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

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4. The mid-IR emission of the discrete peaks

The CVF and filter observations show strong emission peaks which we discuss here. They are ordered by increasing right ascension and named as shown in Fig. 5. This figure shows the LW2 (6.75 µm) contours superimposed to the Digital Sky Survey image of N66. The isolated stars are identified by numbers given in Massey et al. (1989). In Fig. 6 we present the CVF spectra of these peaks. Most of the spectra represent an average of two pixels: spectra of peaks C and E have been obtained averaging four pixels (1 pixel[FORMULA]1.2 pc for the assumed SMC distance). In general the spectra show emission bands and fine structure line on top of a continuum. The wavelengths of the emission bands correspond to those of the Unidentified Infrared Bands already observed before ISO at 6.2, 7.7, 8.6, 11.3 and 12.7 µm (Gillett, Forrest and Merrill 1973, Russell, Soifer and Merrill 1977a, Russell, Soifer and Willner 1977b, Cohen, Tielens and Allamandola 1985, Cohen and Kevin 1989, Jourdain de Muizon et al. 1986, Phillips, Airken and Roche 1984, Roche, Aitken and Smith 1989). They are an universal signature of the ISM in our (Roelfsema et al. 1996, Verstraete et al., 1996, Cesarky et al. 1996a, 1996b, Boulanger et al. 1996, Mattila et al. 1996, Uchida, Sellgren and Werner 1998) and in external galaxies (Boulade et al. 1996, Vigroux et al. 1996, Acosata-Pulido et al. 1996, Metcalfe et al. 1996, Helou et al. 2000). The exact chemical species from which these bands originate have not been yet identified. The best candidates are the Polycyclic Aromatic Hydrocarbons (PAH) (Puget and Lèger 1989), i.e. planar macro-molecules (few hundred atoms) transiently heated by single photon absorption. However, whatever is the exact nature of these carriers, the bands are certainly due to aromatic compounds. For this reason hereafter we will call them Aromatic Infrared Bands (AIBs) carriers. Fig. 6 show the following characteristics:

  • Peaks A, B, C, E, H and I are aligned along the "bar"(Fig. 5). Peak C coincides with the center of the dense star cluster NGC 346. The other peaks are at various distances from this cluster and receive less far-UV radiation except perhaps Peak E. The H[FORMULA] and fine-structure emission line in the direction of Peak C are relatively small, presumably because the gas has been partly expelled by stellar winds from the dense central cluster.

  • Peaks D, F and G lie outside the "bar". F coincides with the compact H II  region N 66A.

[FIGURE] Fig. 5. Map of N 66 in the LW2 filter centered at 6.75 µm (contours), superimposed on the ESO Digital Sky Survey (DSS) image. Coordinates are J2000. Several stars are detected in the LW2 filter: they are HD 5980 (N 346-755) and 2 red stars (N 346-283 and 811: numbers in the catalogue of Massey et al. (1989), see Table 1). The red star N 66-136 is not detected. Some other stars are surrounded by an extended emission: they are blue and probably heat the surrounding interstellar matter.

[FIGURE] Fig. 6. CVF spectra of the 9 main emission peaks in the region of N 66. The peaks are identified on Fig. 5. These spectra have been corrected for zodiacal light as explained in Sect. 2. An estimate for the mean ISRF at 1600 Å normalized to the local ISRF at the same wavelength for each source is given. If dust is mixed with the ionized gas, these values should be decreased by a factor 2.5 (see text for details). The main fine-structure line, the visible H2 line and AIBs (A) are identified in the spectrum of Peak I. All spectra show the [Ne III ] 15.6 µm and [S IV ] 10.5 µm line. The AIBs exhibit a variety of shapes and relative intensities. The broad 10 µm silicate band is seen in emission in the spectrum of Peak C and B and less obviously of Peak F.

The CVF spectra of all these emission peaks show [Ne III ] 15.6 µm and [S IV ] 10.5 µm line emission (Fig. 6).

Even if emission bands are observed at the typical wavelengths of the most intense AIBs (6.2, 7.7, 8.6, 11.3 and 12.8 µm), these are very different in their shape and relative intensities from the AIBs observed in the galactic reflection nebulae, to which hereafter we will refer as the "classical" AIBs.

Peak A shows a broad AIB at 7.7 µm, a 11.3 µm AIB not very intense and faint 12.7 (possibly blended with a [Ne II ] line at 12.8 µm), 13.5 and 14.5 µm bands.

Peak B shows very faint AIBs, if any, and a broad silicate emission at [FORMULA] 10 µm. Note that there are a few faint hot stars in Peak A (N 346-320 and 325), as well as in Peak B (N 346-347, 352, 353 and 357: Massey et al. 1989).

Peak C, in the direction of the center of the young star cluster, has a spectrum very similar to that of Peak B but with a stronger continuum. It exhibits only faint AIBs and a broad 10 µm silicate band is clearly seen in emission. The spectrum of Peak C is discussed in more detail by Contursi et al. (2000).

The spectrum of peak D is characterized by broad emission near 8 µm where the usual AIBs are partly merged. Note the short-wavelength continuum, also seen towards Peaks C and E. This region contains at least 3 hot stars (N 346-466, 469 and 478) the brightest of which is the evolved or reddened N 346-466 (V=15.91, U-B=-0.54:, B-V=0.27, Massey et al. 1989)

Peak E contains the relatively bright, reddened O8V star N 346-549 with V=15.26, U-B=-0.96, B-V=0.22 (Massey et al. 1989). The continuum near 5 µm is the strongest in the whole map (see Fig. 2). It is too strong to be the photospheric emission of the star, but it can be due at least in part to circumstellar dust or to a red companion. The most conspicuous feature in the spectrum of Peak E is a very broad emission feature centered near 7.7 µm in which the usual AIBs are even less identifiable than in the spectrum of peak D. Both the continuum at 5 µm and the presence of the broad band at 7.7 µm are characteristics of AGN spectra like that of Centaurus A (Mirabel et al. 1999). The origin of the 7.7 µm broad feature has not yet been established: it may be due to coal-like grains. However, it is not clear whether these types of grains normally exist in the ISM of galaxies and become visible only when destruction of classic AIBs carriers occurs, or if they form through hard UV photons processing on the classical AIB carriers. The 6.2 and 11.3 µm bands are surprisingly weak. The peculiar appearance of the 7.7 µm brad feature and the faintness of the 11.3 µm band might be due to some amount of silicate absorption, but the [S IV ] line at 10.5 µm, which should also be affected, does not seem particularly weak. Moreover, the presence of a certain amount of silicate absorption cannot explain the weakness of the 6.2 µm AIB. Note also the features at 13.5 and 14.5 µm which can arise from the out-of-plane C-H bending vibrations on aromatic rings with 3 and 4 contiguous H atoms (trio and quarto ).

The spectrum of Peak F (N 66A) shows probable silicate emission and weak AIBs. Peak F contains at least 7 hot stars, the brightest of which is the O5.5V star N 346-593 with V=14.96, U-B=-1.01, B-V=-0.16 (Massey et al. 1989).

Peak G coincides with two hot stars, N 346-628 and 635 (Massey et al. 1989). This peak is on the molecular cloud not associated with the main HII region (Fig. 12). Its spectrum is the closest to the typical Galactic AIB spectra, e.g. those of NGC 7023 (Cesarsky et al. 1996a).

Peak H has faint bands and peak I displays intense AIB bands. Both show a classical AIB spectrum. They contain a few faint hot stars, respectively N 346-640, 641, 648, 654 and N 346-696 and 697 and in fact it has a steep continuum rising toward long wavelength. Moreover, Peak I contains the bright late O or early B star N 346-690 with V=15.70, U-B=-0.75, B-V=0.00 (Massey et al. 1989) and it has the brightest emission in both CO(1-0) and H2 among the MIR peaks (Rubio et al. 2000). The column density in this peak, relative to the others region, is thus sufficiently high to explain the strength of AIBs.

As the AIBs are believed to be excited mainly by far-UV photons in the hard radiation field of N 66, we have built a rough map of the radiation density at 160 nm using the stellar photometry from Massey et al. (1989) (Fig. 7). Details about how we built this map are given in Appendix A. There are two sources of uncertainties in this calculation. 1) Extinction has not been taken into account (except for determining the intrinsic stellar UV flux). Extinction in N 66 is known to be very small for stars (E(B-V)=0.14, Massey et al. 1989) and the Balmer decrement value of 3.05 [FORMULA] 0.15 (Ye et al. 1991) is close to the unreddened value of 2.86. If dust is mixed with the ionized gas, our values for the UV fluxes are upper limits and may be too high by [FORMULA] 1 mag. (a factor 2.5). If dust is outside the ionized gas regions our values are unaffected. 2) The other uncertainty is due to errors in the assignment of the stellar spectral types. However, changing the luminosity class in the most ambiguous cases changes the radiation density by only 30[FORMULA].

[FIGURE] Fig. 7. Projected distribution of the radiation field at 160 nm in the region of N 66 (grey scale), superposed on the LW2 (6.75 µm) contours. Coordinates are J2000. Gray scale values are in Local ISRF units.

The average values of the ISRF at 1600 Å normalized to the local ISRF (LISRF) at the same wavelength (Gondhalekar et al. 1980) are indicated in Fig. 6 and they range from 2 to 9 [FORMULA] 105 the LISRF. They correspond to the values obtained per DSS pixel (=1.7") averaged over a circular area of 2.8 pc radius (= 5.6 DSS pix with an assumed distance for SMC=61 kpc). This is the approximate resolution of the ISO data, thus the same aperture was used to obtained the LW3, LW2 and the 160 nm fluxes reported later in Fig. 14. Note that if dust is mixed with gas inside the HII region, the UV flux values still remain very high, ranging from 5.3 104 (peaks A and I) to 2.5 105 (peak C) times that of the solar neighborhood. In Fig. 6 we have not labeled the ISRF average value of peak G because the new CO(2-1) data show that this cloud and probably the "spur" visible as diffuse emission (see Sect. 5) are not associated with the N 66 bar (Rubio et al. 2000).

From the collection of CVF spectra that we have just discussed, several conclusions can be derived:

  • Silicate emission is clearly visible in Peak C and B and more marginally in Peak F. Interstellar silicate emission has been detected in the Orion nebula and a few other H II  regions, and must be due to relatively big grains (size [FORMULA] 0.01 µm) heated to [FORMULA] 100 K or more, since it is only seen when the radiation field is very high (Cesarsky et al. 2000).

  • In three peaks (C, D and E), there is clear continuum emission at all the studied wavelengths down to the shortest one, 5 µm. While a part of this continuum may be associated with the AIBs, it is clear that they cannot account for all: classic Galactic AIB spectra as those of NGC 7023 or M 17 (Cesarsky et al. 1996a, 1996b) show a negligible contribution of the AIBs at 5 µm. The 5 µm continuum should then be due mostly to grains rather close to the hot stars contained in each of these regions. In Peak C relatively small silicate or other grains can be heated to sufficient temperatures to emit at 5 µm (see Contursi et al. 2000). The radiation field is lower in peaks E and D. But both peaks contain reddened stars, and the emission might be due to dust around these stars.

  • The AIBs emission shows differences in the studied peaks. The AIB spectra are sometimes very different from the "typical" Galactic AIB spectra of e.g. NGC 7023 (Cesarsky et al. 1996a). The 7.7 µm band towards Peak E is much broader and the 8.6 µm band is not visible, perhaps merged into the 7.7 µm feature (but this might be partly due to silicate absorption). Peak D displays an intermediate case. "New" features near 13.5 and 14.5 µm are visible in several spectra. Although faint, these features are likely to be real. Residual from glitches could result in artificial features only for few pixels. Moreover, the same bands are also visible in other regions like M 17 and NGC 7023 (Cesarky et al. 1996a, 1996b, Klein et al. 1999). Fig. 6 shows that there is a spectral evolution from peak C, in the center of the star cluster, to peaks F, I, D and G, where the AIBs are stronger with respect to the continuum and more similar to the "classical" Galactic AIBs. This suggests that UV radiation has a crucial role on the grain processing. It destroys the classical AIB carriers, favoring the broader-band emission of relatively big carbonaceous grains which are heated to sufficient temperatures (peak E) or of smaller carbonaceous grains heated transiently by absorption of single photons. There might be transformations from PAH-like 2-D molecules responsible for the classical AIBs into 3-D grains or vice-versa. The broad band emission of peak E might indicate phenomena occurring close to the reddened O stars contained in these regions.

  • The strongest AIB/continuum ratio and the most "classical" AIB spectrum is observed towards Peak G to the North of the CVF map. This peak is at the southern edge of a spur well visible in the LW2 filter images (5.0-8.0 µm) of Fig. 5. Since this spur and the molecular cloud (Fig. 9) are not associated with the N 66 star cluster (Rubio et al. 2000), grains here are heated by an ISRF lower than in the bar, providing a spectrum more similar to those observed in relatively quiescent regions of our Galaxy.

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Online publication: October 30, 19100
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