ⫹0.07 ⫺0.09

[Ni/Zn] ⫺0.77 ⫺1.48 ⫺1.77 ⫺1.11 ⫺0.88 ⫺0.76

⫹0.05 ⫺0.05 ⫹0.10 ⫺0.13 ⫹0.11 ⫺0.17 ⫹0.03 ⫺0.04 ⫹0.03 ⫺0.04 ⫹0.10 ⫺0.14

[Zn/H]c … … … ⫺1.05⫹0.15 ⫺0.22 ⫺0.74⫹0.05 ⫺0.06 ⫺0.69⫹0.13 ⫺0.17

a

Uncertainties are 1 j (photon noise and continuum uncertainty); values are logarithmic. b Numbers in parentheses indicate components included (in units of kilometers per second). c ⫹0.05 ⫹0.11 Using log N(H) p 21.54⫹0.15 ⫺0.22 for Sk 155, 20.93⫺0.05 for Sk 78, and 20.54⫺0.14 for Sk 108 (Fitzpatrick 1985). d Values for Mg, S, Fe, and Ni from new fits to STIS spectra described by Koenigsberger et al. 2001; values for Si, Cr, Mn, and Zn from fits to IUE spectra using component structures obtained from STIS spectra; and value for Ti from fit to equivalent widths reported by Roth et al. 1995. e Welty et al. 1997 (updated for newer f-values and meteoritic reference abundances); Mallouris et al. 2001 (S only).

4. DISCUSSION

4.1. Abundances and Depletions In principle, observed gas-phase interstellar abundances reflect the combined effects of the nucleosynthetic history of the material and any depletion into dust grains. In the local Galactic ISM, observed subsolar abundances of elements such as Mg, Si, Fe, and Ni are usually taken to indicate that these more refractory elements are partly (or largely) locked in the dust, under the assumptions that the ISM is both well mixed and characterized by the abundance pattern found in the solar photosphere and in meteorites. Considering abundances relative to the typically undepleted S or Zn (rather than to H) minimizes uncertainties as to the overall (gas ⫹ dust) abundances in the ISM, which may be slightly subsolar (Sofia, Cardelli, & Savage 1994; Snow & Witt 1996). Several representative depletion patterns have been associated with different environments in the Galactic ISM (Jenkins 1987; Savage & Sembach 1996; Welty et al. 1997 and references therein; see Fig. 2): the “cold, dense cloud” (C) pattern, for which [Mg, Si/Zn] ∼ ⫺0.8 to ⫺0.9 dex and [Fe, Ni/Zn] ∼ ⫺1.8 dex; the “warm, diffuse cloud” (W) pattern, for which [Mg, Si/Zn] ∼ ⫺0.2 to ⫺0.4 dex and [Fe, Ni/Zn] ∼ ⫺1.2 dex; and the “halo cloud” (H) pattern, for which [Mg, Si/Zn] ∼ ⫺0.2 dex and [Fe, Ni/Zn] ∼ ⫺0.5 dex. While those three depletion patterns may only be representative of a continuum of patterns, the depletions of mod-

erately depleted elements such as Mg and Si generally do seem to be monotonically related to the depletions of the more severely depleted elements such as Fe and Ni. Toward Sk 108 [with much lower SMC E(B⫺V) and N(H)], all the SMC components/clouds seem to be characterized by very mild depletions, broadly similar to those found for Galactic halo clouds (Welty et al. 1997; Mallouris et al. 2001). Toward Sk 155, however, there are significant component-to-component differences—the overall depletions of Fe and Ni in the SMC range from mild ([Fe, Ni/Zn] ∼ ⫺0.3 to ⫺0.8 dex in group A and some of the lower column density components) to severe ([Fe, Ni/Zn] ∼ ⫺1.7 dex in group C) (see Table 2 and the left panel of Fig. 2). Moreover, the pattern of gas-phase abundances in the higher column density clouds toward Sk 155 does not agree with any of the patterns found in our Galaxy. Both Mg and Si appear to be essentially undepleted in all SMC components—even in groups B and C, where Fe and Ni are moderately to severely depleted; Ti is less severely depleted than would have been expected from [Fe, Ni/Zn], especially in group C. Even for the lowest N(Si ii) consistent with the l1808 profile, [Si/Zn] ∼ 0.0 dex and [Si/S] ∼ ⫺0.2 dex in groups B and C—so that Si would be depleted by at most 0.2 dex. This apparent lack of significant depletion for Mg and/or Si may be characteristic of the SMC, as similar [Mg, Si/Zn] are also found for the SMC gas toward Sk 108 and Sk 78, both located in the main SMC “bar” (see Table 2 and the right panel of Fig. 2).

Fig. 2.—Gas-phase abundance ratios [X/Zn] (relative to solar ratios) for the three strongest SMC component groups toward Sk 155 (left) and for all SMC components toward Sk 155, Sk 78, and Sk 108 (right), compared with typical values found for Galactic cold, dense clouds (C), warm, diffuse clouds (W), and halo clouds (H). Uncertainties (plotted only if 10.05 dex) are 1 j (Table 2). Since Zn is typically only mildly depleted, these relative abundances reflect primarily the depletions of elements X. In the SMC gas toward Sk 155, significant component-to-component variations are seen for the typically more heavily depleted elements Cr, Mn, Fe, and Ni. The elements Mg, Si, and S appear to be essentially undepleted in all three Sk 155 component groups and in all three lines of sight—even when Fe (for example) is heavily depleted. Such depletion patterns have not been seen in the Galactic ISM.

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It is unlikely that the differences in gas-phase abundance patterns are a result of overall (gas ⫹ dust) abundance differences between the SMC (or SMC wing) and our Galaxy. While the SMC is characterized by a lower overall metallicity, the relative elemental abundances determined for SMC stars and nebulae are quite similar to those found for analogous objects in our Galaxy (Russell & Dopita 1992; Welty et al. 1997 and references therein; Venn 1999)—at least for the elements considered here. In particular, the significant enhancement of the a-particle elements relative to the “Fe-peak” elements, found for Galactic stars of comparable metallicity, is not seen in the SMC. No significant differences in either overall abundances or abundance ratios are seen for four supergiants in the wing region analyzed by Luck & Lambert (1992) and Venn (1999). The [O/H] ratios found for the wing H ii regions N83 and N84 also are quite consistent with those found throughout the rest of the SMC (Pagel et al. 1978). Finally, for all the SMC components toward Sk 155, [S/Zn] ∼ 0.1–0.2 dex—consistent with the ratios seen in diffuse clouds in the Galactic disk and halo. The different gas-phase abundance patterns observed in the various components toward Sk 155 therefore appear to be due to differences in depletions, as assumed above. 4.2. Dust Composition The apparent differences in depletions between the SMC and the Milky Way imply corresponding (and significant) differences in dust composition. Most models of interstellar dust grains assume some combination of carbonaceous grains (for ˚ extinction bump), with the size distribution exthe 2175 A tended into the molecular regime (to account for the observed near- and mid-IR emission), and silicate grains (e.g., Mathis, Rumpl, & Nordsieck 1977; Draine & Lee 1984; Li & Draine 2001). Because three of the four well-determined SMC ex˚ tinction curves exhibit a steep far-UV rise with no 2175 A bump (Gordon & Clayton 1998), silicates have figured more prominently in models of SMC dust (Pei 1992; Zubko 1999; Weingartner & Draine 2001). It is generally assumed that all the available SMC Si is in the dust and that silicates account for most of the extinction. While IR emission attributed to silicate dust has been detected in the SMC star-forming region N66 (near Sk 78; Contursi et al. 2000), Si is depleted by no more than about 0.2 dex in the three SMC sight lines with accurately determined interstellar abundances (assuming [Zn/H] ⲏ ⫺0.2 dex). If Si is generally very mildly depleted (at most) in the SMC, the dust models will have to be reassessed. From comparisons of the abundances of the most abundant elements in Galactic dust, Sofia et al. (1994) concluded that “oxides [e.g., Fe2O3, Fe3O4, MgO] and/or metallic Fe must be a substantial fraction of the grain core population.” The combination of minimal Si depletion and moderate to severe Fe depletion in the SMC gas toward Sk 155 (and Sk 78) suggests that such metallic grains might be even more prominent in the

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SMC. Wayte (1990) associated the 170 km s⫺1 component (our group C)—seen in Ca ii absorption toward Sk 151 and Sk 155 but not toward other stars in that region—with the stronger polarization measured toward those two stars. Toward Sk 155, group C has both the highest N(H) [inferred from N(Zn ii) and N(S ii)] and the most severe depletions of Cr, Mn, Fe, and Ni—and thus (presumably) the largest population of metallic grains. 4.3. QSO Absorption-Line Systems There are several intriguing similarities between properties of the SMC ISM and those observed or inferred for starburst galaxies and/or for QSO absorption-line systems (ALSs). On average, the QSOALSs have metallicities roughly 1/10 solar (although a wide range has been observed; see, e.g., Pettini et al. 1997; Lu et al. 1996; Prochaska & Wolfe 2000); on average, the relative gas-phase abundances in the QSOALSs are similar to those seen for SMC clouds (and Galactic halo clouds) characterized by mild depletions (Welty et al. 1997). Moreover, SMC-like extinction has been inferred both for QSOALSs (Pei, Fall, & Bechtold 1991) and for starburst galaxies (Gordon & Clayton 1998)—suggesting that the physical environments and/ or the dust composition may be similar in all three cases. If the dust in the QSOALSs contains little Si, as seems to be the case for at least some of the dust in the SMC, then the implications of various QSOALS abundance ratios involving Si would need to be reconsidered: (1) If Si is not depleted, then it can be used as a metallicity indicator. The advantages are that Si can be observed more easily than Zn for systems with lower N(H) and for systems at higher redshift. For 21 observed QSOALSs (from the literature as of 2000 December), [Si/Zn] ∼ ⫺0.07 Ⳳ 0.18 dex; however, some systems clearly have nonzero [Si/Zn]. (2) From that same sample, the average [Fe/Zn] ∼ ⫺0.54 Ⳳ 0.29 dex (29 systems) is suggestive of depletions similar to those seen for Galactic halo clouds. The corresponding [Si/Fe] ∼ ⫹0.43 Ⳳ 0.18 dex (35 systems), however, is slightly larger than the value of ⫹0.3 dex expected for halo cloud depletions—which could be taken to imply an overabundance of a-particle elements relative to Fe-peak elements. Conversely, that [Si/Fe] would also be consistent with no depletion of Si (and no enhancement of a-particle abundances). Disentangling the effects of depletion and nucleosynthetic history in the QSOALSs (Lauroesch et al. 1996; Kulkarni, Fall, & Truran 1997; Vladilo 1998) will require abundances for a variety of elements—especially if the depletion patterns are different for different systems. We thank Kathy Roth for providing Ti ii equivalent widths for Sk 78. Support for this work has been provided by NASA through grant HST-GO-08145.01-A (administered by STScI) and through LTSA grant NAG5-3228 (both to the University of Chicago).

REFERENCES Anders, E., & Grevesse, N. 1989, Geochim. Cosmochim. Acta, 53, 197 Ardeberg, A., & Maurice, E. 1977, A&AS, 30, 261 Contursi, A., et al. 2000, A&A, 362, 310 Draine, B. T., & Lee, H. M. 1984, ApJ, 285, 89 Fitzpatrick, E. L. 1985, ApJS, 59, 77 Gordon, K. D., & Clayton, G. C. 1998, ApJ, 500, 816 Howk, J. C., & Sembach, K. R. 2000, AJ, 119, 2481 Jenkins, E. B. 1987, in Interstellar Processes, ed. D. J. Hollenbach & H. A. Thronson (Dordrecht: Reidel), 533 Koenigsberger, G., et al. 2001, AJ, 121, 267

Kulkarni, V. P., Fall, S. M., & Truran, J. W. 1997, ApJ, 484, L7 Lauroesch, J. T., Truran, J. W., Welty, D. E., & York, D. G. 1996, PASP, 108, 641 Lehner, N., Sembach, K. R., Dufton, P. L., Rolleston, W. R. J., & Keenan, F. P. 2001, ApJ, 551, 781 Lennon, D. J. 1997, A&A, 317, 871 Li, A., & Draine, B. T. 2001, ApJ, in press (astro-ph/0011319) Lu, L., Sargent, W. L. W., Barlow, T. A., Churchill, C. W., & Vogt, S. S. 1996, ApJS, 107, 475 Luck, R. E., & Lambert, D. L. 1992, ApJS, 79, 303

No. 1, 2001

WELTY ET AL.

Mallouris, C., et al. 2001, ApJ, in press (astro-ph/0102279) Mathis, J. S., Rumpl, W., & Nordsieck, K. H. 1977, ApJ, 217, 425 Morton, D. C. 1991, ApJS, 77, 119 Pagel, B. E. J., Edmunds, M. G., Fosbury, R. A. E., & Webster, B. L. 1978, MNRAS, 184, 569 Pei, Y. C. 1992, ApJ, 395, 130 Pei, Y. C., Fall, S. M., & Bechtold, J. 1991, ApJ, 378, 6 Pettini, M., King, D. L., Smith, L. J., & Hunstead, R. W. 1997, ApJ, 478, 536 Prochaska, J. X., & Wolfe, A. M. 2000, ApJ, 533, L5 Roth, K. C., Blades, J. C., & Albert, C. E. 1995, BAAS, 186, 35.02 Russell, S. C., & Dopita, M. A. 1992, ApJ, 384, 508 Savage, B. D., & Sembach, K. R. 1996, ARA&A, 34, 279 Snow, T. P., & Witt, A. N. 1996, ApJ, 468, L65

L79

Sofia, U. J., Cardelli, J. A., & Savage, B. D. 1994, ApJ, 430, 650 Venn, K. A. 1999, ApJ, 518, 405 Vladilo, G. 1998, ApJ, 493, 583 Wayte, S. R. 1990, ApJ, 355, 473 Weingartner, J. C., & Draine, B. T. 2001, ApJ, 548, 296 Welty, D. E., Frisch, P. C., Sonneborn, G., & York, D. G. 1999a, ApJ, 512, 636 Welty, D. E., Hobbs, L. M., Lauroesch, J. T., Morton, D. C., Spitzer, L., & York, D. G. 1999b, ApJS, 124, 465 Welty, D. E., Lauroesch, J. T., Blades, J. C., Hobbs, L. M., & York, D. G. 1997, ApJ, 489, 672 Zubko, V. G. 1999, ApJ, 513, L29