3.1.2. Spectroscopic Extinction and Polarization Features
It is the extinction (absorption) and emission spectral lines instead of the overall shape of the extinction curve provides the most diagnostic information on the dust composition.
1. The 2175 Å Extinction Hump
"It is frustrating that almost 3 decades after
the identity of this (2175 Å) feature remains uncertain!"
-- B.T. Draine 
The strongest spectroscopic extinction feature is the 2175 Å hump. Observations show that its strength and width vary with environment while its peak position is quite invariant. Its carrier remains unidentified 37 years after its first detection (Stecher 1965). Many candidate materials, including graphite (Stecher & Donn 1965), amorphous carbon (Bussoletti et al. 1987), graphitized (dehydrogenated) hydrogenated amorphous carbon (Hecht 1986; Goebel 1987; Sorrell 1990; Mennella et al. 1996; Blanco et al. 1999), (16) nano-sized hydrogenated amorphous carbon (Schnaiter et al. 1998), quenched carbonaceous composite (QCC; Sakata et al. 1995), coals (Papoular et al. 1995), PAHs (Joblin et al. 1992; Duley & Seahra 1998; Li & Draine 2001b), and OH- ion in low-coordination sites on or within silicate grains (Duley, Jones & Williams 1989) have been proposed, while no single one is generally accepted (see Draine 1989b for a review).
Graphite was the earliest suggested and the widely adopted candidate in various dust models (Gilra 1972; Mathis et al. 1977; Hong & Greenberg 1980; Draine & Lee 1984; Dwek et al. 1997; Will & Aannestad 1999). However the hump peak position predicted from graphite particles is quite sensitive to the grain size, shape, and coatings (Gilra 1972; Greenberg & Chlewicki 1983; Draine 1988; Draine & Malhotra 1993) which is inconsistent with the observations. It was suggested that very small coated graphite particles (a 0.006 µm) could broaden the hump while keeping the hump peak constant (Mathis 1994). (17) However, this seems unlikely because the proposed particles are so small that temperature fluctuations will prevent them from acquiring a coating (Greenberg & Hong 1974a; Aannestad & Kenyon 1979). Furthermore, it was noted by Greenberg & Hong (1974b) that if the very small particles as well as the large particles accrete mantles the RV value would decrease in molecular clouds, in contradiction with astronomical observations. Rouleau, Henning, & Stognienko (1997) proposed that the combined effects of shape, clustering, and fine-tuning of the optical properties of graphite could account for the hump width variability.
Another negative for graphite, is that the dust grains in circumstellar envelopes around carbon stars which are the major sources of the carbon component of interstellar dust are in amorphous form rather than graphitic (Jura 1986). It is difficult to understand how the original amorphous carbonaceous grains blown out from the star envelopes are processed to be highly anisotropic and evolve to the layer-lattice graphitic structures in interstellar space. Instead, it is more likely that the interstellar physical and chemical processes should make the carbonaceous grains even more highly disordered.
Recently, the PAH proposal is receiving increasing attention. Although a single PAH species often has some strong and narrow UV bands which are not observed (UV Atlas 1966), a cosmic mixture of many individual molecules, radicals, and ions, with a concentration of strong absorption features in the 2000-2400 Å region, may effectively produce the 2175 Å extinction feature (Li & Draine 2001b). This is supported by the correlation between the 2175 Å hump and the IRAS 12 µm emission (dominated by PAHs) found by Boulanger, Prévot, & Gry (1994) in the Chamaeleon cloud which suggests a common carrier. Arguments against the PAHs proposal also exist (see Li & Draine 2001b for references).
So far only two lines of sight toward HD 147933 and HD 197770 have a weak 2175 Å polarization feature detected (Clayton et al. 1992; Anderson et al. 1996; Wolff et al. 1997; Martin, Clayton, & Wolff 1999). Even for these sightlines, the degree of alignment and/or polarizing ability of the carrier should be very small (if both the hump excess polarization and the hump extinction are produced by the same carrier); for example, along the line of sight to HD 197770, the ratio of the excess polarization to the hump extinction is P2175 Å / A2175 Å 0.002 while the polarization to extinction ratio in the visual is PV / AV 0.025, thus (P2175 Å / A2175 Å) / (PV / AV) is only ~ 0.09. Therefore, it is reasonable to conclude that the 2175 Å carrier is either mainly spherical or poorly aligned.
The 2175 Å hump polarization was predicted by Draine (1988) for aligned non-spherical graphite grains. Wolff et al. (1993) and Martin et al. (1995) further show that the observed 2175 Å polarization feature toward HD 197770 can be well fitted with small aligned graphite disks.
Except for the detection of scattering in the 2175 Å hump in two reflection nebulae (Witt, Bohlin, & Stecher 1986), the 2175 Å hump is thought to be predominantly due to absorption, suggesting its carrier is sufficiently small to be in the Rayleigh limit.
2. The 9.7 µm and 18 µm (Silicate) Absorption Features
"In my opinion, the only secure identification
is that of the 9.7 µm and 18 µm IR features."
-- B.T. Draine 
The strongest IR absorption features are the 9.7 µm and 18 µm bands. They are respectively ascribed to the Si-O stretch and O-Si-O bending modes in some form of silicate material, perhaps olivine Mg2xFe2-2xSiO4. The shape of the interstellar silicate feature is broad and featureless both of which suggest that the silicate is amorphous. (18) The originating source of the interstellar silicates is in the atmospheres of cool evolved stars of which the emission features often show a 9.7 µm feature consistent with amorphous silicates and sharper features arising from crystalline silicates (Waters et al. 1996). (19) How can interstellar crystalline silicate particles become amorphous? This is a puzzle which has not yet been completely solved. A possible solution may be that the energetic processes (e.g. shocks, grain-grain collision) operated on interstellar dust in the diffuse ISM have disordered the periodic lattice structures of crystalline silicates. Another possible solution may lie in the fact that, according to Draine (1990), only a small fraction of interstellar dust is the original stardust, i.e., most of the dust mass in the ISM was condensed in the ISM, rather than in stellar outflows. The re-condensation at low temperatures most likely leads to an amorphous rather than crystalline form.
First detected in the Becklin-Neugebauer (BN) object in the OMC-1 Orion dense molecular cloud (Dyck et al. 1973), the silicate polarization absorption feature in the 10 µm region is found to be very common in heavily obscured sources; some sources also have the 18 µm O-Si-O polarization feature detected (see Aitken 1996; Smith et al. 2000 for summaries). In most cases the silicate polarization features are featureless, indicating the amorphous nature of interstellar silicate material (see Aitken 1996) except AFGL 2591, a molecular cloud surrounding a young stellar object, has an additional narrow polarization feature at 11.2 µm, generally attributed to annealed silicates (Aitken et al. 1988; Wright et al. 1999).
Reasonably good fits to the observed 10 µm Si-O polarization features can be obtained by elongated (bare or ice-coated) "astronomical silicate" grains (Draine & Lee 1984; Lee & Draine 1985; Hildebrand & Dragovan 1995; Smith et al. 2000). High resolution observations of the BN 10 µm and 18 µm polarization features provided a challenge to the "astronomical silicate" model since this model failed to reproduce two of the basic aspects of the observations: (1) the 10 µm feature was not broad enough and (2) the 18 µm feature was too low by a factor of two relative to the 10 µm peak (Aitken, Smith, & Roche 1989). Attempts by Henning & Stognienko (1993) in terms of porous grains composed of "astronomical silicates", carbon and vacuum were not successful. In contrast, it is shown by Greenberg & Li (1996) that an excellent match to the BN 10 µm and 18 µm polarization features in shape, width, and in relative strength can be obtained by the silicate core-organic mantle model using the experimental optical constants of silicate and organic refractory materials. It seems desirable to re-investigate the 18 µm O-Si-O band strength of "astronomical silicates" (Draine & Lee 1984) by a combination of observational, experimental and modelling efforts.
In the mid-IR, interstellar grains of submicron size are in the Rayleigh limit. The scattering efficiency falls rapidly with increasing wavelength. Therefore the effects of scattering are expected to be negligible for the silicate extinction and polarization.
3. The 3.4 µm (Aliphatic Hydrocarbon) Absorption Feature
"When the 3.4 µm feature was detected in
VI Cyg #12, the problem
of why no H2O was observed would appear to have been resolved."
-- J. Mayo Greenberg 
Another ubiquitous strong absorption band in the diffuse ISM is the 3.4 µm feature. Since its first detection in the Galactic Center toward Sgr A W by Willner et al. (1979) and IRS 7 by Wickramasinghe & Allen (1980), it has now been widely seen in the Milky Way Galaxy and other galaxies (Butchart et al. 1986; Adamson, Whittet, & Duley 1990; Sandford et al. 1991; Pendleton et al. 1994; Wright et al. 1996; Imanishi & Dudley 2000; Imanishi 2000). Although it is generally accepted that this feature is due to the C-H stretching mode in saturated aliphatic hydrocarbons, the exact nature of this hydrocarbon material remains uncertain. Nearly two dozen different candidates have been proposed over the past 20 years (see Pendleton & Allamandola 2002 for a review). The organic refractory residue, synthesized from UV photoprocessing of interstellar ice mixtures, provides a perfect match, better than any other hydrocarbon analogs, to the observed 3.4 µm band, including the 3.42 µm, 3.48 µm, and 3.51 µm subfeatures (Greenberg et al. 1995). (20) But, at this moment, we are not at a position to rule out other dust sources as the interstellar 3.4 µm feature carrier. This feature has also been detected in a carbon-rich protoplanetary nebula CRL 618 (Lequeux & Jourdain de Muizon 1990; Chiar et al. 1998) with close resemblance to the interstellar feature. However, after ejection into interstellar space, the survival of this dust in the diffuse ISM is questionable (see Draine 1990).
The 3.4 µm feature consists of three subfeatures at 2955 cm-1 (3.385 µm), 2925 cm-1 (3.420 µm), and 2870 cm-1 (3.485 µm) corresponding to the symmetric and asymmetric C-H stretches in CH3 and CH2 groups in aliphatic hydrocarbons which must be interacting with other chemical groups. The amount of carbonaceous material responsible for the 3.4 µm feature is strongly dependent on the nature of the chemical groups attached to the aliphatic carbons. For example, each carbonyl (C = O) group reduces its corresponding C-H stretch strength by a factor of ~ 10 (Wexler 1967). Furthermore, not every carbon is attached to a hydrogen as in saturated compounds. Aromatic hydrocarbons do not even contribute to the 3.4 µm feature although they do absorb nearby at 3.28 µm. The fact that the 3.4 µm absorption is not observed in molecular cloud may possibly be attributed to dehydrogenation or oxidation (formation of carbonyl) of the organic refractory mantle by accretion and photoprocessing in the dense molecular cloud medium (see Greenberg & Li 1999a) - the former reducing the absolute number of CH stretches, the latter reducing the CH stretch strength by a factor of 10 (Wexler 1967) - the 3.4 µm feature would be reduced per unit mass in molecular clouds.
Very recently, Gibb & Whittet (2002) reported the discovery of a 6.0 µm feature in dense clouds attributed to the organic refractory. They found that its strength is correlated with the 4.62 µm OCN- (XCN) feature which is considered to be a diagnostic of energetic processing.
Attempts to measure the polarization of the 3.4 µm absorption feature (PC-HIRS7-obs) was recently made by Adamson et al. (1999) toward the Galactic Center source IRS 7. They found that this feature was essentially unpolarized. Since no spectropolarimetric observation of the 10 µm silicate absorption feature (PsilIRS7) has yet been carried out for IRS 7, they estimated PsilIRS7 from the 10 µm silicate optical depth silIRS7, assuming the IRS 7 silicate feature is polarized to the same degree as the IRS 3 silicate feature; i.e., PsilIRS7 / silIRS7 = PsilIRS3 / silIRS3 where silIRS7, silIRS3 and PsilIRS3 were known. Assuming the IRS 7 aliphatic carbon (the 3.4 µm carrier) is aligned to the same degree as the silicate dust, they expected the 3.4 µm polarization to be PC-HIRS7-mod = PsilIRS7 / silIRS7 × C-HIRS7. They found PC-HIRS7-obs << PC-HIRS7-mod (Adamson et al. 1999). Therefore, they concluded that the aliphatic carbon dust is not in the form of a mantle on the silicate dust as suggested by the core-mantle models (Li & Greenberg 1997; Jones, Duley, & Williams 1990). (21) We note that the two key assumptions on which their conclusion relies are questionable: (1) PsilIRS7 / silIRS7 = PsilIRS3 / silIRS3; (2) PC-HIRS7 / C-HIRS7 = PsilIRS7 / silIRS7 (see Li & Greenberg 2002 for details). We urgently need spectropolarimetric observations of IRS 7.
Hough et al. (1996) reported the detection of a weak 3.47 µm polarization feature in BN, attributed to carbonaceous materials with diamond-like structure, originally proposed by Allamandola et al. (1992) based on the 3.47 µm absorption spectra of protostars. (22)
4. The 3.3 µm and 6.2 µm (PAH) Absorption Features
Recently, two weak narrow absorption features at 3.3 µm and 6.2 µm were detected. The 3.3 µm feature has been seen in the Galactic Center source GCS 3 (Chiar et al. 2000) and in some heavily extincted molecular cloud sight lines (Sellgren et al. 1995; Brooke, Sellgren, & Geballe 1999). The 6.2 µm feature, about 10 times stronger, has been detected in several objects including both local sources and Galactic Center sources (Schutte et al. 1998; Chiar et al. 2000). They were attributed to aromatic hydrocarbons (Schutte et al. 1998; Chiar et al. 2000). The theoretical 3.3 µm and 6.2 µm absorption feature strengths (in terms of integrated optical depths) predicted from the astronomical PAH model are consistent with observations (Li & Draine 2001b). Note the 7.7, 8.6, and 11.3 µm PAH features are hidden by the much stronger 9.7 µm silicate feature, and therefore will be difficult to observe as absorption features.
The aromatic absorption bands allow one to place constraints on the PAH abundance if the PAH band strengths are known. But one should keep in mind that the strengths of these PAH absorption bands could vary with physical conditions due to changes in the PAH ionization fraction. In regions with increased PAH ionization fraction, the 6.2 µm absorption feature would be strengthened and the 3.3 µm feature would be weakened (see Li & Draine 2001b).
Although it was theoretically predicted that the PAH IR emission features can be linearly polarized (Léger 1988), no polarization has been detected yet (Sellgren, Rouan, & Léger 1988).
5. The Diffuse Interstellar Bands
In 1922, Heger observed two broad absorption features centering at 5780 Å and 5797 Å, conspicuously broader than atomic interstellar absorption lines. This marked the birth of a long standing astrophysical mystery - the diffuse interstellar bands (DIBs). But not until the work of Merrill (1934) were the interstellar nature of these absorption features established. So far, over 300 DIBs have been detected from the near IR to the near UV. Despite ~ 80 years' efforts, no definite identification (including the recent neutral/charged PAHs [Salama & Allamandola 1992], C60+ [Foing & Ehrenfreund 1994], and C7- [Tulej et al. 1998] proposals) of the carrier(s) of DIBs has been found. We refer the readers to the two extensive reviews of Krelowski (1999, 2002).
No polarization has been detected for the DIBs (Martin & Angel 1974, 1975; Fahlman & Walker 1975; Adamson & Whittet 1992, 1995; see Somerville 1996 for a review).
6. The Ice Absorption Features
The formation of an icy mantle through accretion of molecules on interstellar dust is expected to take place in dense clouds (e.g., the accretion timescale is only ~ 105 yrs for clouds of densities nH = 103-105 cm-3; Schutte 1996). The detection of various ice IR absorption features (e.g., H2O [3.05, 6.0 µm], CO [4.67 µm], CO2 [4.27, 15.2 µm], CH3OH [3.54, 9.75 µm], NH3 [2.97 µm], CH4 [7.68 µm], H2CO [5.81 µm], OCN- [4.62 µm]; see Ehrenfreund & Schutte 2000 for a review) have demonstrated the presence of icy mantles in dark clouds (usually with an visual extinction > 3 magnitudes; see Whittet et al. 2001 and references therein). In comparison with that for a gas phase sample, the IR spectrum for the same sample in the solid phase is broadened, smoothed, and shifted in wavelength due to the interactions between the vibrating molecule with the surroundings, the suppression of molecular rotation in ices at low temperatures, and the irregular nature of the structure of (amorphous) solids (see Tielens & Allamandola 1987).
Note not only are the relative proportions of ice species variable in different regions but also the presence and absence of some species. In almost all cases, however, water is the dominant component. An important variability is in the layering of the various molecular components. Of particular note is the fact that the CO molecular spectrum is seen to indicate that it occurs sometimes embedded in the H2O (a polar matrix) and sometimes not. This tells a story about how the mantles form. As was first noted by van de Hulst, the presence of surface reactions leads to the reduced species H2O, CH4, NH3. Since we now know that CO is an abundant species as a gas phase molecule, we expect to find it accreted along with these reduced species - at least initially.
The two approaches to understanding how the grain mantles evolve are: (1) the laboratory studies of icy mixtures, their modification by UV photoprocessing and by heating; (2) theoretical studies combining gas phase chemistry with dust accretion and dust chemistry. In the laboratory one creates a cold surface (10K) on which various simple molecules are slowly deposited in various proportions. The processing of these mixtures by UV photons and by temperature variation is studied by IR spectroscopy. This analog of interstellar dust mantles is used to provide a data base for comparison with the observations (see Schutte 1999 for a review).
The 3.1 µm ice polarization has been detected in various molecular cloud sources (see Aitken 1996 and references therein). The detection of the 4.67 µm CO and 4.62 µm OCN- polarization was recently reported by Chrysostomou et al. (1996). The BN ice polarization feature was well fitted by ice-coated grains (Lee & Draine 1985), suggesting a core-mantle grain morphology. However, the AFGL 2591 molecular cloud shows no evidence for ice polarization (Dyck & Lonsdale 1980; Kobayashi et al. 1980) while having distinct ice extinction and silicate polarization (Aitken et al. 1988). Perhaps only the hot, partly annealed silicate grains close to the forming-star are aligned (say, by streaming of ambipolar diffusion) while the ice-coated cool grains in the outer envelope of the cloud are poorly aligned.
16 Mennella et al. (1996) reported that a stable peak position can be obtained by subjecting small hydrogenated amorphous carbon grains to UV radiation. However, the laboratory produced humps are too wide and too weak with respect to the interstellar one. In a later paper (Mennella et al. 1998) they proposed that the 2175 Å carrier can be modelled as a linear combination of such materials exposed to different degrees of UV processing. Back.
17 Hecht (1981) investigated the effect of coatings on graphite particles and concluded that small coatings could be present on spherical a 200 Å particles and still cause the 2175 Å feature. Back.
18 Very recently, Li & Draine (2001a) estimated that the abundances of a < 1 µm crystalline silicate grains in the diffuse ISM is < 5% of the solar Si abundance. Back.
19 Crystalline silicates have also been seen in six comets (see Hanner 1999 for a summary), in dust disks around main-sequence stars (see Artymowicz 2000 for a summary), young stellar objects (see Waelkens, Malfait, & Waters 2000 for a summary), in interplanetary dust particles (IDPs) (Bradley et al. 1999), and probably also in the Orion Nebula (Cesarsky et al. 2000). Back.