Astron. Astrophys. 362, 310-324 (2000)

6. Discussion and conclusions

We have presented CVF and multi-filter ISO observations of the region of N 66 in the SMC. They reveal a wide variety of phenomena that are not always easy to interpret. The following results have been obtained:

i) Emission in the fine structure line [Ne III ] 15.6 µm and [S IV ] 10.5 µm is present throughout the region. These line are very strong compared to line from singly-ionized ions like [Ne II ], due to excitation by the very hot stars of N 66. There are considerable differences between the space distributions of the [Ne III ] and [S IV ] line, that we attribute to density effects in the photoionized regions and to shock excitation in a supernova remnant.

ii) AIB emission is generally weak but present in many places of the field. This general weakness, already noted by Sauvage et al. (1990), can be related to the low carbon abundance in the SMC, which is 14-20 times smaller than in our Galaxy (Pagel 1993; Garnett et al. 1995). However, the mid-IR spectrum of a quiescent region in the SMC (Reach et al. 2000) is similar to the galactic "cirrus" emission, suggesting that the differences observed in N66 are principally produced by the extremely high and hard ISRF. Analogous ISOCAM observations of another HII region in the SMC (SMCB1), not yet completely reduced, will help us to clarify which parameters affect the dust properties. There is AIB emission probably coming from the surface of a molecular cloud, like in the region of N 4 in the LMC (Contursi et al. 1998). Most of the AIB spectra we have obtained (see examples in Fig. 6) are different from the classical "Galactic" AIB spectra, e.g. those of the reflection nebula NGC 7023 (Cesarsky et al. 1996a). The 7.7 µm feature is almost always broader than in the ISM of our Galaxy and the 8.6 µm AIB is not always visible (merged with the 7.7 µm feature?). The 11.3 µm band can be strong with respect to the other AIBs, but it can also be quite weak. Also interesting is the fact that the bands at 13.5 and 14.5 µm seem stronger than what is observed in the Galaxy. A similar wide variety of AIB spectra is seen in compact or ultra-compact H II  regions (Cesarsky et al. 1996b, Roelfsema et al. 1996).

This variety is presumably due to the co-existence of several forms of AIB carriers, one of which dominates depending on the conditions. The spectrum in the direction of Peak G (close to the edge of the molecular cloud) is not conspicuously different from the classical Galactic AIB spectrum. This cloud is not associated with the N 66 bar. However the 11.3 µm band is somewhat stronger and overall the spectrum is very similar to the spectrum of the molecular cloud M 17N (Henning et al. 1998). The spectrum of peak E is similar to that of very small 3-D carbonaceous grains like semi-anthracite, which also reproduces well the spectra of a few Galactic proto-planetary nebulae (Guillois et al. 1996). None of our spectra matches well that of nanoparticules produced by laser pyrolysis of hydrocarbon (Herlin et al. 1998), at least in the 11-14 µm spectral range. Together with other data, our data will allow the study of the dependence of the AIB spectra on the far-UV radiation density and spectral hardness. Maps of the radiation field like the one presented in Fig. 7 will be useful for such studies.

iii) The ISRF in the bar of N66 is at least 10⁵ times the local ISRF at the same wavelength. We have evidence that such strong and hard ISRFs are able to significantly destroy AIB carriers and to a lesser extent also the VSGs.

iv) Aside from the carbonaceous grains just discussed, our observations show continuum emission by silicate grains at several MIR peaks.

Our observations also shed light on the evolution of the N 66 region. Its general optical and mid-IR morphology (Fig. 5) suggests that star formation has arisen in an arc of material compressed by shocks, probably caused by previous supernovae explosions. There are many examples of similar phenomena in both the SMC and the LMC: HI bubbles (Staveley-Smith et al. 1997, Kim et al. 1998), secondary star formation on the edge of these super bubbles (Parker et al. 1992). This secondary star formation itself is probably not coeval. In fact, peak C contains only unredenned OB stars and it is more evolved than other peaks. Peaks E, H and I for example, have reddened stars suggesting that the surrounding material has not been yet spread out. Following the model proposed by Elmegreen (Elmegreen 1995), we suggest that star formation along the N66 bar has taken place in a sequential way, starting from the OB stars associated to the peak C. The spatial separation between the sub-groups of stars associated to the different MIR peaks is 8-10 pc, comparable to that predicted by numerical simulations (Elmegreen 1995). Also the nature of dust changes along the bar (Fig. 6) probably because it forms at different times and in different environments. Some of the spectra presented in Fig. 6 show weak S(0)(9.6 µm) and perhaps S(3) (7.0 µm) line of H₂ in emission. This indicates that some molecular gas is still present in the HII region. New CO(2-1) data seem also support this scenario and show that the MIR peaks correspond to molecular clumps with different velocities (Rubio et al. 2000).

Several processes can explain the variety in the observed AIBs strengths and shapes. The AIBs faintness can be ascribed to a significant destruction of their carriers, either by the harsh ISRF or by shocks produced by stellar winds. Broader than normal AIBs can arise from different grains excited by the ISRF of N66; their emission can be generally hidden by the AIBs where these carriers are not significantly destroyed. Other possibilities are that the original grain size distribution is modified by photo-processing on grains, grains shattering, or that the grain composition was originally different due to the low metallicity of SMC. However, this last hypothesis seems to be discarded from the presence of classical Galactic AIBs in a quiscent region of SMC. Finally, the observed variation in the MIR spectra can be related to the the fact that the dust formed at different times and in different environments according to the idea that in N66 the star formation evolved in a sequential way.

The results obtained from the analysis of ISO observations of nearby objects (HII regions, PDRs, molecular clouds, etc.) are useful to understand the dust properties of the more distant galaxies. It is thus important to give the global properties of these nearby regions on scale lengths comparable with the ISOCAM resolution of at least moderately distant galaxies. As LW2 and LW3 are by far the most used filters for ISOCAM observations of external galaxies, we give in Table 1 the 15/6.75 µm (LW3/LW2) ratio of N 66 for two different fields of view: 3´3´ (CVF field of view) and 7.8´7.8´ (broad band images), with their corresponding physical sizes. This ratio has been calculated directly on the broad-band images. For comparison we also give in Table 1 the same results for the HII region N 4 in the LMC, for which ISOCAM broad-band data have been already published by Contursi et al. (1998). We also show on Fig. 17 the spectrum of N 66 averaged over the total CVF field of view. Note that N 4 is bathed in a ISRF 100 times lower than the N 66 ISRF. At a scale length of 100 pc there are no significant differences in the mid-IR colors of these HII regions despite their different ISRF. Moreover, on the 100 pc scale, the LW3/LW2 ratios of both N 4 and N 66 are not very different from those found in regions bathed in ISRFs similar to the Local one. The ISOCAM resolution of 7" corresponds to a 100 pc region for an object at a distance of 3 Mpc, that of galaxies in the Sculptor group for example.

Fig. 17. The average zodiacal-light subtracted CVF spectrum of N66 averaged over the whole CVF (3´3´) field, corresponding to a physical size of 53 pc 53 pc.

Table 1. The global 15/6.75 µm ratios integrated over the CVF and filter fields of view (3´3´ and 7´7´) of the N66 (SMC) and N4 (LMC) HII regions

© European Southern Observatory (ESO) 2000

Online publication: October 30, 19100
helpdesk.link@springer.de