5 Inception and evolution

Kerogen-like materials are bound to evolve, the faster the higher the ambient temperature (see Sect. 2). One cannot, therefore, be content with the comparison of one kerogen model with one particular line of sight. Fortunately, there are enough data from both sides to pursue the analogy. However, while the 3.4 feature is conspicuous in all kerogens (see Figs. 2 and 3), the open kerogen literature does not exhibit this feature on such a scale as to allow its sub-features to be quantitatively compared with the corresponding astronomical sub-features. The only exceptions to this are to be found in the astronomical literature cited in the introduction (Ehrenfreund et al. 1991, Fig. 4). This is probably due to the fact that the 3.4-$\mu$m blend had already been analyzed and discussed at length by coal experts (see e.g. Painter et al. 1981).

Moreover, there is a dearth of data on kerogens more evolved than $H/C\sim 0.5$. This is because kerogens are collected in sedimentary layers, which lay near the earth's surface and are relatively young. Also, the oil industry is only interested in the less evolved samples, which seem to be geographically associated with oil fields. On the other hand, there is plenty of data on coals from $H/C\sim 1$ down to 0 (semi-graphite), because coal has been searched for down to below 5000 m deep. Besides, there is a very rich collection of IR spectra of coals from various parts of the world. Now, kerogens and coals in evolved stages are very similar: they occupy the same quasi-vertical strip in the van Krevelen diagram (see Sect. 2). Thus, if one looks for a parallel between astronomical dust and natural carbonaceous materials, one can and should take advantage of both kerogen and coal data according to the needs of modelling.

Even so, some interpolation between available terrestrial materials was found to be necessary. This is made possible by the fact that heat treatment has long been known to mimic the natural evolution of coal and kerogen in the earth (see Papoular et al. 1996 for references). Figures 5 and 6 were constructed with the above in mind.

Figure 5 shows 5 spectra from different types of astronomical objects: a) the GC (IRS 6E, in absorption); b-d) 3 post-AGB nebulae; e) a reflection nebula. The evolution is obvious and parallel to that in Fig. 6, which exhibits the spectra of 3 standard coals from progressively deeper mines (see Papoular et al. 1996): a) Vouters (O/C=0.06, H/C=0.75), b) Mericourt (O/C=0.028, H/C=0.59), c) Escarpelle (O/C=0.018, H/C=0.46), and spectra of an Escarpelle sample annealed successively to $225\, ^\circ$C (d) and $300\, ^\circ$C (e). The Vouters sample is the youngest coal and is akin to, but in a slightly more advanced evolutionary stage than, the model IIb kerogen chosen above. Its spectrum fits Fig. 5a better than do the kerogen samples in Fig. 4 of Ehrenfreund et al. (1991). The changes in the relative intensities of the aliphatic and aromatic stretch features clearly suggest a progressive aromatization of the carrier structure from top to bottom of Fig. 5.

In the case of Fig. 5c (IRAS 0.4296+3429), the 11-13 $\mu$m range exhibits out-of-plane C-H bend features which could be modelled, too, using the optical properties of the same (Escarpelle) coal as in Fig. 6c (Guillois et al. 1996). The same material was used successfully to model the near- and mid-IR spectra of two PPNe: IRAS 22272+5435 and 07134+1005. This indicates that dust present in the winds of AGB and post-AGB stars is already aromatic enough to exhibit a strong 3.3-$\mu$m feature. The material in Fig. 6e is still more evolved and has lost most of its heteroatoms. The same can be expected for its counterpart in 5e, which must, therefore, have spent a very long time under photon irradiation in the DISM.

In novae, aromatization is much faster and nearly completed close to the dust birthplace, as, for instance, in Nova Cen 1986 (Hyland & Mc Gregor 1988; Smith et al. 1994). This, of course, is due to the dust temperature being much higher than in the case of red giant or AGB stars.

Similarly, parallel evolutions occur in the mid-IR. For IS dust, this is illustrated by similar "transition" objects observed by Buss et al. (1990, 1993). In well-developed PNe (e.g. NGC 7027), the UIBs at 6.2, 7.7, 8.6 and 11.3 $\mu$m stand out distinctly and are relatively narrow, according to the usual pattern in RNe. By contrast, the transition objects exhibit a much greater spectral variety, with the common characteristic of reduced visibility of the bands, which are blended in two prominent massifs peaking at about 8 and 12 $\mu$m, respectively (see, e.g., IRAS 22272+5435). This can be interpreted qualitatively by comparison with kerogen spectra (see Fig. 3 and assignments in Sect. 2) and coal spectra (Guillois et al. 1996): in the transition stages, oxygen and aliphatic carbon sites are still abundant, and responsible for the strong 5.85, 6.9 and 6-10-$\mu$m features; the aromatic clusters are still small, allowing many polysubstitutions around the rings and, hence, many overlapping out-of-plane C-H vibrations. The loss of heteroatoms upon ageing "purifies" the spectrum.

Goto et al. (2000) have recently produced a similar parallelism between near-IR astronomical spectra and spectra of QCC, using thermal processing of the initially aliphatic material up to 873 K. However, they did not discuss the fingerprint region (mid-IR), which is so sensitive to oxygen content (see Figs. 2 and 3). But Sakata et al. (1987) have previously shown that unoxidized QCC is unable to mimic mid-IR celestial spectra correctly.

The structural changes which accompany these spectral changes are vividly illustrated in the sketches of the type shown in our Fig. 4, produced by Behar & Vandenbroucke (1986).

In the framework of the present model, the absence of the usual carbonaceous features in red giant envelopes, their appearance in post-AGB circumstellar shells and the progressive strengthening of aromatic features and weakening of aliphatic ones through the PPN stage to the PN stage can only mean that, at the end of the PNe stage, the dust born in AGBs is already very much processed and almost as aromatic as the carrier of the UIBs (Unidentified IR Bands) observed in reflexion nebulae.

This scenario raises a number of issues. First, how is it that such an unevolved (aliphatic) dust as in Fig. 5a is observed towards the GC, seeming to imply that IS dust is mainly aliphatic? Indeed, if this dust were entirely distributed in the DISM, the fraction of the total available carbon, $C_{{\rm avail}}$, that should have to be locked into it would be

$$C_{{\rm dust}}/C_{{\rm avail}}=1.7\times 10^{-2} \frac{H_{{\rm avail}}}{C_{{\rm avail}}}\frac{\tau(3.4)}{A_{{\rm v}}}$$ (3)

using the cross-section per atom deduced from Sect. 4 ( $2.9\times 10^{-20}$ cm² per C atom) and $N(H)=2\times 10^{21} A_{{\rm v}}$; here, $(C/H)_{{\rm avail}}$ represents the relative number of carbon atoms available for dust, and lies between 80 and 200 per million, depending on authors. Now, according to Pendleton et al. (1994), $A_{{\rm v}}/\tau(3.4)\sim 150$ towards the GC. The locked fraction then turns out to lie between 0.6 and 1.4. This is hardly available, especially in view of the fact that relevant kerogen-like material is not aromatic enough to exhibit the right extinction in the visible/UV (see Khare et al. 1989) and should be complemented with some other carbonaceous dust. The same conclusion was reached on the basis of HAC models (Duley et al. 1998; Furton et al. 1999). Thus, only part of the feature intensity can be ascribed to dust in the DISM and in molecular clouds. Based on $A_{{\rm v}}/\tau(3.4)\sim 150$ towards the CG and $\sim$270 in the local ISM, this fraction may rougly be estimated at 1/2. The rest is likely to reside in the CS shells of the very abundant population of AGB stars revealed by the IRAS satellite towards the GC, as is the case for CRL618 (Lequeux & de Muizon 1990).

This conclusion may seem to be at variance with the non-detection of the 3.4 feature in nearby AGB stars, either in emission or in absorption. Among others, Chiar et al. (1998) and Goto et al. (2000) also faced this paradox in the framework of their dust models and scenarios. Goto et al. suggested that an adequate chemical mechanism is suddenly switched on at the very end of the AGB phase. However, since carbon dust in some form clearly condenses even before the AGB phase, and in presence of abundant hydrogen, it seems less far fetched to speculate, on the ground of simple radiative theory, that, there and then, the dust temperature, density and distance from the star are such as to prevent the feature from showing up against the strong star light background. This would be in line with the very low contrast of the 3.3 and 3.4 features observed in the spectra of PPNe but not in those of PNe. Admittedly, however, this has yet to be worked out quantitatively.

Again, if the above implies that an important part of the IS carbonaceous dust is aromatic, then why is the 3.3-$\mu$m band hardly seen in absorption through the DISM (see Schutte et al. 1998; Sellgren 2001), by contrast with the high visibility of the 217 nm band which is also assigned to aromatic material? This may be due to a combination of two factors: a) the intrinsic intensity of this band is so weak that it could not be measured accurately in kerogens; Wexler (1965) puts it at about 1/10 of the aliphatic one; b) for the model material to finally exhibit the very aromatic UIB spectrum of reflection nebulae, H/C must have fallen to about 0.3, i.e. less than 1/4 of the value used above to model the GC 3.4-$\mu$m band.

Finally, both aliphatic and aromatic dusts expelled from CS shells are likely to ultimately settle in dense molecular clouds. How is it, then, that the carbonaceous dust of young stellar objects and dense molecular clouds is so different from that of the GC (Brooke et al. 1999, 1996)? Of course, any refractory dust inside the cloud is bound to be coated with ice. Pendleton (1999) considered the likelihood of the 3.4 feature being hidden by such a coating, but discarded it on the basis of laboratory experiments by Baratta & Strazzulla (1990). Another solution was then proposed by Munoz Caro et al. (2001) and Mennella et al. (2001): they speculate that, somehow, there are enough FAR UV photons and/or cosmic rays inside the cloud for the composite grains to be processed in such a way as to dehyrogenate the carbonaceous component. At the same time, they fully acknowledge the need for theoretical modelling to validate their scenario.

At this point, it may be noted that the astronomical observations of dense clouds do not imply a complete suppression of signatures in the 3.4-$\mu$m range: a broad 3.47-$\mu$m feature is still clearly detected along many sight lines; its red wing even seems to be nearly coincident with that of the diffuse ISM feature (Pendleton 1999). Now, Baratta & Strazzulla's experiments show that a layer 0.1 $\mu$m thick of water ice over an aliphatic organic residue considerably attenuates the 3.4 feature with respect to the ice feature, the former still peaks near 3.4 $\mu$m and has only a shoulder near 3.5 $\mu$m. However, one may speculate that grains in dense clouds do not grow in distinct chemical layers. Another extreme case would be a thoroughly mixed combination of refractory and ice. Because of the very strong dielectric constant of water, the IR spectrum of such a mixture cannot be considered as the sum of the component spectra: the aliphatic feature is likely to be perturbed in shape and position in a non-trivial way. This has to be measured or be computed using, for instance, Maxwell Garnet's theory. Pending such developments, the issue remains open.

The emerging overall picture is that the abundant H and O atoms present in the envelopes of extreme C-stars and C-rich AGB stars favour the formation of highly aliphatic kerogen-like dust (cf. Fig. 5a). Afterwards, in the envelopes of PNe, the temperature and/or integrated photon radiation dose are high enough that aromatisation can nearly be completed (cf. Fig. 5e). However, some of the early dust is liable to escape into the ISM before this is indeed the case. Both aliphatic and aromatic dusts can therefore be observed in the ISM without implying that they are formed there.

More generally, the whole pattern of spectral evolution and diversity is less easily understood in terms of different specific, small, molecules (which cannot evolve spontaneously and continuously) than in terms of a family of materials like kerogens which differ continuously by the relative abundances of a small number of different classes of functional groups attached to a more or less aliphatic skeleton. Moreover, in the latter model, evolution does not require extreme conditions such as shocks, very low or high temperatures, nor high energy photons; ionization is not likely to make a big spectral difference.