Astron. Astrophys. 358, 708-716 (2000)

4. Tracing the silicate emission

To delineate the spatial extent of the 10 µm -silicate emission conspicuously visible in Figs. 3 and 4, we proceed as follows. We start with the spectrum towards  Ori A, which shows the most conspicuous silicate emission and we represent the AIBs by Lorentz profiles, see Fig. 6 (top). Next we subtract them from the CVF spectra. The remaining continuum has the generic shape of a blackbody on top of which we see the broad bands corresponding to the silicate emission, Fig. 6 (middle). Finally, we subtract a second order polynomial from the continuum thus obtaining the well known silicate emission profile at that position, Fig. 6 (bottom). The profile thus obtained is then used as a scalable template to estimate the emission elsewhere, see Appendix C for more details.

Fig. 6. Top panel: CVF spectrum towards  Ori A (solid line). The ordinates give fluxes in Jy per pixel. A Lorentz fit to the AIBs is shown as the dotted line. Middle panel: result of the subtraction of the AIBs from the CVF spectrum. The fit to the continuum is shown by the dotted line. Bottom panel: Residual from the middle figure, i.e. the suspected amorphous silicate emission profile; notice the narrower bump near 9.6 µm .

On top of the broad band of amorphous silicate centered near 9.7 µm   we see a band centered at nearly 9.6 µm , which we ascribe to crystalline silicates (Jaeger et al.  1998). This band was also used as a scalable template as explained above and in Appendix C. Finally, the S(5) rotation line of H₂ at 6.91 µm  is present and is probably blended with the [Ar II ] line at 6.99 µm .

In Fig. 7, we see that the spatial distribution of the 9.7 µm -feature of amorphous silicate is quite similar to that of the 15.5 µm -continuum. The 15.5 µm  continuum emission includes a strong contribution from silicates (see Fig. 4), but a peak in the silicate emission around  Ori A is also evident. The silicate emission is thus predominantly due to larger grains. The narrower 9.6 µm  feature is mapped in Fig. 8. We note its similarity to the distribution of the 9.7 µm  broad band: this fact lends support to our assignation of this band to crystalline silicate.

Fig. 7. Map of the intensity of the broad 9.7 µm  band of amorphous silicates (contours) superimposed on the 15.5 µm  continuum map (grey scale). Note the bright silicate emission around  Ori A (cross). The contours correspond to integrated band intensities from 0.25 to 0.7 erg s- 1 cm-2 sr-1 by steps of 0.05; the gray image spans from 1 to 80 Jy/pixel.

Fig. 8. Map of the 9.6 micron feature map (contours) superimposed on the map of the broad 9.7 µm  band of amorphous silicates (grey scale spanning 0.1 to 10 erg s- 1 cm-2 sr-1. The contours correspond to integrated band intensities from 0.02 to 0.11 erg s-1 cm- 2 sr-1 by steps of 0.001. The shift with respect to the position of  Ori A (cross) is by less than one pixel and may not be significant.

Due to the low spectral resolution of the CAM-CVF, however, the 9.6 µm  feature will certainly blend with the S(3) pure rotational line of molecular hydrogen - if present. To check this we have compared our 9.6 µm  map to that of molecular hydrogen in its fluorescent vibrational line 1 S(1) (2.12 µm). Courtesy of P.P. van der Werf (van der Werf et al.  1996), we reproduce in Fig. 9 the map of the fluorescent molecular hydrogen emission. This latter correlates better with the AIB emission as traced by the 6.2 µm -feature (bottom figure) than it does with the tentative crystalline silicate emission (top), namely they both peak along the bar. This is not surprising because the H₂ and AIB emitters require shielding from far-UV radiation to survive. Conversely, the 9.6 µm  silicate feature is stronger where H₂ is weak as can be seen around  Ori A. In addition, the H₂ S(3) rotational line at 9.66 µm  is detected in the ISO-SWS spectrum of the Orion bar presented in Verstraete et al.  (1999, in preparation) with an intensity of W m^-2 sr^-1. This value is a factor of 16 below the median flux of the 9.6 µm  feature in our map, namely W m^-2 sr^-1. We can thus safely conclude that our 9.6 µm -emission predominantly originate from silicates. A confirmation of the identification of the 9.6 µm  band with a crystalline silicate dust component would be possible if a second signature band were seen in our spectra. The SWS spectrum (Fig. 4) shows only broad emission bands that are difficult to characterise, snd additionally, the chacteristic crystalline olivine band in the 11.2-11.4 µm  region (e.g.  Jaeger et al.  1998), if present, is blended with the 11.2 µm  aromatic hydrocarbon feature. Additionally, most of the chacteristic crystalline bands fall longward of the CVF spectra. Thus, it is difficult to self-consistently confirm the 9.6 µm  band identification with the presented data.

Fig. 9. The 9.6 µm  feature (top) and 6.2 µm -AIB (bottom) both in contours superimposed to the S(1) line emission of molecular hydrogen taken from van der Werf et al.  (1996) (grey scale). The contours correspond to integrated band intensities from 0.02 to 0.11 by steps of 0.01 (top figure) and 0.045 to 0.27 by steps of 0.025 (bottom figure) in units of erg s-1 cm- 2 sr-1.

In summary, emission in the 9.7 µm  band of amorphous silicate emission exists everywhere inside the Orion H II  region. Previously, amorphous silicate emission had only been seen in the direction of the Trapezium (Stein & Gillett 1969; Forrest et al.  1975; Gehrz et al.  1975). We may assume that the 18 µm  band is also widely present in the region, as witnessed by the single SWS spectrum (Fig. 4) and by the generally rising long wavelength end of ISOCAM spectra; the two spectra shown, Figs. 3 and 10 are quite representative of the steeply rising continuum longward of 15 µm .

Fig. 10. CVF spectrum towards Ori A (heavy solid line). The ordinates give fluxes in Jy per pixel. The fit to the continuum is shown by the thin solid line. The fit (see Sect. 3) to these data comprises, from top to bottom on the right-hand axis: 100-K amorphous astronomical silicate (dot-dashed line), 110-K amorphous carbon (dashed line), 235-K amorphous astronomical silicate (dot-dashed line) and 330-K amorphous carbon emission (dashed line).

4.1. The interstellar silicate and H₂ emission around  Ori A

The case of  Ori A is particularly interesting because the geometry is simple and therefore allows quantitative calculations. Moreover, the thermal radio continuum, the recombination lines and the fine-structure lines are faint in the neighbourhood of this star (Felli et al.  1993; Pogge et al.  1992; Marconi et al.  1998, and the present paper, Fig. 2).  Ori A is classified as an O9.5Vpe star and shows emission lines (see e.g. Weaver & Torres-Dodgen 1997). It is a spectroscopic binary and an X-ray source. There is little gas left around the star and the observed silicate dust (Fig. 8) is almost all that is visible of the interstellar material left over after its formation. Indeed, O stars are not known to produce dust in their winds which are probably much too hot, so that the silicates we see here must be of interstellar origin.

The mid-IR continuum observed towards  Ori A can be accounted for by combining emission of warm silicate and carbon grains (see Fig. 10). The model continuum was obtained in the same way as for the SWS observation (see Fig. 4) and with the same assumptions. The grain temperatures are consistent with the heating of interstellar grains by the strong radiation field of the star.

As discussed above, the band near 9.6 µm (Fig. 6 bottom and Fig. 8) may be due to crystalline silicates, any contribution of the S(3) H₂ line to this band is minor. Another band at 14 µm (see Figs. 4 and 10) might also be due to crystalline silicates. Amongst the crystalline silicates whose mid-IR absorption spectra are shown by Jaeger et al.  (1998), synthetic enstatite (a form of pyroxene) might perhaps match the Ori A spectrum. The interest in the possible presence of crystalline silicates around this star is that they would almost certainly be of interstellar origin, pre-dating the formation of the star. Observations at longer wavelengths are needed for a definitive check of the existence of crystalline silicates and for confirming their nature. Such observations do not exist in the ISO archives and should be obtained by a future space telescope facility.

© European Southern Observatory (ESO) 2000

Online publication: June 8, 2000
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