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A&A 517, A45 (2010) DOI: 10.1051/0004-6361/200912817
c ESO 2010

Astronomy & Astrophysics

From interstellar abundances to grain composition: the major dust constituents Mg, Si, and Fe
N. V. Voshchinnikov1
1 2 3

,2,3

and Th. Henning1

Max-Planck-Institut fÝr Astronomie, KÆnigstuhl 17, 69117 Heidelberg, Germany e-mail: voshchinnikov@mpia.de Sobolev Astronomical Institute, St. Petersburg University, Universitetskii prosp. 28, 198504 St. Petersburg, Russia e-mail: nvv@astro.spbu.ru Isaac Newton Institute of Chile, St. Petersburg Branch, Russia

Received 3 July 2009 / Accepted 6 April 2010
ABSTRACT

We analyse observational correlations for three elements entering into the composition of interstellar silicate and oxide grains. Using current solar abundances, we converted the gas phase abundances into dust phase abundances for 196 sightlines. We deduce a sharp difference in abundances for sightlines located at low (|b| < 30 ) and high (|b| > 30 ) galactic latitudes. For high-latitude stars, the ratios Mg/Si and Fe/Si in dust are close to 1.5. For disk stars they are reduced to Mg/Si 1.2and Fe/Si 1.05. The derived numbers indicate that 1) the dust grains cannot be the mixture of silicates with olivine and pyroxene composition only, and some amount of magnesium or iron (or both) should be in another population and that 2) the destruction of Mg-rich grains in the warm medium is more effective than for Fe-rich grains. We reveal a decrease in dust phase abundances and correspondingly an increase in gas phase abundances with distance D for stars with D > 400 pc. We attribute this to an observational selection effect: a systematic trend toward lower observed hydrogen column density for distant stars. We find differences in abundances for disk stars with low (E ( B - V ) < 0.2) and high (E ( B - V ) > 0.2) reddenings that reflect the distinction between the sightlines passing through diffuse and translucent interstellar clouds. For Scorpius-Ophiuchus, we detect a uniform increase in dust phase abundances of Mg and Si with an increase in the ratio of total to selective extinction RV and a decrease in the strength of the far-UV extinction. This is the first evidence of growth of Mg-Si grains due to accretion in the interstellar medium.
Key words. ISM: abundances ­ dust, extinction

1. Introduction
Interstellar space is filled with gas and dust. Both components interact with each other. Atoms and molecules collide with solid particles and cause grain growth or destruction (sputtering). It depends on the relative velocity. An examination of these processes is based on analysis of the observed gas phase abundances. Spectroscopic studies of interstellar UV absorption lines started in the 1970s have revealed a deficit of heavy elements in the ISM in comparison with cosmic (solar reference) abundances (Spitzer & Jenkins 1975). The missing atoms were assumed to be tied up in solid particles that opened an indirect way to investigate the element composition of interstellar dust. Modelling of interstellar extinction (e.g., Li & Greenberg 1997; Zubko et al. 2004; Voshchinnikov et al. 2006) has demonstrated that cosmic abundance constraints might be crucial in deciding on modern dust models. Cosmic abundances of heavy elements obtained from spectroscopic studies of ordinary stars (Snow & Witt 1996; Przybilla et al. 2008) and a decrease in the estimates of metal abundances in the solar atmosphere over the past years (Asplund et al. 2005, 2009) essentially limited the number of atoms incorporated into dust particles. In this situation, accurate determination and analysis of gas phase abundances are especially important.
Table A.1 is only available in electronic form at http://www.aanda.org

A quantitative theory of element depletions is lacking. First phenomenological models showed a possible dependence of depletions on the element equilibrium condensation temperature (Field 1974) and the first or second element ionization potential (Snow 1973; Tabak 1979). The dependence of gas phase abundances on hydrogen column density, fraction of molecular hydrogen, distance, location, etc., were also considered (Tarafdar et al. 1983;Harris etal. 1984; Jenkins et al. 1986; see also references in Jensen 2007; and Jenkins 2009). A fresh approach to the problem of gas phase abundances has been devised by Jenkins (2004, 2009, hereafter J09) who investigated general patterns in depletions of 17 elements. He finds that the propensity of an element X to convert from gas to solid phase can be described by a linear equation with coefficient AX . The values of AX are assumed to be the same for all sightlines, while the individual sightlines can be characterized by a depletion factor F which is common to all elements. This means that all elements are depleted in unison independently of local physical conditions. In this paper, we exploit a more traditional approach by trying to keep individual features of separate sightlines. We investigate correlations in abundances of three elements Mg, Si, and Fe, which are classified as major dust constituents (Jones 2000). These elements along with the primary element O can be incorporated in solid phase in the form of Mg-Fe silicates, metal particles, or oxides. We outline oxygen abundances, which will be fully discussed elsewhere. Another primary element, C, Page 1 of 15

Article published by EDP Sciences


A&A 517, A45 (2010)

50 30

Fe (135) Mg (149) Si (40) O (120)

b, deg.

10 -10 -30 -50 180 90 0 270 180

l, deg.
Fig. 1. Galactic distribution of stars studied in this paper. Sightlines with measured Fe, Mg, Si, and O are shown by different symbols. Number of stars considered is indicated in parentheses in the legend.

cannot be studied in such detail since the number of sightlines with measured carbon abundances does not exceed 20 (Sofia et al. 2004; J09), and at least in six directions they have been revised downward in comparison with earlier estimates (Sofia & Parvahti 2009). Our main goal is to establish what could be the real dust phase abundances and whether they could rule out ambiguity in modelling interstellar extinction, polarization, and spectral IR features.

2.2. Data

2. Sample of stars and first analysis
2.1. Definitions

The abundance of an element in the interstellar medium is determined as a number of atoms relative to that of hydrogen, [X/H], where X (or N (X)) and H (or N (H) = N (HI) + 2 N (H2 )) are the column densities of an element X and hydrogen in a given direction. The abundances by number are often expressed as the number of X atoms per 106 hydrogen nuclei (parts per million, ppm, hereafter). Usually, the gas phase abundances of most elements [X/H]g are smaller than the corresponding "cosmic" (reference, solar) abundances. The depletion of an element X is defined by DX = X H X H ·
cosmic

(1)

g

The logarithmic quantities X = log DX = log are also used1 .
The bracketed notation [X/H] and the units ppm are traditionally utilized when the dust phase abundances are studied (e.g., Greenberg 1978;Mathis 1996; Voshchinnikov 2004).
1

X H

- log
g

X H

(2)
cosmic

We assume that interstellar atoms are in the dominant ionization stage for HI regions: OI, MgII, SiII, and FeII. The contribution of neutral atoms to the total column density of magnesium, silicon, and iron can be neglected (see, e.g., data of Savage & Bohlin ´ 1979; and Gnacinski & Krogulec 2006). The presence of the HII regions on the line of sight may lead to some fraction of atoms in a stage of ionization above the preferred one. To exclude this effect, J09 deduces his fits for the sightlines with N (H) > 3 â 1019 cm-2 , although later he finds small departure from a linear trend between the depletion factors F and the logarithm of the average density for sightlines with hydrogen column densities lower than 3 â 1019 cm-2 . Our list of stars includes different targets with measured gas phase abundances of oxygen, magnesium, silicon, or iron. We transform all data into the unified (standard) system of oscillator strengths as given in Table 1 of J09. We also use the hydrogen column density from J09 when available2 . The final sample contains 196 sightlines with 1 errors (see Table A.1 in Appendix). Observational data are taken from J09, Cartledge et al. (2004, 2008), and Jensen et al. (2005) for oxygen (120 sightlines); J09, Jensen & Snow (2007b), Cartledge et al. (2006), Howk et al. ´ (1999), and Gnacinski & Krogulec (2006) for magnesium (149); ´ Jensen (2007), Gnacinski & Krogulec (2006), and J09 for silicon (40); and J09, Jensen & Snow (2007a), Snow et al. (2002), and Miller et al. (2007) for iron (135). For two sightlines (HD 93521 and HD 215733), where the data for separate velocity components are obtained, we take the total column densities. The galactic distribution of all 196 stars is plotted in Fig. 1. Overlapping symbols indicate that, for a given sightline measurements were made for more than one element. Figure 2 illustrates the distance distribution of stars in the projection on the galactic plane. The abundances were supplemented with stellar distances D, colour excesses E ( B - V ) and characteristics of the extinction
2

N (H) for star CPD ­59 2603 is taken from Jensen & Snow (2007a).

Page 2 of 15


N. V. Voshchinnikov and Th. Henning: From interstellar abundances to grain composition

F e (135) M g (149) S i (40) O (120)






10 10 0 1 00 0











Fig. 3. Distance distribution of colour excess E ( B - V ) for stars studied in this paper. Open squares show stars located at galactic latitudes |b| > 30 . polar reprerom the Sun Fe, Mg, Si, is indicated

Fig. 2. Distance distribution of stars studied in this paper in sentation. The Galactic centre is to the right. The distance f is plotted in logarithmic scale. Sightlines with measured and O are shown with different symbols. Number of stars in parentheses.

curves: the ratio of total to selective extinction RV = AV /E ( B - V ) and the parameters of UV extinction as suggested by Fitzpatrick & Massa (1990, 2007). Extinction data are collected from papers of Fitzpatrick & Massa (2007), Valencic et al. (2004), Wegner (2002, 2003), Patriarchi et al. (2003) and the papers with published abundances cited above. We recalculated all colour excesses E ( B - V ) published by J09 using normal colours of stars from Straizys (1992) and Landolt-BÆrnstein and spectral types from Bowen et al. (2008) and J09. This permits avoiding the difficulties with the negative values of E ( B - V ) obtained by J09. Using the 2MASS K -magnitudes from the Simbad database, correction to the Johnson system given by Bowen et al. (2008) and normal colours (V - K )0 from Straizys (1992) and Winkler (1997), we estimated colour excesses E (V - K ) for several stars. Next we found the ratio RV with the aid of the relation RV = 1.1 E (V - K )/E ( B - V ) (Voshchinnikov & Il'in 1987). From the distance distribution of colour excesses plotted in Fig. 3, it follows that the major part of stars have E ( B - V ) < 0.6. This means that they are located behind diffuse atomic and molecular clouds or translucent interstellar clouds (see Table 1 in Snow & McCall 2006, for classification). In Fig. 3 we separate 16 stars located at high galactic latitudes |b| > 30 3 . Previous considerations (e.g., Savage & Sembach 1996) indicate that the gas has smaller depletions in the lower halo in comparison with gas located in the galactic disk. The total hydrogen column density is smaller for stars observed at high galactic latitudes (Fig. 4). Figure 4 allows one to study the gas-to-dust ratio in the direction of stars from our sample. Ryter (1996) deduce the average gas-to-dust ratio N (H) = 6.56 â 10 E(B - V )
21

Fig. 4. Total hydrogen column density as function of the colour excess for stars studied in this paper. The dashed line shows the average dependence between N (H) and E ( B - V ) deduced by Ryter (1996).

atoms cm

-2

mag

-1

,

(3)

where the contribution of ionized hydrogen is neglected. From Fig. 4 it can be seen that the dependence of N (H) on E ( B - V ) does not strongly deviate from the average dependence if E ( B - V ) > 0.2. Near this value there is the border between dif fuse and translucent interstellar clouds (Snow & McCall 2006). Observational errors are larger if the sightline crosses a diffuse cloud. For stars seen through the translucent clouds, the ratio N (H)/E ( B - V ) lies within rather narrow limits from 3 â 1021 cm-2 mag-1 to 1 â 1022 cm-2 mag-1 . For almost all sightlines considered by us, the criterion of J09 (N (H) > 3 â 1019 cm-2 ) is satisfied. Two points in the lower left corner of Fig. 4 correspond to HD 34029 (Capella)4 and HD 48915
4 J09 gives the observed colour B - V not E ( B - V ) for Capella in Table 2.

3 We include HD 38666 (b = -27.1 ) where anomalous high gas phase abundances of Mg, Si, and Fe are observed in this list of stars.

Page 3 of 15


A&A 517, A45 (2010)

(Sirius). The colour excesses for them are very small and uncertain, therefore we exclude these two stars from further analysis.
smic

1.0

2.3. Reference abundances

0.9

We use, as a starting point, the solar abundances from Asplund et al. (2009) as cosmic abundances. These data are justified and are rather close to the modern solar system abundances recommended by Lodders et al. (2009). Taking [O/H] = 490 ppm, [Mg/H] = 39.8 ppm, [Si/H] = 32.4 ppm, and [Fe/H] = 31.6 ppm, and neglecting the errors in the reference abundances, we converted the gas phase abundances into the dust phase abundances X X X X ( 1 - DX ) = - = Hd H cosmic Hg H cosmic = X H 1 - 10
cosmic X

[Mg/H]d/[Mg/H]co

0.8

0.7

stars with |b|300, E(B-V)0.2 (48) stars with |b|300, E(B-V) >0.2 (85) stars with |b|>300 (14)

0.6

.

(4)

0.5 1.0

2.4. Choice of external parameter and selection effects

E ( B - V ) = AB - AV 1.086 Cext,B - Cext,V Nd , where Nd is the dust column density and Cext,B , Cext,V are average extinction cross sections. Sightlines with low reddening may be the result of the absence of dust (low Nd -value) or similar cross sections in the B and V bands. The latter can be interpreted as the presence of large grains producing neutral extinction. Perhaps, such sightlines have higher hydrogen column density than average and appear as points above the curve at the left upper part of Fig. 4. However, one should keep in mind that the error of the colour excess is larger for lower values of E ( B - V )5 . Relative dust phase abundances of Mg, Si, and Fe are plotted in Fig. 5 in the form [X/H]d /[X/H]cosmic = 1 - DX as a function of E ( B - V )/D. Using Eq. (3), the average reddening can be translated into average gas density: n(H) cm-
3

[Fe/H] /[Fe/H]cosmi

Previous considerations reveal correlations between element depletion and an "external" parameter characterizing the gas density in the line of sight: N (H), n(H) = N (H)/D, f (H2 ) = 2N (H2 )/N (H) (see Jensen 2007; and J09 for detailed discussion). Clear trends toward increasing gas depletion (and correspondingly growth of dust phase abundances) are found for Mg, Si, and Fe with respect to n(H) . Sometimes, variations in depletion with distance are also considered (Cartledge et al. 2006, 2008; Jensen 2007). In deciding on the parameter connecting dust grains and abundances, we are restricted by reddening E ( B - V ), extinction AV and average reddening E ( B - V )/D or extinction AV /D. In order to determine extinction, we must know the ratio RV , which is not easily detected from observations (see discussion in Straizys 1992). It is also important that we only know total (summary) extinction or reddening and cannot separate individual clouds on the line of sight. For further analysis, we choose the average reddening E ( B - V )/D as an external parameter because E ( B - V )and D are usually well known for different sightlines. Colour excess E ( B - V ) (or reddening) characterizes the amount of dust on the line of sight and the properties of dust particles. It can be calculated as

cosmic

0.8

[Si/H]d/[Si/H]

0.6

stars with |b|300, E(B-V)0.2 (14)

0.4

stars with |b|300, E(B-V)>0.2 (15) stars with |b|>300 (10)

1.00

c

0.95

0.90

d
0.85

stars with |b|300, E(B-V) 0.2 (47) stars with |b|300, E(B-V) >0.2 (78) stars with |b|>300 (9) supergiants

0.80 0.01 0.1 1

E(B-V)/D, mag/kpc
Fig. 5. Relative dust phase abundances of Mg, Si, and Fe with 1 error bars in dependence on average reddening E ( B - V )/D. Circles and squares show data for stars with |b| 30 and |b| > 30 , respectively. The number of stars is indicated in parentheses. Crosses in lower panel show data for supergiants.

2.13 â E ( B - V )/D[mag/kpc].

5 The polarimetric data obtained for stars with low reddening give the polarization efficiency close to average one P/E ( B - V ) = 3%/mag (Berdyugin et al., in prep.).

Page 4 of 15


N. V. Voshchinnikov and Th. Henning: From interstellar abundances to grain composition
1.00

[Fe/H] /[Fe/H]cosmi

c
0.95
HD 62542

d

stars with D400 pc stars with D>400 pc

0.90 0.01 0.1 1

Fig. 6. Distance dependence of total hydrogen column density for stars studied in this paper. The dashed line is the fit for stars with D > 400 pc.

E(B-V)/D, mag/kpc
Fig. 7. Relative dust phase abundances of Fe as a function of average reddening E ( B - V )/D for stars with |b| 30 and E ( B - V ) > 0.20. Rhombuses and crosses show data for stars with D 400 pc and D > 400 pc, respectively. The dashed line is the fit.

Stars in our sample are divided into two groups depending on their position in the Galaxy: stars at high galactic latitudes (|b| > 30 ) and stars located near the disk (|b| 30 ), which in turn are separated as sightlines passing through diffuse (E ( B - V ) 0.20) and translucent (E ( B - V ) > 0.20) clouds. It can be seen that the abundances for high-latitude and disk sightlines are quite different. For high-latitude stars, the element fraction in dust is lower and does not depend on the average reddening. This is well established and was interpreted in the homework of the Spitzer (1985) model of the different grain composition in the warm and cold phases of the ISM. For disk stars, independent of reddening, there is a smooth decrease in the dust phase abundances of Mg, Si, and Fe when E ( B - V )/D grows smaller. This is expected in the context of previous findings (Jensen 2007; J09). The reason for such behaviour may be some physical processes lower cosmic abundances outside the solar environment or observational selection. Selection effects can appear for distant stars as they may be more luminous and produce more UV flux that can modify gas abundances. To check this we mark out in Fig. 5 (lower panel) 19 supergiants. It is evident that the directions to supergiants are not peculiar and fall well into the general pattern of disk and high-latitude stars. Therefore, differences in the types of material probed along more and less luminous stars seem cannot cause the trends observed in Fig. 5. The true reason of decreasing dust phase abundances (and correspondingly increasing gas phase abundances) for smaller E ( B - V )/D and n(H) is evident from Fig. 6 (see also Fig. 3). A systematic tendency exists for a decrease in the observed hydrogen column density and colour excess for distant stars. In our sample (178 sightlines with |b| 30 ), the decrease is observed for stars with D > 400 pc with the gas-to-dust ratio re maining almost constant for clouds located at different distances. As a result of lower hydrogen column density for distant stars, we obtain a clear trend for abundances as a function of any parameter related to the gas or dust density or distance. Figure 7 illustrates this observational selection effect. It shows the ratio [Fe/H]d /[Fe/H]cosmic in dependence on E ( B - V )/D for stars observed through translucent interstellar clouds. Apparently, many

previous correlations of interstellar abundances reflect a decrease in N (H) and E ( B - V ) with increasing D. It looks like we observe less and less dense clouds when the distance grows. Evidently, this problem requires further investigation, especially observations of high-density longer sightlines.

3. Results and discussion
3.1. Mean values and deviations

Mean element abundances locked up in dust are given in Table 2. The data are presented for all sightlines and separately for highlatitude and disk stars having low and high reddening, respectively. Distinctions between the various groups are noticed. For all three elements the following inequality is valid: [X/H]d
|b|>30

< [X/H] < [X/H]

d |b|<30 , E(B-V )0.2 d |b|<30 , E(B-V )>0.2

.

As follows from Fig. 5, halo (z > 150 pc) and disk stars can be easily distinguished by the average reddening. For halo stars we have E ( B - V )/D < 0.05 mag/kpc. Exceptions are three short high-latitude sightlines for HD 116658 (b = +50.8, D = 80 pc), HD 203532 (b = -31.7, D = 211 pc), and HD 210121 (b = -44.4, D = 223 pc). They cross relatively dense clouds (see also Fig. 3) and have abundances similar to those of halo stars. The upper right corner of Fig. 5 is mainly filled by reddened nearby stars. A dividing line between stars with distances greater or less than 400 pc passes near the value E ( B - V )/D 1mag/kpc (Fig. 7). Because the data for stars with D > 400 pc seems to be "infected" with observational bias, we also calculated the mean abundances for nearby stars. A noticeable growth of abundances occurs for reddened disk stars. The results are shown in Table 2 in italics. There are several sightlines where the dust phase abundances of Mg and Fe are significantly lower than the general trends
Page 5 of 15


A&A 517, A45 (2010) Table 2. Mean value of element abundance in dust phase in ppm with 1 error. [Mg/H] all sightlines stars with |b| 30 E ( B - V ) 0.20 E ( B - V ) > 0.20 stars with |b| > 30 Notes.
()

d

[Mg/H]d [Mg/H]cosm

[Si/H]
ic

d

[Si/H]d [Si/H]cosm

[Fe/H]
ic

d

[Fe/H]d [Fe/H]cosm

ic

33.13 ± 2.35 32. 34. 36. 29. 31 11 90 98 ± ± ± ± 2. 1. 1. 3. 60 96 44 92

0.832 ± 0.059 (147) 0. 0. 0. 0. 812 857 927 753 ± ± ± ± 0. 0. 0. 0. 065 049 036 099 ( ( ( ( 48) 85) 14) 14)

25.01 ± 2.91 23. 29. 30. 20. 97 30 25 03 ± ± ± ± 3. 1. 0. 4. 03 48 76 90

0.772 ± 0.090 (39) 0. 0. 0. 0. 740 904 934 618 ± ± ± ± 0. 0. 0. 0. 093 046 023 151 ( ( ( ( 14) 15) 6) 10)

30.64 ± 0.41 30. 30. 31. 28. 65 86 16 68 ± ± ± ± 0. 0. 0. 1. 40 31 31 30

0.970 ± 0.013 (134) 0. 0. 0. 0. 970 977 986 908 ± ± ± ± 0. 0. 0. 0. 013 010 010 041 ( ( ( ( 47) 78) 17) 9)

Number of sightlines. Results for stars with D 400 pc and |b| 30 , E ( B - V ) > 0.20 are given in italics.

clearly seen in Fig. 5. The major part of "peculiar" sightlines is related to stars with E ( B - V ) 0.2. Only two objects (HD 62542 and HD 99890) are observed through translucent clouds, but the observational errors are quite large in these cases. Note also that the shape of the UV extinction curve in the direction of HD 62542 is very peculiar (Voshchinnikov & Das 2008). It should be remembered that almost all stars with anomalous low dust phase abundances are located in Carina-Centaurus at the galactic longitudes l 290 -330 . We compared our results with data presented in Table 4 of J09, which gives element depletion parameters corresponding to the cases of "full depletion" (F = 1) and "no depletion" (F = 0). We transformed these parameters into our reference system and find the relative dust phase abundances of Mg, Si, and Fe. They are equal to [Mg/H]d/[Mg/H]cosmic = 0.94 and 0.44, [Si/H]d /[Si/H]cosmic = 0.94 and 0.25, and [Fe/H]d /[Fe/H]cosmic = 0.99 and 0.88 for the cases F = 1and F = 0, respectively. The data for the "full depletion" case are well within our results (see Table 2 and Fig. 5). For another case ("no depletion"), the abundances of J09 seem to be too low even for high-latitude stars. This may be a result of observational selection and smaller number of sightlines in comparison with our sample.
3.2. Correlations

1.4

rcorr.=­0.73
HD 99890

d

[Fe/H]d/[Mg/H]

1.2

stars with |b|300, E(B-V)0.2 (39) stars with |b|300, E(B-V)>0.2 (57) stars with |b|>300 (8)

1.0

0.8

0.01

0.1

1

E(B-V)/D, mag/kpc
Fig. 8. Ratio of dust phase abundances of Fe and average reddening E ( B - V )/D. Open and filled disk stars with E ( B - V ) 0.20 and E ( B - V ) Squares correspond to sightlines with |b| > 30 . is the linear regression fit for stars with E ( B - [Mg/H]d = (-0.074 ± 0.010) log[E ( B - V )/D] + Mg in dependence on circles show data for > 0.20, respectively. The bold dashed line V ) > 0.20: [Fe/H]d / (0.830 ± 0.003).

Using our data it is possible to plot the dust phase abundance of one element against that of another element. Such diagrams clearly show the existence of strong correlations between the abundances of Mg, Si, and Fe for low reddened and distant stars (see, e.g., Cartledge et al. 2006; Miller et al. 2007). However, these correlations trace the behaviour of the hydrogen column density discussed in Sect. 2.4. To exclude the effect of NH we plot the ratio of the dust phase abundances of Fe to Mg in dependence on the average reddening. The result is shown in Fig. 8 for 104 sightlines with the linear regression fit for 56 disk sightlines with E ( B - V ) > 0.20 (HD 99890 was excluded) as derived by a 2 minimization that takes error bars into account. The Pearson correlation coefficient for these sightlines is rcorr = -0.73. As follows from Fig. 8, the amount of iron grows slightly in comparison with the amount of magnesium when the average reddening decreases. Twelve stars have distances D < 400 pc. They are located in the bottom right part of Fig. 8 and have an almost constant ratio Fe/Mg 0.84. A similar behaviour can be observed if we compare dust phase abundances of Mg or Fe and Si. However, in this case the number of sightlines is three times less, because it is dictated by the measurements of silicon (see Figs. 1 and 2). For almost all sightlines with measured Si we have measurements of Mg and Fe. Therefore, one can consider the composition of grains.
Page 6 of 15

3.3. Olivines, pyroxenes, and...?

All elements considered by us are constituents of cosmic silicate grains showing a pronounced 9.7 m feature observed in spectra of a wide variety of objects (see Henning 2009, for a recent review). The origin of this feature is related to the stretching of the Si-O bond in amorphous silicates with olivine (Mg2 x Fe2-2 x SiO4 ) or pyroxene (Mgy Fe1-y SiO3 ) stoichiometry, where 0 x,y 1. If we assume that Mg, Si, and Fe are incorporated only into Mg-Fe silicates, the ratio of (Fe+Mg)/Si in the dust phase must be in the range from 1 to 2, while the ratios Mg/Si and Fe/Si may vary from 0 to 2. Figure 9 shows the ratios as a function of average reddening for 31 sightlines where joint measurements of abundances of Mg, Si, and Fe are available. As follows from Fig. 9 (upper panel), the ratio (Fe + Mg)/Si > 2 for all stars. This ratio may exceed 3 for low-reddened and high-latitude stars. For high reddened stars (excluding CPD ­59 2603), the ratio (Fe + Mg)/Si lies in a narrow range from 2.15 to 2.35. The average value for 13 sightlines is (Fe + Mg)/Si = 2.25 ± 0.14.


N. V. Voshchinnikov and Th. Henning: From interstellar abundances to grain composition
4. 5

10
stars with |b|300, E(B-V)0.2 (4) stars with |b|300, E(B-V)>0.2 (13) stars with |b|>300 (1)

4. 0

H D 3 866 6

[M g + F e / H ] d /[ S i /H ]

d

8

[O/H]d/[Si/H]

3. 5
C P D - 5 9 26 03

d
6 4 2 0 0.01 0.1 1

3. 0

2. 5

2. 0

E(B-V)/D, mag/kpc

[M g / H ] d /[ S i / H ]

d

2. 0

C P D - 5 9 26 03

Fig. 10. Ratio of dust phase abundances of O and Si with 1 error bars reduced twice in dependence on average reddening E ( B - V )/D. Open and filled circles show data for disk stars seen through diffuse and translucent interstellar clouds, respectively. Squares correspond to sightlines with |b| > 30 .

1. 5

1. 0

stars with |b|300, E(B-V)0.2 (10) stars with |b|300, E(B-V)>0.2 (14) stars with |b|>300 (7)

2.0

[Fe/H]d/[Si/H]

The middle and low panels of Fig. 9 demonstrate that magnesium and iron are incorporated into dust particles in unequal parts for high and low reddened stars and high-latitude stars. For disk stars with E ( B - V ) > 0.20, the composition of grains averaged over 13 targets is Mg1.22 Fe1.04 SiOz1 . This composition cannot be reproduced by the mixture of olivine and pyroxene silicates alone with any value of x and y. It indicates that some amount of magnesium or iron (or both) should be embedded in another population of dust grains, probably, metal oxides. When we consider high-latitude stars, averaging over 6 targets (excluding HD 38666) gives the "grain composition" Mg1.50 Fe1.47 SiOz2 6 ; i.e., the ratio Fe/Mg is greater for highlatitude stars than for disk stars. This suggests that the destruction of Mg-rich grains in the warm medium is more effective than of Fe-rich grains.
3.4. ... + Problematic O

d

CPD -59 2603

1.5

1.0 0.01 0.1 1

E(B-V)/D, mag/kpc
Fig. 9. Ratio of dust phase abundances of (Fe+Mg)/Si, Mg/Si, and Fe/Si with 1 error bars in dependence on average reddening for 31 sightlines with joint measurements of three elements. Open and filled circles show data for sightlines with diffuse and translucent interstellar clouds, respectively. Squares correspond to high-latitude stars.

The sample of stars discussed in Sect. 3.3 includes 20 sightlines where the abundances of OI have also been measured. Thus, we can compare abundances of four elements. We plot oxygen abundances in Figs. 10 and 11. The error bars are twice reduced in comparison with those observed in these figures. Unfortunately, the uncertainties in determining of the gas phase oxygen abundances are too large to allow any definitive answer about the trends in oxygen depletion. Figure 9 clearly shows the excess of both iron and magnesium over silicon, so if we assume that all Si atoms are incorporated into olivine, the O to Si ratio must be equal to or exceed 4. This ratio is plotted in Fig. 10 for 18 sightlines7 . It is interesting
6 The values of z1, z2 are not the same and must lie between 3 and 4 for a mixture of pyroxene and olivine grains. 7 Two sightlines (towards HD 38666 and HD 141637) with [O/H]d < 0 are omitted.

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A&A 517, A45 (2010)
500
120 90 60

400

|b|<30 , E(B-V) 0.2 (21) |b|<300, E(B-V) >0.2 (82) stars with |b|>300 (4)
0

[O /H ]d 100 p p m ( 59) [O /H ]d > 100 p p m ( 48)
150 30

[O/H]d, ppm

300

200

180 100 100 0

0

100
210 330

0

-100 0.01 0.1 1

240 270

300

E(B-V)/D, mag/kpc
Fig. 11. Dust phase abundances of O in ppm with 1 error bars reduced in twice in dependence on average reddening E ( B - V )/D. Open and filled circles show data for disk stars seen through diffuse and translucent interstellar clouds, respectively. Squares correspond to sightlines with |b| > 30 .

Fig. 12. Distance distribution of stars with [O/H]d 100 ppm (open circles) and [O/H]d > 100 ppm (filled circles) in polar representation. The Galactic centre is to the right. The distance from the Sun is plotted in logarithmic scale. Number of stars is indicated in parentheses.

that, for 10 of 13 stars observed through translucent clouds, the ratio [O/H]d /[Si/H]d 4. This means that we have enough O for silicate particles and, perhaps, some additional oxygen for producing oxides and ices. However, this effect may be related to the regional variations in oxygen abundances as all these stars are located in the bottom part of Fig. 2, i.e., at the galactic longitudes l 180 -360. In Fig. 11 we plot the oxygen dust phase abundances for our sample of stars. These abundances were found as the difference between solar oxygen abundance (490 ppm) and observed gas phase abundances. The data are shown for 107 sightlines (12 sightlines where [O/H]d < -150 ppm were omitted). The absence of trends and correlations as seen for Mg, Si, or Fe (Fig. 5) is obvious. For many directions the dust phase abundances of O are negative and primarily related to large observational errors. We can make a tentative inference about the deficit (not excess!) of oxygen in the dust phase opposite to the conclusion of J09. In the case of using proto-Sun oxygen abundance ([O/H] = 575 ppm from Lodders 2003), things will not get much better, so it seems too early to search for the "missing oxygen" in the dust phase (Whittet 2010). We divided the stars into two groups: with [O/H]d > 100 ppm and [O/H]d 100 ppm. This border value is calculated as the product 4 â 25, assuming the mean value for Si from Table 2 and assuming that all Si atoms are tied up into olivinetype silicates8 . Such a division reveals an interesting galactic distribution of the "O-rich" and "O-poor" sightlines shown in Fig. 12. It can be seen that the open and filled circles are not well mixed. There are areas on the sky where the symbols of one type concentrate. Particularly striking is the zone between l 70 and l 140 (Cygnus, Cassiopea, Perseus) where 23 of 59 stars with reduced O abundance in dust are located. Perhaps, this behaviour
Models of dust evolution predict a dust phase abundance of oxygen at a level of about 130 ppm at the modern time (see, e.g., Fig. 16 in Zhukovska et al. 2008). Page 8 of 15
8

does not merely reflect the observational errors and indicate the existence of an another non-solar cosmic standard with enhanced metal abundances in this area. This hypothesis is partially supported by a higher fraction of metal-rich Cepheids found at these galactic longitudes (see Fig. 5 in Pedicelli et al. 2009).
3.5. Correlation with extinction and regional variations

The elements of C, O, together with Mg, Si, and Fe as studied here contribute to most of the mass of the interstellar dust. Therefore, it is interesting to study a dependence of dust phase abundances on the ratio RV . The ratio of total to selective extinction RV characterizes the visual extinction produced by dust particles with radii r > 0.05 m (e.g., Voshchinnikov 2004). The search for a correlation between depletions DO, Mg, Fe and RV has been attempted by Jensen (2007), who finds no correlation for Mg and Fe and a slight trend to increasing OI depletion with increasing RV . In our sample there are 164 stars with known or calculated values of RV . We found no significant correlation of dust phase abundances with RV either for the total sample or for the disk stars with low and high reddening. Very probably, the absence of correlation is a consequence of the mixture of short and long sightlines in different galactic directions. We investigated the regional variations of depletions and RV . The difference in iron depletion was found by Savage & Bohlin (1979) for sightlines in Cygnus and Scorpius-Ophiuchus, while Patriarchi et al. (2003) and Wegner (2003) discovered the difference in RV between stars in Cygnus and Carina and stars belonging to separate associations, respectively. To exclude the distance effects discussed in Sect. 2.4 we only considered reddened stars with D < 450 pc. From this subsample of 25 sightlines, we separated two groups of stars more or less closely located on the sky (see also Figs. 1 and 2 and Table A.1): four stars in Perseus (b = -13 Â -17, l = 160 -173, D 220-420 pc, N (H) = (1.26-1.95) â 1021 cm-2 , E ( B - V ) = 0.27-0.36) and seven stars in Scorpius-Ophiuchus


N. V. Voshchinnikov and Th. Henning: From interstellar abundances to grain composition
1.00

1.00

[Mg, Si, Fe/H]d/[Mg, Si, Fe/H]cosmic

0.95

[Mg, Si, Fe/H]d/[Mg, Si, Fe/H]cosmic

0.95

0.90
Mg Si Fe

0.90
Mg Si Fe

0.85 2.5 3.0 3.5 4.0 4.5

0.85 2.5 3.0 3.5 4.0 4.5

RV
Fig. 13. Relative dust phase abundances of Mg, Si, and Fe with 1 error bars in dependence on RV for stars located in Perseus. The data correspond to the stars (from left to right): HD 24398, HD 27778, HD 24912, and HD 23180. The values of RV were taken from Fitzpatrick & Massa (2007) for HD 23180 and HD 27778, Wegner (2003) for HD 24398 and Valencic et al. (2004) for HD 24912.

RV
Fig. 14. The same as in Fig. 13 but now for stars located in ScorpiusOphiuchus. The data correspond to the stars (from left to right): HD 144217, HD 143275, HD 147165, HD 144470, HD 147888, HD 148184, and HD 147933. The values of RV were taken from Fitzpatrick & Massa (2007) for HD 144470, HD 147165, HD 147888, and HD 147933 and Lewis et al. (2009) for HD 143275, HD 144217, and HD 148184.
1.00

(b = 17 -23 , l = 350 -358, D 120-200 pc, N (H) = (1.38-5.50) â 1021 cm-2 , E ( B - V ) = 0.21-0.51). Three of them (HD 147165, HD 147888, and HD 147933) belong to the Eastern Group of the Oph cloud, while four other stars belong to the Northern Group (Snow et al. 2008). The measured abundances of Mg, Si, and Fe are compared with RV in Figs. 13 and 14. For stars in Perseus, the value of RV is lower, on average, than the mean galactic value RV = 3.1. In this case the correlation between dust phase abundances and RV is absent (Fig. 13). For stars in Scorpius-Ophiuchus, RV varies from 2.6 to 4.4. Figure 14 shows a clear growth of silicon and magnesium abundances in dust with increasing RV 9 . The Pearson correlation coefficients are rcorr = 0.81 for Mg and rcorr = 0.86 for Si. This is the first evidence of a correlation of dust phase abundances with the ratio RV . Although the correlation coefficients are relatively large, one should keep in mind the small number of sightlines. Since RV is considered as a measure of grain size (Whittet 2003), we conclude that accretion of Si and Mg atoms on large grains takes place. Note also that obviously the abundance of iron is independent of the value of RV . For three stars in Perseus and four stars in ScorpiusOphiuchus the extinction curve in the far-UV is also known. Fitzpatrick & Massa (2007) fit the observed far-UV extinction using two parameters c4 and c5 , entering into a formula for the entire UV extinction. These parameters characterize the departure, in the far-UV, from the extrapolated bump-plus-linear components and indicate the strength of the far-UV curvature (c4 ) and the wavenumber in m-1 from which the far-UV extinction starts to grow (c5 ). We interpret parameter c4 as a measure of the relative amount of small grains (the slope of the size distribution function, e.g., the power index q in the powerlaw size distribution n(r) r-q ) and parameter c5 as a measure of the minimum particle size rmin in the dust ensemble.
9

[Mg, Si, Fe/H]d/[Mg, Si, Fe/H]

cosmic
0.95 0.90
Mg Si Fe

0.85 0.2 0.3 0.4 0.5 0.6

c

4

Fig. 15. Relative dust phase abundances of Mg, Si, and Fe with 1 error bars as a function of UV extinction curve parameter c4 (the strength of the far-UV curvature) for stars located in Scorpius-Ophiuchus. The data correspond to the stars (from left to right): HD 147933, HD 147888, HD 147165, and HD 144470.

A similar trend is seen if we replace RV by N (H).

Unfortunately, the limited data for Perseus do not allow a careful analysis. For stars located in Scorpius-Ophiuchus the normalized far-UV extinction is lower than the average Galactic extinction curve, which points to a deficit of the particles of small sizes. Figures 15 and 16 show the anticorrelation of the dust phase
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A&A 517, A45 (2010)
1.00

0.95

0.90
Mg Si Fe

0.85 5.5 6.0 6.5 7.0 7.5 8 .0

c5, m-1
Fig. 16. Relative dust phase abundances of Mg, Si, and Fe with 1 error bars as a function of UV extinction curve parameter c5 (wavenumber from which the far-UV extinction starts to grow) for stars located in Scorpius-Ophiuchus. The data correspond to the stars (from left to right): HD 147933, HD 147888, HD 144470, and HD 147165.

and pyroxene silicates. Some amount of magnesium or iron (or both) should be embedded into another population, probably oxides. 2. We reveal a clear distinction in abundances for nearby (D < 400 pc) and distant (D > 400 pc) stars: decrease in dust phase abundances and correspondingly increase in gas phase abundances with growth in D. We attribute this distinction to an observational selection effect: a systematic trend toward lower observed hydrogen column density for distant stars. As a result, we obtain a clear trend for abundances as a function of any parameter related to the gas or dust density or distance. 3. The less pronounced difference in abundances is found for disk stars with low (E ( B - V ) < 0.2) and high (E ( B - V ) > 0.2) reddenings. This reflects the distinction between sightlines passing through diffuse and translucent interstellar clouds. 4. Regional variations of abundances of Mg, Si, and Fe are not evident. However, for Scorpius-Ophiuchus, we established an uniform increase of dust phase abundances in Mg and Si with an increase in the ratio of total to selective extinction RV and an decrease in the strength of the far-UV extinction. Thus it is valid to say that there is a growth in Mg-Si grains due to accretion. The uncertainties in determing of the oxygen abundances are large and do not allow one to make definitive conclusions about the oxygen depletion. We can only indicate a possible regional peculiarity in the zone between l 70 and l 140 (Cygnus, Cassiopea, Perseus) where many stars with reduced O abundance in dust are located.

abundances of Mg and Si with parameters c4 and c5 (the correlation coefficients are from rcorr = -0.70 to rcorr = -0.99). The shift from right to left in these figures indicates the decrease in q and increase in rmin (see, e.g., Voshchinnikov & Il'in 1993; and Fig. 24 in Voshchinnikov 2004). This flattening of the size distribution function and growth of minimum grain size is accompanied by the increase in the dust phase fraction of Mg and Si; i.e., smaller grains are built up due to accretion of atoms from the gas. Apparently, this occurs in clouds with the hydrogen column density N (H) > 2 â 1021 cm-2 resulting in after the propagation of a low-velocity shock (Meyers et al. 1985). Our arguments in favour of grain growth by accretion are supported by the results of a detailed interpretation of extinction for two stars made by Das et al. (2010). They find that for silicate grains rmin = 0.07 m, q = 2.0 for HD 147933 and rmin = 0.04 m, q = 2.2 for HD 147165. The uniform variations in abundances of Mg and Si with RV , N (H), c4 , and c5 cannot be explained by grain coagulation because in this case the gas and dust phase abundances of elements are kept constant.

[Mg, Si, Fe/H]d/[Mg, Si, Fe/H]

cosmic

Acknowledgements. We thank Adam Jensen for sending data for silicon, Andrei Berdyugin for sending the polarimetric results, and Vladimir Il'in, ´ Jacek Krelowski, Piotr Gnacinski, Ralf Siebenmorgen, Svitlana Zhukovska, and Ted Snow for stimulating discussions. Special thanks go to the anonymous referee for helpful comments and suggestions. The work was partly supported by grants RFBR 07-02-00831, RFBR 10-02-00593a, NTP 2.1.1/665 and NSh 1318.2008.2.

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4. Conclusions
We investigated the differences in interstellar dust phase abundances of Mg, Si, and Fe entering into the composition of silicate and oxide grains. The distinctions in abundances can be separated into the following groups. 1. A sharp distinction in abundances is observed for sightlines located at low (|b| < 30 ) and high (|b| > 30 ) galactic latitudes. This is well known from previous studies. For highlatitude stars the ratios Mg/Si and Fe/Si in dust are close to 1.5. For disk stars these ratios are reduced to 1.2and 1.05 for Mg and Fe, respectively. The derived numbers indicate that the dust grains cannot be just a mixture of only olivine
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Appendix A
Table A.1. Dust phase abundances of Mg, Si, and Fe in ppm. N (1) 1 2 3 4 5 6 7 8 9 10 11 12 13 14 15 16 17 18 19 20 21 22 23 24 25 26 27 28 29 30 31 32 33 34 35 36 37 38 39 40 41 42 43 44 45 46 47 48 49 50 51 52 53 54 55 56 57 58 59 60 61 Star (2) HD 1383 HD 5394 HD 12323 HD 13268 HD 13745 HD 14434 HD 15137 HD 18100 HD 21856 HD 22586 HD 22928 HD 22951 HD 23180 HD 23478 HD 24190 HD 24398 HD 24534 HD 24760 HD 24912 HD 27778 HD 28497 HD 30614 HD 34029 HD 34816 HD 34989 HD 35149 HD 35715 HD 36486 HD 36822 HD 36861 HD 37021 HD 37043 HD 37061 HD 37128 HD 37367 HD 37468 HD 37903 HD 38087 HD 38666 HD 38771 HD 40111 HD 40893 HD 41117 HD 41161 HD 42087 HD 43384 HD 43818 HD 46056 HD 46202 HD 47839 HD 48915 HD 52266 HD 53367 HD 53975 HD 54662 HD 57060 HD 57061 HD 62542 HD 63005 HD 64760 HD 65575 l (3) 119.02 123.58 132.91 133.96 134.58 135.08 137.46 217.93 156.32 264.19 150.28 158.92 160.36 160.76 160.39 162.29 163.08 157.35 160.37 172.76 208.78 144.07 162.59 214.83 194.62 199.16 200.09 203.86 195.40 195.05 209.01 209.52 208.92 205.21 179.04 206.82 206.85 207.07 237.29 214.51 183.97 180.09 189.69 164.97 187.75 187.99 188.49 206.34 206.31 202.94 227.23 219.13 223.71 225.68 224.17 237.82 238.18 255.92 242.47 262.06 266.68 b (4) ­0.89 ­2.15 ­5.87 ­4.99 ­4.96 ­3.82 +7.58 ­62.73 ­16.75 ­50.36 ­5.77 ­16.70 ­17.74 ­17.42 ­15.90 ­16.69 ­17.14 ­10.09 ­13.10 ­17.39 ­37.40 14.04 +4.57 ­26.24 ­15.61 ­17.86 ­17.22 ­17.74 ­12.29 ­12.00 ­19.38 ­19.58 ­19.27 ­17.24 ­1.03 ­17.34 ­16.54 ­16.26 ­27.10 ­18.50 +0.84 +4.34 ­0.86 +12.89 +1.77 +3.53 +3.87 ­2.25 ­2.00 +2.20 ­8.89 ­0.68 ­1.90 ­2.32 ­0.78 ­5.37 ­5.54 ­9.24 ­0.93 ­10.42 ­12.32 Spectrum (5) B1II B0IVpe ON9V O8V O9.7IIn O6.5V O9.5II-IIIn B5II-III B1V B2III B5III B0.5V B1IVSB B3IV B2V B1Ib O9.5pe B0.5IV O7V B3V B2Vne O9.5Ia G8III+G0III B0.5IV B1V B1V B2IV O9.5II B0.5IV-V O8IIIf B0V O9III B0.5 V B0Iab B2 V S B O9.5V B1.5V B5V O9.5V B0Iab B0.5II B0IV: B2Ia O8Vn B2.5Ib B3Ia B0II O8V(n) O9V O7Ve A1V O9IV B0IV:e O7.5V O6.5V O7Iabfp O9II B5V O6Vf B0.5Ib B3IVp D,pc (6) 2702 188 3586 2391 1900 4108 3300 3100 500 2020 160 320 219 240 550 301 590 165 421 262 483 963 13 260 490 295 370 281 330 550 678 406 476 412 273 370 719 315 397 221 480 2632 909 1400 1578 1100 1623 1670 1670 313 3 1735 780 1400 1220 1870 980 396 5200 510 140 E( B - V ) (7) 0.47 0.12 0.21 0.36 0.37 0.48 0.31 0.05 0.19 0.06 0.05 0.19 0.29 0.28 0.30 0.27 0.59 0.11 0.35 0.36 0.05 0.29 0.01 0.05 0.10 0.12 0.04 0.08 0.08 0.09 0.48 0.07 0.53 0.04 0.38 0.06 0.33 0.31 0.06 0.12 0.18 0.45 0.40 0.21 0.37 0.58 0.52 0.49 0.47 0.07 ­0.01 0.26 0.74 0.185 0.26 0.14 0.10 0.36 0.28 0.08 0.05 [Mg/H]d (8) 32.56 ± 3.23 34.18 ± 1.98 34.18 ± 1.90 33.19 ± 2.23 33.34 ± 2.44 33.49 ± 2.13 31.29 ± 3.37 29.08 ± 4.67 ··· 36.92 ± 1.14 39.18 ± 0.23 38.14 ± 1.21 37.71 ± 1.40 ··· ··· 36.49 ± 1.42 38.54 ± 0.21 35.53 ± 1.85 35.53 ± 1.64 38.70 ± 0.20 32.80 ± 3.15 33.04 ± 3.04 36.08 ± 0.79 32.56 ± 3.34 ··· 29.57 ± 4.88 ··· 31.86 ± 1.43 34.18 ± 2.75 32.04 ± 3.10 38.02 ± 0.75 32.04 ± 1.81 38.63 ± 0.41 33.19 ± 1.83 35.12 ± 1.94 29.57 ± 3.22 38.29 ± 0.55 ··· 21.60 ± 1.72 33.04 ± 1.22 32.21 ± 3.29 33.63 ± 1.77 ··· ··· ··· ··· 33.04 ± 1.65 ··· ··· 25.67 ± 5.38 35.63 ± 4.12 34.55 ± 1.61 ··· 30.68 ± 2.46 36.71 ± 2.66 29.80 ± 5.19 25.35 ± 4.91 ··· 34.67 ± 1.19 30.25 ± 3.92 39.45 ± 0.00 [Si/H]d (9) ··· ··· ··· ··· ··· ··· ··· 16.91 ± 6. 23.49 ± 6. ··· ··· ··· ··· ··· ··· ··· 31.35 ± 0. ··· 30.78 ± 0. ··· 23.34 ± 2. ··· 23.28 ± 2. ··· ··· 25.94 ± 0. ··· 24.64 ± 1. ··· ··· 29.95 ± 1. ··· 30.26 ± 0. 24.99 ± 1. ··· ··· ··· ··· 12.45 ± 1. ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 19.81 ± 3. ··· ··· ··· ··· ··· ··· 23.07 ± 1. ··· ··· ··· ··· [Fe/H]d (10) ··· 30.71 ± 0. 30.86 ± 0. ··· 30.15 ± 0. ··· 30.12 ± 0. 28.29 ± 1. 30.69 ± 0. 30.02 ± 0. 31.49 ± 0. ··· 31.23 ± 0. ··· ··· 31.43 ± 0. 31.43 ± 0. 30.95 ± 0. 31.33 ± 0. 31.49 ± 0. ·· · 30.55 ± 0. 29.94 ± 0. ··· 31.13 ± 0. 31.04 ± 0. 31.08 ± 0. 30.86 ± 0. 31.18 ± 0. 30.95 ± 0. 30.98 ± 0. 30.45 ± 0. 31.09 ± 0. 31.02 ± 0. ··· 30.50 ± 0. 31.33 ± 0. 31.14 ± 0. 28.78 ± 0. 31.05 ± 0. 30.81 ± 0. 31.00 ± 0. 30.88 ± 0. 30.81 ± 0. 31.13 ± 0. 30.53 ± 0. ··· 30.99 ± 0. 31.28 ± 0. ·· · ··· ··· 31.32 ± 0. 31.10 ± 0. 31.28 ± 0. 30.31 ± 0. 30.45 ± 0. 29.77 ± 1. ··· 30.34 ± 0. 31.47 ± 0.

44 29 62 86 53 72 93 04 30 13 08 35 10 06 67 35 38 25 42 18 24 29 24 45 14 20 70 14 16 20 14 60 21 38 22 08 65 10 07

75 51

97 49 13 60 97 21 53 78 88

88

55

24

12 14 18 63 39 27 65 03

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N. V. Voshchinnikov and Th. Henning: From interstellar abundances to grain composition Table A.1. continued. N (1) 62 63 64 65 66 67 68 69 70 71 72 73 74 75 76 77 78 79 80 81 82 83 84 85 86 87 88 89 90 91 92 93 94 95 96 97 98 99 100 101 102 103 104 105 106 107 108 109 110 111 112 113 114 115 116 117 118 119 120 121 122 123 124 125 Star (2) HD 65818 HD 66788 HD 66811 HD 68273 HD 69106 HD 71634 HD 72754 HD 73882 HD 74375 HD 75309 HD 79186 HD 79351 HD 88115 HD 90087 HD 91316 HD 91597 HD 91651 HD 91824 HD 91983 HD 92554 HD 93030 HD 93205 HD 93222 HD 93521 HD 93843 HD 94493 HD 99857 HD 99890 HD 100340 HD 103779 HD 104705 HD 106490 HD 108248 HD 108639 HD 109399 HD 110432 HD 111934 HD 113904 HD 114886 HD 115071 HD 116658 HD 116781 HD 116852 HD 119608 HD 121263 HD 121968 HD 122879 HD 124314 HD 127972 HD 135591 HD 136298 HD 137595 HD 138690 HD 141637 HD 143018 HD 143118 HD 143275 HD 144217 HD 144470 HD 144965 HD 147165 HD 147683 HD 147888 HD 147933 l (3) 263.48 245.43 255.98 262.80 254.52 273.32 266.83 260.18 275.82 265.86 267.36 277.69 285.32 285.16 234.89 286.86 286.55 285.70 285.88 287.60 289.60 287.57 287.74 183.14 228.24 289.01 294.78 291.75 258.85 296.85 297.45 298.23 300.13 300.22 301.71 301.96 303.20 304.67 305.52 305.76 316.11 307.05 304.88 320.35 314.07 333.97 312.26 312.67 322.77 320.13 331.32 336.72 333.19 346.10 347.21 338.77 350.10 353.19 352.76 339.04 351.33 344.86 353.65 353.68 b (4) ­10.28 +2.05 ­4.71 ­7.68 ­1.33 ­11.52 ­5.82 +0.64 ­10.86 ­1.90 +2.25 ­7.37 ­5.53 ­2.13 +52.77 ­2.37 ­1.72 +0.07 +0.05 ­2.02 ­4.90 ­0.71 ­1.02 +62.15 ­0.90 ­1.18 ­4.94 +4.43 +61.23 ­1.02 ­0.34 +3.79 ­0.36 +1.95 ­9.88 ­0.20 +2.51 ­2.49 ­0.83 +0.15 +50.84 ­0.07 ­16.13 +43.13 +14.19 +55.84 +1.79 ­0.42 +16.67 ­2.64 +13.82 +18.86 +11.89 +21.70 +20.23 +11.01 +22.49 +23.60 +22.76 +8.42 +17.00 +10.09 +17.71 +17.70 Spectrum (5) B2II/IIIn O8V O5Ibnf WC8+O9I B0.5II B5III B2Ia:pshe O8.5V B1.5III B2Ib/II B5Ia B2IV-V B1.5IIn B2/B3III B1Iab B7/B8IV/V O9V:n O7V((f)) B1III O5III B0V O3V O7III((f)) O9Vp O6III B0.5Iab B1Ib B0.5V: B1V B0.5II B0.5III B2IV B0.5IV B1III B1Ib B0.5IIIe B2Ib WC5+B0Ia O9IIIn B0.5V B1III-IV B0IIIe O9III B1Ib B2.5IV B1V B0Ia O6Vnf B1.5Vne O7.5IIIf B1.5IV B3Vn B2IV B2.5Vn B1V B2.5IV B0.3IVe B0.5V B1V B2Vne B1IIISB,V B4V B3V:SB B1.5V D,pc (6) 290 4200 330 350 3076 400 690 759 440 2924 980 140 3654 2716 1754 6400 2964 2910 2910 6795 140 3187 2201 1760 2548 3888 3058 3070 3000 3061 2082 110 100 110 1900 301 2525 2660 1000 1200 80 1492 4832 4200 120 3800 2265 1100 90 1250 210 400 130 160 141 140 123 163 183 290 137 280 195 118 E( B - V ) (7) 0.06 0.20 0.05 0.03 0.20 0.13 0.36 0.67 0.14 0.28 0.40 0.10 0.20 0.30 0.04 0.27 0.28 0.25 0.29 0.39 0.04 0.38 0.33 0.04 0.27 0.23 0.33 0.24 0.04 0.21 0.28 0.06 0.20 0.35 0.26 0.51 0.51 0.21 0.29 0.44 0.14 0.34 0.21 0.12 0.05 0.15 0.36 0.46 0.11 0.22 0.07 0.25 0.07 0.18 0.07 0.02 0.21 0.21 0.22 0.35 0.41 0.39 0.51 0.47 [Mg/H]d (8) 31.67 ± 5.56 32.72 ± 3.36 26.92 ± 2.49 30.68 ± 4.64 34.55 ± 1.09 36.98 ± 1.53 ··· ··· ··· 33.77 ± 2.30 34.79 ± 1.71 34.05 ± 0.00 ··· ··· 25.35 ± 6.66 33.04 ± 1.83 27.21 ± 2.92 30.25 ± 2.22 30.89 ± 3.37 27.50 ± 4.53 31.29 ± 3.44 31.48 ± 1.30 30.03 ± 1.92 25.67 ± 9.98 31.86 ± 1.87 30.03 ± 1.53 32.72 ± 2.67 23.20 ± 6.72 29.33 ± 3.95 30.47 ± 2.33 29.57 ± 2.55 30.25 ± 1.31 22.02 ± 5.66 ··· 33.91 ± 1.91 ··· 33.49 ± 2.70 35.12 ± 2.93 ··· ··· 34.29 ± 3.11 31.86 ± 2.51 32.04 ± 2.65 29.08 ± 2.75 34.67 ± 0.66 ··· 32.04 ± 2.72 33.34 ± 1.61 22.02 ± 2.09 ··· 22.82 ± 2.00 ··· 36.17 ± 0.80 34.67 ± 2.38 36.98 ± 2.18 36.56 ± 0.72 36.08 ± 1.08 36.00 ± 0.64 35.91 ± 1.38 ··· 36.08 ± 3.03 ··· 37.85 ± 1.02 38.70 ± 1.03 [Si/H]d (9) ··· ··· 13.78 ± 9. 21.44 ± 3. ··· ··· ··· ··· ··· ··· ··· ··· ··· 29.80 ± 1. 17.95 ± 4. ··· ··· ··· ··· ··· ··· ··· ··· 15.42 ± 4. ··· ··· ··· ··· 26.78 ± 1. ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 19.52 ± 8. ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 30.00 ± 0. 27.83 ± 0. ··· 29.65 ± 1. 29.38 ± 0. ··· ··· 30.26 ± 1. ··· 30.11 ± 1. 31.68 ± 0. [Fe/H]d (10) ·· · 30.84 ± 0. 30.28 ± 0. 29.86 ± 0. 31.15 ± 0. ·· · ·· · 31.08 ± 0. 30.95 ± 0. ·· · ·· · ·· · ·· · 30.74 ± 0. ·· · 30.37 ± 0. 29.98 ± 0. ·· · ·· · 29.90 ± 0. 30.73 ± 0. 30.37 ± 0. 30.25 ± 0. 26.35 ± 2. 30.25 ± 0. 29.36 ± 0. 30.67 ± 0. 29.51 ± 0. 27.97 ± 0. 30.67 ± 0. 29.94 ± 0. 29.90 ± 0. 29.15 ± 0. ·· · 30.50 ± 0. 31.47 ± 0. ·· · ·· · ·· · ·· · 29.41 ± 1. 30.53 ± 0. 29.94 ± 0. ·· · 30.77 ± 0. ·· · 30.79 ± 0. 31.02 ± 0. 28.91 ± 0. 31.16 ± 0. 29.56 ± 0. ·· · 31.12 ± 0. 31.12 ± 0. 31.01 ± 0. 30.15 ± 0. 31.21 ± 0. 31.13 ± 0. 31.26 ± 0. ·· · 31.04 ± 0. ·· · 31.02 ± 0. 31.44 ± 0.

48 43

33 29 57 45 19 54

90 01

15 41 36 67 32 24 47 31 33 58 30 92 88 30 54 23 94 37 07

36

37

65

50 48 48 12 36 19 29 41 32 06 43 26 23 20 35 17 52 28 16

87 97 01 40 22 29 23

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A&A 517, A45 (2010) Table A.1. continued. N (1) 126 127 128 129 130 131 132 133 134 135 136 137 138 139 140 141 142 143 144 145 146 147 148 149 150 151 152 153 154 155 156 157 158 159 160 161 162 163 164 165 166 167 168 169 170 171 172 173 174 175 176 177 178 179 180 181 182 183 184 185 186 187 188 189 Star (2) 148184 148594 149404 149757 149881 151804 151805 152236 152590 154368 155806 156110 157246 157857 158926 160578 164740 165024 165955 167264 167756 168076 170740 175360 177989 179406 184915 185418 186994 188209 190918 192035 192639 195965 197512 198478 198781 199579 201345 202347 202904 203374 203532 206267 206773 207198 207308 207538 208440 208947 209339 210121 210809 210839 212791 214680 214993 215733 218376 218915 219188 220057 224151 224572 l (3) 357.93 350.93 340.54 6.28 31.37 343.62 343.20 343.03 344.84 349.97 352.59 70.99 334.64 12.97 351.74 351.04 5.97 343.33 357.41 10.46 351.47 16.94 21.06 12.53 17.81 28.23 31.77 53.60 78.62 80.99 72.65 83.33 74.90 85.71 87.89 85.75 99.94 87.50 78.44 88.22 80.98 100.51 309.46 99.29 99.80 103.14 103.11 101.60 104.03 106.55 104.58 56.88 99.85 103.83 101.64 96.65 97.65 85.16 109.96 108.06 83.03 112.13 115.44 115.55 b (4) +20.68 +13.94 +3.01 +23.59 +36.23 +1.94 +1.59 +0.87 +1.83 +3.22 +2.87 +35.91 ­11.48 +13.31 ­2.21 ­4.72 ­1.17 ­13.82 ­7.43 ­1.74 ­12.30 +0.84 ­0.53 ­11.29 ­11.88 ­8.31 ­13.29 ­2.17 +10.06 +10.09 +2.06 +7.76 +1.48 +5.00 +4.63 +1.49 +12.61 ­0.30 ­9.54 ­2.08 ­10.05 +8.62 ­31.74 +3.74 +3.62 +6.99 +6.82 +4.67 +6.44 +9.00 +5.87 ­44.46 ­3.13 +2.61 ­4.30 ­16.98 ­16.18 ­36.35 ­0.79 ­6.89 ­50.17 +0.21 ­4.64 ­6.36 Spectrum (5) B1.5Ve B9:V O9Ia O9.5Vnn B0.5III O8Iab B1Ib B1Ia O7.5V O9Ib O7.5Ve B3Vn B1Ib O7V B2IV B1.5III O7.5V(n) B2Ib B1Vnp B0.5Ia B0.5Iab? O5V B2V B6III B2II B3IVvar B0.5IIIne B0.5 V B0III O9.5Ib WN4+O9.7Iab B0III-IVn O8V B0V B1V B3Ia B2IV B0.5V O9V B1V B2Vne B0IVpe B5V O6V B0V O9II B0.7III-IVn O9.5V B1V B2V B0IV B9V O9Ib O6Iab B8 O9V B1.5IIIn B1II B1III O9.5Iabe B0.5III B2IV B0.5II-III B1V D, pc (6) 160 134 908 146 2100 1254 6009 612 1800 1046 860 720 348 1902 220 142 1330 250 1640 1514 4230 1820 235 270 5021 227 700 1027 2500 2210 2290 2800 999 1300 1614 890 768 990 2570 1300 276 820 211 814 597 1216 1470 880 620 500 980 223 3961 1260 370 610 610 2900 383 3660 1064 1421 1355 340 E( B - V ) (7) 0.44 0.21 0.62 0.31 0.11 0.30 0.43 0.60 0.46 0.76 0.28 0.03 0.06 0.43 0.10 0.08 0.86 0.05 0.21 0.30 0.07 0.76 0.47 0.12 0.22 0.31 0.22 0.47 0.16 0.15 0.41 0.35 0.61 0.22 0.29 0.57 0.31 0.33 0.17 0.19 0.13 0.60 0.28 0.52 0.45 0.54 0.52 0.64 0.34 0.19 0.35 0.38 0.31 0.57 0.06 0.08 0.10 0.10 0.23 0.26 0.08 0.24 0.42 0.19 [Mg/H]d (8) 37.76 ± 1.50 37.80 ± 0.54 ··· 37.56 ± 0.49 24.30 ± 1.40 33.34 ± 3.92 ··· ··· 34.43 ± 2.03 ··· 33.04 ± 3.38 30.89 ± 0.78 32.88 ± 2.65 34.18 ± 2.21 ··· 31.67 ± 4.37 ··· 33.91 ± 2.00 31.48 ± 2.25 36.08 ± 2.83 31.67 ± 2.49 ··· ··· 35.82 ± 1.18 34.18 ± 1.62 ··· 36.49 ± 2.44 36.25 ± 1.09 ··· ··· 31.09 ± 2.25 36.25 ± 1.11 34.43 ± 1.62 34.05 ± 1.04 ··· 37.29 ± 1.32 35.73 ± 1.22 ··· 30.68 ± 2.75 34.90 ± 2.17 ··· 34.30 ± 2.01 36.17 ± 3.46 36.56 ± 1.15 34.30 ± 1.22 37.29 ± 0.52 36.71 ± 1.00 36.71 ± 1.01 34.18 ± 1.98 ··· 33.77 ± 1.34 ··· 31.86 ± 3.01 35.82 ± 1.09 34.30 ± 3.98 30.68 ± 4.09 35.01 ± 2.88 30.89 ± 5.87 34.79 ± 4.56 ··· 33.34 ± 6.30 36.56 ± 1.41 32.88 ± 1.70 36.00 ± 2.16 [Si/H]d (9) ··· ··· ··· 30.45 ± 0.40 18.27 ±11.41 ··· ··· ··· 30.45 ± 0.79 ··· ··· ··· ··· ··· 25.16 ± 0.95 26.90 ± 2.42 ··· ··· ··· ··· 20.65 ± 4.78 ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 27.61 ± 4.47 ··· ··· ··· ··· ··· ··· 27.83 ± 4.10 28.93 ± 3.01 ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· ··· 21.93 ± 5.74 ··· ··· 27.72 ± 2.70 ··· ··· ··· [Fe/H]d (10) 31.28 ± 0. ··· 30.98 ± 0. 31.43 ± 0. 28.05 ± 2. 30.77 ± 0. ··· 31.35 ± 0. 31.08 ± 0. 31.08 ± 0. 30.94 ± 0. ··· 30.71 ± 0. ··· 30.15 ± 0. 30.86 ± 0. 31.48 ± 0. 30.86 ± 0. ··· ··· 30.05 ± 0. 31.08 ± 0. 31.41 ± 0. ··· 30.86 ± 0. 31.17 ± 0. ··· 31.21 ± 0. 30.34 ± 0. 30.58 ± 0. ··· ··· 31.14 ± 0. 30.87 ± 0. 31.18 ± 0. ··· ··· ··· ··· 30.79 ± 0. 30.82 ± 0. 30.82 ± 0. ··· 31.33 ± 0. ··· 31.32 ± 0. 31.14 ± 0. 31.32 ± 0. ··· ··· 30.92 ± 0. 29.86 ± 0. ··· 31.13 ± 0. ··· 31.22 ± 0. 30.05 ± 1. 29.41 ± 1. 30.73 ± 0. ··· ··· ··· 30.45 ± 0. 31.00 ± 0.

HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD HD

23 11 04 47 46 14 10 26 55 44 22 49 03 34 50 45 12 27 33 21 42 71 21 10 12

34 76 31 13 09 19 19 18 88 16 31 18 04 85

37 53

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N. V. Voshchinnikov and Th. Henning: From interstellar abundances to grain composition Table A.1. continued. N (1) 190 191 192 193 194 195 196 Star (2) HD 232522 HD 303308 HD 308813 BD +35 4258 BD +53 2820 CPD -59 2603 CPD -69 1743 l (3) 130.70 287.59 294.79 77.19 101.24 287.59 303.71 b (4) ­6.71 ­0.61 ­1.61 ­4.74 ­1.69 ­0.69 ­7.35 Spectrum (5) B1II O3V O9.5V B0.5 Vn B0IV:n O7V B1Vn D,pc (6) 5438 3631 2398 3093 4506 2630 4700 E( B - V ) (7) 0.21 0.45 0.28 0.25 0.37 0.46 0.30 [Mg/H]d (8) 31.48 ± 1.94 32.56 ± 2.16 32.56 ± 2.08 32.72 ± 2.97 ··· 32.72 ± 1.56 ··· [Si/H]d (9) ··· ··· ··· ··· ··· 18.91 ± 3.64 ··· [Fe/H]d (10) ··· 30.28 ± 0.41 ··· 30.22 ± 0.64 ··· 30.37 ± 0.32 ···

Notes. Dust phase abundances are calculated as difference between solar abundances ([Mg/H] = 39.8 ppm, [Si/H] = 32.4 ppm, [Fe/H] = 31.6 ppm, Asplund et al. 2009) and gas phase abundances. () For these stars only gas phase abundances of O are measured.

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