A&A 381, 571-605 (2002)
E.Peeters1,2 - N. L. Martín-Hernández2 - F. Damour3 - P. Cox4 - P. R. Roelfsema1 - J.-P. Baluteau3 - A. G. G. M. Tielens1,2 - E. Churchwell5 - M. F. Kessler6 - J. S. Mathis5 - C. Morisset3 - D. Schaerer7
1 - SRON-Groningen, PO Box 800, 9700 AV Groningen, The Netherlands
2 - Kapteyn Institute, PO Box 800, 9700 AV Groningen, The Netherlands
3 - Laboratoire d'Astrophysique de Marseille, CNRS & Univ. de Provence, BP 8, 13376 Marseille Cedex 12, France
4 - Institut d'Astrophysique Spatiale, Bât. 121, Université de Paris XI, 91405 Orsay, France
5 - Department of Astronomy, 475 North Charter Street, University of Wisconsin, Madison, WI 53706, USA
6 - ISO Data Centre, Astrophysics Division, ESA, Villafranca, Spain
7 - Laboratoire d'Astrophysique, Observatoire Midi-Pyrénées, 14, Av. E.Belin, 31400 Toulouse, France
Received 3 August 2001 / Accepted 19 October 2001
Infrared spectra between 2.3 and 196 were taken towards a sample of 45 compact Hii regions using the two spectrometers (SWS and LWS) on board ISO. The primary goal is to determine the distribution of element abundances in the Galaxy, although there are also many other uses of this database. The spectra contain a wealth of information on the ionized gas and the associated photodissociation regions through the atomic fine-structure lines and on the dust properties via the dust emission bands and the continuum. Significant variations are found from source to source in both spectral shape and content. The sample of Hii regions spans a wide range in galactocentric distance (from 0 to 22 kpc) enabling to investigate the variations of the nebular properties across the Galactic plane. The observations and the data reduction are described in detail in the present paper. The ISO spectral catalogue of compact Hii regions contains the combined SWS-LWS spectra for each of the sources, the fluxes of the atomic fine-structure lines and hydrogen recombination lines, and an inventory of the spectra in terms of molecular lines, dust and ice bands.
Key words: catalogues - Hii regions - Galaxy: abundances - infrared: ISM: lines and bands - infrared: ISM: continuum
The two spectrometers on board the Infrared Space Observatory (ISO, Kessler et al. 1996), the Short Wavelength Spectrometer (SWS, de Graauw et al. 1996) and the Long Wavelength Spectrometer (LWS, Clegg et al. 1996), provided a unique opportunity to measure spectra from 2.5 to 196 with resolving powers between 150 and 500, towards 45 Galactic Hii regions. The spectral coverage gives access to nearly all the atomic fine-structure and hydrogen recombination lines in the infrared range. In addition to the atomic emission lines, the ISO spectra reveal the shape of the strong dust continuum and several emission features. Absorption bands from molecular ice species were detected toward some sources. The observed Hii regions were carefully selected to sample the Galactic disk out to 22 kpc from the center, permitting a study of Galactic abundance variations over a larger range than in previous studies. The information contained in the combined SWS-LWS grating spectra provides a basis for a broad range of studies including the properties of Hii regions, the characteristics of associated dust and the abundance gradients in the Galaxy.
This Paper presents for each source the combined SWS-LWS spectrum, the measured line fluxes for all the atomic emission lines and the inventory of the detected molecular lines and the dust and ice bands. An accompanying Paper (Martín-Hernández et al. 2002, hereafter PaperII) will discuss the first results of the properties of the compact Hii regions and the element abundance as derived from the present data.
The sample of compact Hii regions with their observations are discussed in Sect. 2. Section 3 gives a detailed description of the reduction of the SWS and LWS data; the line flux determination, the error determination and the influence of the source structure on the observed fluxes are discussed in Sect. 4. Section 5 compares the ISO data with corresponding KAO and IRAS observations. Characteristics of the Hii regions such as the kinematical distance, the luminosity and the morphology are described in Sect. 6. The ISO spectra, the line fluxes, the detected molecular lines and dust and ice bands are presented and discussed in Sect. 7 for each Hii region of the sample. This section contains the core of the paper. A summary is given in Sect. 8.
|IRAS 01045+6506||01 07 50.7||+65 21 21.7||18-Mar.-98||85303602||-||G124.64+2.54||WB 380|
|IRAS 01420+6401||01 45 39.6||+64 16 02.1||21-Jul.-97||61301076||61301075||G128.77+2.01||WB 399|
|IRAS 02219+6125||02 25 44.6||+62 06 11.3||23-Aug.-97||64600609||-||G133.70+1.20||W3 A, RAFGL 326|
|IRAS 02383+6241||02 42 19.8||+62 53 51.8||3-Mar.-98||83901404||83901403||G135.18+2.69||WB 436|
|IRAS 02575+6017||03 01 31.3||+60 29 13.5||27-Mar.-98||86300968||86300969||G138.29+1.55||WB 463, RAFGL 4029|
|IRAS 04025+5313||04 06 25.5||+53 21 50.0||26-Sep.-97||68100312||68100305||G149.59+0.90||WB 529|
|IRAS 05167+3858||05 20 11.1||+39 01 19.7||22-Aug.-97||64501216||64501206||G168.68+1.08||WB 625|
|IRAS 05221+4139||05 25 39.8||+41 41 50.3||22-Aug.-97||64501104||64501107||G167.06+3.46||WB 640|
|IRAS 05302+3739||05 33 40.1||+37 41 07.3||22-Aug.-97||-||64501371||G171.26+2.54||WB 656|
|IRAS 05335+3609||05 36 52.6||+36 11 00.3||8-Oct.-97||-||69201373||G172.87+2.26||WB 668|
|IRAS 06158+1517||06 18 44.8||+15 16 43.4||13-Mar.-98||84901804||84901903||G195.65-0.10||Sh 266, WB 794|
|IRAS 10589-6034||11 00 59.8||-60 50 27.1||11-Aug.-96||26800760||26800723||G289.88-0.79||RAFGL 4122|
|IRAS 11143-6113||11 16 33.8||-61 29 59.4||5-Aug.-96||26200509||26200510||G291.86-0.68||RAFGL 4127|
|IRAS 12063-6259||12 09 01.1||-63 15 54.7||2-Aug.-96||25901414||60900713||G298.19-0.78||RAFGL 4144|
|IRAS 12073-6233||12 10 00.3||-62 49 56.5||2-Aug.-96||25901572||25901573||G298.23-0.33|
|IRAS 12331-6134||12 36 01.9||-61 51 03.9||10-Sep.-96||29900470||29900475||G301.11+0.97|
|IRAS 15384-5348||15 42 17.1||-53 58 31.5||10-Sep.-96||29900661||29900628||G326.44+0.91|
|IRAS 15502-5302||15 54 06.0||-53 11 36.4||5-Aug.-97||27301117||62802718||G328.31+0.43|
|IRAS 16128-5109||16 16 39.3||-51 16 58.3||6-Sep.-96||29402233||29402234||G332.15-0.45|
|IRAS 17160-3707||17 19 26.1||-37 10 53.8||6-Oct.-96||32400821||32400822||G350.10+0.09|
|IRAS 17221-3619||17 25 31.7||-36 21 53.5||12-Oct.-96||33100380||33100376||G351.46-0.44|
|IRAS 17279-3350||17 31 18.0||-33 52 49.4||3-Oct.-96||32200877||32200878||G354.20-0.05||RAFGL 5347|
|Sgr C||17 44 35.6||-29 27 29.3||5-Mar.-98||84100301||-||G359.43-0.08|
|17 44 35.9||-29 27 54.3||11-Mar.-98||-||84700220|
|IRAS 17455-2800||17 48 41.5||-28 01 38.3||29-Aug.-96||28701327||28701328||G1.13-0.11||Sgr D|
|IRAS 17591-2228||18 02 13.2||-22 27 58.9||14-Apr.-97||51500580||51500579||G7.47+0.06|
|IRAS 18032-2032||18 06 13.9||-20 31 43.2||14-Apr.-97||51500478||51500477||G9.61+0.20B||RAFGL 5436|
|IRAS 18116-1646||18 14 35.2||-16 45 20.6||13-Apr.-96||14801733||14801732||G13.88+0.28|
|IRAS 18162-2048||18 19 12.0||-20 47 31.1||13-Apr.-96||14802136||14802135||G10.84-2.59||RAFGL 2121|
|IRAS 18317-0757||18 34 24.9||-07 54 47.9||8-Mar.-97||47801040||47801041||G23.96+0.15||RAFGL 2194|
|IRAS 18434-0242||18 46 04.0||-02 39 20.5||17-Apr.-96||15201383||15201381||G29.96-0.02||RAFGL 2245|
|IRAS 18469-0132||18 49 33.0||-01 29 03.7||26-Oct.-97||71100888||71100887||G31.40-0.26|
|IRAS 18479-0005||18 50 30.8||-00 01 59.4||17-Apr.-96||15201791||15201792||G32.80+0.19||RAFGL 5536|
|IRAS 18502+0051||18 52 50.2||+00 55 27.6||17-Apr.-96||15201645||15201645||G33.91+0.11||RAFGL 5541|
|IRAS 19207+1410||19 23 02.4||+14 16 40.6||15-Apr.-96||15001041||15001039||G49.20-0.35||RAFGL 2379|
|IRAS 19442+2427||19 46 20.1||+24 35 29.4||15-Apr.-96||15000444||15000443||G60.90-0.10||RAFGL 2454, Sh 87|
|IRAS 19598+3324||20 01 45.6||+33 32 43.7||11-Apr.-96||14601350||14601348||G70.29+1.60||K3-50 A|
|DR 21||20 39 00.9||+42 19 41.9||17-Apr.-96||15200555||15200786||G81.70+0.54||RAFGL 2624|
|IRAS 21190+5140||21 20 44.9||+51 53 26.5||24-Apr.-96||15901853||15901854||G93.53+1.47||M1-78|
|IRAS 21270+5423||21 28 41.9||+54 36 51.5||13-Feb.-98||82100309||82100310||G96.29+2.59||Sh 127A, WB 85A|
|IRAS 21306+5540||21 32 11.4||+55 53 23.9||15-Feb.-98||82301012||82301011||G97.52+3.18||Sh 128, WB 91|
|IRAS 22308+5812||22 32 45.9||+58 28 21.0||12-May-96||17701258||17701257||G105.63-0.34||Sh 138, WB 191|
|IRAS 23030+5958||23 05 10.6||+60 14 40.6||24-Jun.-96||22000961||22000962||G110.10+0.05||Sh 156, WB 240|
|IRAS 23133+6050||23 15 31.4||+61 07 08.5||24-Jun.-96||22001506||22001505||G111.62+0.37||Sh 159, WB 261|
The Hii regions were selected primarily to be observable with the SWS and the LWS. First, the sources were required to be bright enough (at least a few Jy in the IRAS 12 band) to ensure detectability with both instruments. Second, the nebulae had to be compact with sizes smaller than the SWS and LWS beams (i.e. ) so that all the nebular flux is within both beams. Six sources are point-like for the ISO spectrometers, i.e. with mid-infrared sizes . Third, to minimize confusion problems, isolated regions were prefered over sources lying in complexes. Finally, the sources were chosen to cover as wide a distribution as possible in galactocentric distance in order to derive the Galactic element abundance over larger scales than in previous studies. ISO was pointed at the IRAS positions of the selected IRAS sources. The sample was enlarged by DR 21 and Sgr C.
For two of the sources, IRAS 18316-0602 (also known as GL7009) and IRAS 19110+1045, the 2.3 to 196 spectra were found to be dominated by absorption and they are not considered in this catalogue. A detailed analysis of their spectral content can be found in Dartois et al. (1998). The sample considered for this catalogue consists therefore of 43 Hii regions which are listed in Table 1. For each source, the name, the equatorial coordinates of the position observed with the SWS and LWS, the date of the observation, the Target Dedicated Time (TDT) of the observations (which uniquely identifies each observation) and some other common designations are given in Table 1. For the majority of the sources, the SWS and LWS observations were concatenated, i.e. the LWS spectrum was measured immediately after the SWS observation was completed. For Sgr C, only an SWS spectrum was taken in the UCHii ISO program and the LWS full grating spectrum was extracted from the ISO archive.
The SWS grating spectra were taken in the so-called Astronomical Observation Templates 01 (AOT 01) full scan mode (de Graauw et al. 1996) speed 2. The SWS has two spectrometer sections: the short wavelength section (SW) covering 2.3 to 12 , and the long wavelength section (LW), from 11 to 45 , yielding a total wavelength coverage of 2.3-45 . The combined spectrometer sections are composed of 12 separate grating bands, the AOT bands (see Table 2). Each AOT band is defined through a combination of detector array ( detectors), instrument aperture and grating order. The SWS has three apertures with sizes of for bands 1 and 2, of for bands 3A, 3C and 3D, of for band 3E and of for band 4.
In the AOT 01 mode, spectra are taken by scanning the grating over its full mobility range and back. Thus, for each detector of a detector array and each AOT band two independent scans are obtained; one "up'' and one "down'' scan. As a result, 24 independent measurements of the incident spectrum are obtained (i.e. 2 scans 12 detectors). Due to the sampling strategy of the SWS AOT 01 mode, each of these independent measurements individually is under-sampled. However, because the measurements are obtained at slightly different wavelength offsets they can be combined to yield a factor of four oversampled spectrum. The spectrum nominally has a spectral resolving power of and a wavelength accuracy of about 10% of a resolution element. By comparing the independently reduced up and down scans for a given source, an estimate of the instrumental uncertainty can be obtained.
For calibration purposes, the compact Hii region K3-50A has
also been observed with the SWS in the AOT 01 mode speed 4 yielding a
spectral resolving power of 1600. The results on K3-50A
reported in the catalogue refer to the AOT 01 mode speed 4 data (TDT =
|------- SWS -------||------ LWS ------|
The LWS spectra were measured using the LWS AOT 01 mode (Clegg et al. 1996). The LWS has two spectrometer sections: the short wavelength section (SW) covering 43 to 93 , and the long wavelength section (LW), from 84 to 196 , yielding a total wavelength coverage of 43-196 . Each spectrum consists of ten overlapping sub-spectra, one for each of the ten detectors. The LWS field of view is restricted by an internal field mirror to provide an approximately circular footprint on the sky with a wavelength dependent diameter of 80 (for the exact diameter see the ISO Handbook volume IV: "LWS - The Long Wavelength Spectrometer'' by Gry et al. 2000). The official wavelength ranges are given in Table 2 for the ten LWS detectors.
In the AOT 01 mode, spectra are taken by scanning the grating 3 times over its full mobility range and back. For each target, 6 scans are obtained, 3 in "up'' and 3 in "down'' direction. Combined these scans yield a spectrum sampled at 1/4 of a spectral resolution element. The resolution element is 0.283 for detectors SW1-SW5 (43-93 ) in second spectral order and for detectors LW1-LW5 (84-196 ) in first order, it is 0.584 . The resolving power, , of the LWS grating spectra varies from 140 to 330. Analoguous to SWS, an estimate of the instrumental uncertainty can be obtained by comparing the independently reduced up and down scans for a given source.
For a general description of the SWS instrument and its reduction, we refer to the ISO Handbook volume VI: "SWS - The Short Wavelength Spectrometer'' by Leech et al. (2001). The standard reduction process consist of corrections for the electronics of the system, dark current subtraction, correction for the relative spectral response of the instrument (RSRF) and wavelength and flux calibration. Finally, the wavelengths are corrected for the ISO velocity resulting in an Auto Analysis Result (AAR).
In the following sections, the additional reduction steps applied to the data presented here, the final calibration errors and the unresolved problems are described.
The flux calibration files were updated while reducing the sample sources (corresponding to the final post-mission calibration, OLP10; Shipman & Lahuis 2000). In order to get the final line fluxes, we multiplied the spectra and the measured line fluxes by a factor corresponding to the ratio of the used to the new flux scale calibration factors.
Within filter elements or within the detector material - especially for the band 3 detectors - multiple reflections between the element front and back surfaces may occur. The resulting interference fringes appear as a high frequency modulation of the detector output signal. To some extent all detector bands suffer from fringes. Bands 1, 2 and 4 are only lightly affected, whereas the fringes are more pronounced in band 3. These fringes can be corrected for by fitting cosine functions to the data in wavenumber space and dividing the observed fluxes by those functions. This procedure was applied for bands 3C, 3D and 3E, where the fringes are resolved. Only for the high resolution data of K3-50A fringes are also resolved and hence removed in band 3A.
The final reduction steps are applied at the AAR level. Strong glitches left in the spectra are removed manually (see above). The signals of the 12 detectors of a given AOT band are then all flat-fielded to the average level. Subsequently, for each resolution bin, the median and the standard deviation are calculated. Each point of which the deviation from the mean is more than 2.5 times the standard deviation is removed when situated in line-free spectral regions. Remaining bad data points are manually removed. Finally, all detector data are rebinned to a wavelength grid with a spectral resolution of 450. In case of sources with no or low flux, the reduction process is simplified. Since the deglitching method only influences the flux of strong lines, deglitching is applied to these sources. Furthermore, for each resolution bin, the median and the standard deviation are calculated. Each point that deviates from the mean by more than 2.5 times the standard deviation is removed when situated in all spectral regions.
The final absolute flux scale accuracies for the calibrated SWS spectra are given in Table 2. For each AOT band the 1 accuracy is given including the SWS absolute flux calibration error and the uncertainties in the synthetic stellar models used in establishing the absolute flux scale. The intrinsic wavelength scale accuracy for SWS was established early on in the ISO mission to be better than 5000 (Valentijn et al. 1996). For the resolution of 450 of the data presented here, this corresponds to of a resolution element. However due to the limited accuracy of the satellite pointing and the offset of the given coordinates from the source peak position in some cases, the sources are not always centred in the SWS aperture. Since the observed wavelength directly depends on the angle of incidence of the radiation on the grating, an offset position translates into a slight wavelength shift. Thus, satellite pointing errors lead to an increase in the wavelength error resulting in an overall wavelength accuracy of about 1/3 to 1/6 of a spectral resolution element.
For SWS, a few problems still remain in the calibration. These problems and their consequences for the data presented here are discussed in the following paragraphs.
During the calibration process, the RSRF was fine-tuned to fit three narrow features in band 2C (9-9.3 , 10.1 and 11.0 ) and one feature in band 3A (12.3 ) which are believed to be instrumental since they were seen in all sources with the same strength and width. Although corrected, the shape of the RSRF in these regions is still less reliable than in the rest of the AOT band.
These two problems can lead to the appearances of spurious features which are sometimes mistaken for dust emission features (e.g. the 22 feature, Chan & Onaka 2000).
A sinusoidal wave is present in the reduced spectrum of K3-50A in band 3A (see Fig. 8), as well as in the raw data. The origin of this wave is unclear to us.
The observation with TDT number 64600609 is plotted in the SWS region for W3 since the band overlaps are slightly better in this spectrum. For IRAS 22308, the SWS spectrum with TDT number 17701257 is shown, the LWS part has TDT number 17701258. Again, for the SWS part, this spectrum has a slightly higher quality. For LWS, we took the observation obtained immediately after the SWS observation plotted here. When comparing these two observations of the same source in the SWS wavelength regime, differences can be found since the orientation of the aperture is not the same and the source is extended. For the LWS regime, the aperture position and orientation is the same for both observations.
The LWS observations are calibrated using the pipeline version 8.7. For a detailed description of the instrument and its reduction we refer to Swinyard et al. (1998) and the ISO Handbook volume IV: LWS - The Long Wavelength Spectrometer by Gry et al. (2000).
LWS data was calibrated following the standard reduction steps, which basically include: first-level deglitching, responsity drift and absolute response correction, dark current substraction and finally, wavelength and flux calibration. These processes yield the Auto Analysis Result (AAR).
The dark signal subtraction is a crucial step for very faint sources because a wrong subtraction can lead to negative photocurrent values and to large differences in the flux levels between the sub-spectra. Therefore, the faint sources were reprocessed with the LWS Interactive Analysis Software (LIA) in order to refine this process. LIA also allows us to redo the drift and absolute responsivity corrections, but these tools were not used as their improvement over the pipeline was found to be small.
The post-pipeline analysis of the AAR was performed with the ISO Spectroscopy Analysis Package (ISAP). First of all, the deglitching done in the pipeline removes a large part of the glitches, but some of them or their tails can remain in the data. Furthermore, some glitches may cause latency effects on the detector response which may cause part of a scan for one detector to be higher than the others for some period of time. ISAP permits the removal of these bad data. This step is followed by taking the average of all the 6 scans for every detector and the correction for fringes. Fringes are caused by the interference of two beams (one beam arising from the normal reflection from the LWS field mirror and another from the support structure holding this mirror) which propagate along the instrument with a time delay between them. All detectors are in principle influenced, but the long wavelength detectors are most affected. Fringes occur in the spectra when the source is extended (i.e. > ) and/or structured, or when a point-like source is mispointed or when a point-like source is observed against a strong background. The latter is common for sources that lie in the Galactic plane, such as Hii regions. The fringes are then expected at the longer wavelengths where the intrinsic emission of the source decreases and becomes comparable to the Galactic background emission. These fringes affect the continuum as well as the [OI] and [CII] lines at 145 and 158 . They are corrected for in ISAP by dividing out a sinusoidal function with a wavelength dependent amplitude.
The intrinsic wavelength calibration accuracy achieved in the grating mode is better than 1/4 of a resolution element of the grating, i.e. 0.07 for the SW detectors and 0.15 for the LW detectors.
A 10% absolute flux uncertainty is a "reasonable" value for the majority of the LWS observations for all but the SW1 detector. The calibration of the SW1 detector is worse compared to the other detectors, yielding a higher absolute flux uncertainty. In case of very faint sources (<100 Jy at 100 m), very bright sources (>50000 Jy at 100 m), extended sources with a complex morphology, sources not centred in the beam or sources observed against a strong background, the absolute flux uncertainty can be higher. An indicator for the latter is the scatter between adjacent detectors in the defringed and aperture corrected spectrum.
The SWS calibration (except the flux calibration of band 4) and LWS calibration are entirely based on point sources. To correct for source structure within the instrument beam at any given wavelength, the beam profile has to be convolved with the source brightness distribution. This correction can only be applied when the source structure is known at all wavelengths in the spectrum, both for the continuum emitting regions and for the line emitting regions. Such a detailed knowledge of the source structure is not available for any of the Hii regions discussed in this paper.
In addition, for the LWS detectors, a substantial fraction of the flux from an on-axis point source is diffracted out of the aperture. Because the calibration is applied to point sources observed on axis, these losses are automatically corrected for. However, these diffraction losses do not occur in the case of extended sources and so the flux of these sources is overestimated.
The spectra have therefore been derived assuming that the sources are point-like for both SWS and LWS and on axis for LWS. The jumps occuring at the SWS band edges and/or at the junction of the SWS and the LWS spectra and/or at the overlap of two LWS detectors can be caused by this effect. The line fluxes listed in Tables 5-8, also assume that the sources are point-like and on axis. For extended sources, conversion factors (depending on the wavelength) should be applied to convert the fluxes from Jy to MJy/sr. The conversion factors for an infinitely extended and homogeneously distributed source at the key wavelength of each AOT band and LWS detector are given in Salama (2000) and can be applied, if desired, for both the spectra and the presented fluxes. These extended source correction factors are illustrative so caution should be taken when applying them.
The fluxes of the lines in the SWS and LWS spectra are measured by fitting a Gaussian with the interactive line-fitting routine in ISAP. For the lines in the LWS spectrum, the FWHM of the Gaussian fits was kept fixed with a value equal to the spectral resolution element of the data (0.283 for lines in detectors SW1-5, and 0.584 for detectors LW1-5). The FWHM in the fit of the SWS lines was kept as a free parameter as it changes with wavelength. The line fluxes were measured in the independently reduced "up'' and "down'' scans and in the combined spectrum. The latter measurement gives an estimate of the statistical error of the line flux through the error in the fit. The difference in the line fluxes measured in the "up'' and "down'' scans, on the other hand, provides an estimate of the systematic error coming from memory effects and the uncertainty in the choice of the local continuum. The error listed in Tables 5-8 is the square root of the quadratic sum of these two errors (the calibration uncertainties are not included in these listed errors).
A line is defined as being detected if it is present in both the up and the down scan, its peak intensity is two or more times higher than the rms noise of the local continuum and when the line has a FWHM of the spectral resolution element. For any line which does not fulfil all these criteria, an upper limit for the flux is derived. The upper limits are defined as the flux of a feature with a peak flux three times the continuum rms noise and a width equal to the instrumental resolution element.
A few lines occur in the overlap region of two SWS AOT bands. Usually, one of each pair of lines is incompletely covered
or fall outside the "official'' wavelength range of the AOT band, where
the sensitivity of the detector has dropped dramatically. Hence, the line
in the other band is used. The lines in these overlap regions listed in
Tables 5 and 7 are :
in AOT band 1B, Hi (10-5) and Hi (9-5) in 1D, Hi (8-5) in
1E, Hi (14-6) and Br
in 1E and [Arii] 7
|Figure 1: Comparison of the [OIII] 88.3 fine-structure line present in both the SW5 and the LW1 detector. The error bars correspond to the maximum of the fit error and the up-down error (see text). The grey band corresponds to the 10% calibration uncertainty expected for both detectors.|
|Open with DEXTER|
In order to calculate the final error on the line fluxes, a description of the errors that are involved, is given. First, by fitting the observed line with a Gaussian, a statistical error of the line flux is obtained through the error in the fit. Second, an additional error is obtained by the difference in the line fluxes measured in both the - independently reduced - up and down scans. Third, all SWS AOT bands and LWS detectors have an absolute flux uncertainty. This absolute flux uncertainty is discussed in Sect. 3.1.2 for SWS and in Sect. 3.2.2 for LWS. Fourth, all steps in the reduction process involve additional errors. Recently, a full error propagation is done for each detector flux for the applied SWS reduction process (implemented in the test version of IA by fall 2000). Since, this was not included in the pipeline at the time the reduction was done, this error estimate is not available. The total error on the line flux is generally dominated by the absolute flux uncertainty.
For two lines in the same AOT band (for SWS) or in the same detector (for LWS), the uncertainty in the absolute flux will cancel out when taking a ratio of them. Hence, in this case, the final error on the line flux, used to calculate the error on the considered line ratio, will be the square root of the quadratic sum of the statistical error coming from the Gaussian fit and the error derived from the difference between the up and down scan measurements (i.e. the error listed in the Tables 5-8).
When the considered lines of a line ratio are observed in different
SWS AOT bands or LWS detectors, the
total error on each individual line flux must include the absolute
flux uncertainty of the SWS AOT band or LWS detector in which the line
is observed. In this case, the error, ,
on each individual
line is based on the following formalism :
The extend and structure of some sources (see Sect. 6.2) influence the observed line fluxes in two ways. First, the aperture size changes with wavelength (see Sects. 2.2 and 2.3). Second, the morphology, the size of the source, the offset from the source peak position and the pointing error of the satellite are important for the observed flux due to the instrumental beam profiles (see Sects. 3.1.3, 3.2.3 and 3.3). Several sources, flagged in Tables 5 and 7, show offsets between the SWS pointing and the radio source peak position of more than 10 . Other sources (e.g. IRAS 17279) show complex structure. Hence, care should be taken in the interpretation and analysis of the data.
|Figure 2: Comparison of the LWS and KAO line fluxes for [OIII] 51.8 , [NIII] 57.3 and [OIII] 88.3 fine-structure line. Beam size differences are not taken into account.|
|Open with DEXTER|
|Figure 3: Comparison of the SWS and KAO line fluxes for [SIII] 18.7 , [SIII] 33.5 and [NeIII] 36.0 fine-structure line. Beam size differences are not taken into account.|
|Open with DEXTER|
For a number of Hii regions in the present sample, fine-structure lines were measured with the KAO (see Table 3 for references). For the [Oiii] 51.8 and 88.3 and the [NIII] 57.3 , the LWS line fluxes are systematically higher compared to the KAO fluxes (see Fig. 2). However, their ratio is smaller than the aperture ratio suggesting that the sources are more extended than the KAO beam but smaller than the LWS beam at these wavelengths. Analogous, for the [SIII] 18.7 and 33.5 and [NeIII] 36.0 , the fluxes obtained by the KAO are higher compared to SWS suggesting that the sources are more extended than the SWS beam at those wavelengths (see Fig. 3).
|IRAS 01045||Rudolph et al. (1997)|
|IRAS 12073||Simpson et al. (1995)|
|IRAS 17455||Megeath et al. (1990); Simpson et al. (1995)|
|IRAS 17591||Afflerbach et al. (1997)|
|IRAS 18032||Afflerbach et al. (1997)|
|IRAS 18317||Simpson et al. (1995)|
|IRAS 18434||Herter et al. (1981); Megeath et al. (1990);|
|Simpson et al. (1995)|
|IRAS 18479||Afflerbach et al. (1997)|
|IRAS 19110||Herter et al. (1981)|
|IRAS 19598||Megeath et al. (1990); Colgan et al. (199)|
|IRAS 21270||Rudolph et al. (1997)|
|IRAS 21306||Rudolph et al. (1997)|
|IRAS 23030||Herter et al. (1982); Simpson et al. (1995)|
|Figure 4: Comparison of the SWS and IRAS line fluxes for [SIV] 10.5 , [NeII] 12.8 , [NeIII] 15.5 and [SIII] 18.7 fine-structure lines. The IRAS fluxes are taken from Simpson & Rubin (1990). Beam size differences are not taken into account.|
|Open with DEXTER|
|Figure 5: Comparison of the IRAS PSC fluxes and the flux seen by ISO within the IRAS filters centred at 12, 25, 60 and 100 . Beam size differences are not taken into account.|
|Open with DEXTER|
Analogously, the [SIV] 10.5 , [NeII] 12.8 , [NeIII] 15.5 and [SIII] 18.7 fine-structure lines are compared with IRAS observations (see Fig. 4). The observed IRAS fluxes are systematically higher compared to the ISO fluxes due to its larger aperture. However, their ratio is smaller than the aperture ratio suggesting that the sources are more extended than the SWS beam but smaller than the IRAS beam at these wavelengths.
In addition, we compared the flux seen by ISO in the four IRAS bands at 12, 25, 60 and 100 with the corresponding IRAS PSC fluxes. The combined ISO SWS-LWS spectra are therefore used as they are observed. Hence, no correction is applied to the spectra. Figure 5 shows clearly that the 60 and 100 filters give similar results for the IRAS PSC and ISO. For the 12 and 25 band filters, large differences are noted implying that the corresponding emission regions are larger then the ISO-SWS aperture.
Kinematic distances to the program Hii regions were derived from velocities taken from the literature as detailed in Table 4. A flat Galactic rotation curve and the IAU standard galactic constants ( kpc and ) were adopted. The kinematic distance to an Hii region is best determined using velocities of radio hydrogen recombination lines. Where such data are available, they have been used. Where no recombination line data were available, the velocity information derived from molecular studies was used. In case of different molecular observations and/or various velocity components observed for one molecular species, we selected the velocity from a high density tracer with the following order of preference: CS, NH, HCO and CO.
Table 4 gives the selected tracer (Col. 1), its velocity (Col. 2), the derived galactic distance (Col. 3), the distance from the Sun, including the near and far distance for the sources in the inner solar circle, the latter one in brackets (Col. 5) and the list of references for the observed velocity (Col. 6). For sources within the solar circle there is always an ambiguity in determining the distance to the Sun due to the fact that there is a "near'' and a "far'' intersection of the line of sight with the orbit of the source around the Galactic Centre giving the same radial velocity.
|Figure 6: Distribution over the Galactic plane of the Hii regions listed in Table 1 (using the kinematic distances from Table 5). Distances to the Galactic Centre are indicated by concentric circles from 4 to 24 kpc by steps of 4 kpc. The position of the Sun (at 8.5 kpc) is shown by the symbol. : sources beyond the solar circle; : sources in the solar circle using the near distance; : sources in the solar circle using the far distance.|
|Open with DEXTER|
For IRAS 12331, IRAS 18502 and IRAS 19207, the Galactic rotation model
is not giving the measured
indicating that these Hii regions deviate from the model. Here, the most probable
estimate of the solar distances is given, i.e. the distance of the tangent point
providing the closest
to the observed one.
The radio recombination line obtained towards IRAS 17591 has a very
complex profile and the derived distance is very sensitive to the
exact velocity. Hence, we have followed Bronfman et al. (1996) and assigned
a distance of 5.5 kpc derived from CS observations.
|Figure 7: Distribution of the Hii regions as a function of galactocentric radius ( ).|
|Open with DEXTER|
Figure 6 shows the location of the Hii regions from Table 1 projected into the Galactic plane and Fig. 7 shows their distribution as a function of galactocentric distance ( ), using the adopted distances from Table 4. The nebulae are distributed in the Galactic plane from the centre out to 22 kpc. Whereas the Galactic disk is well sampled between 3 and 12 kpc, fewer sources were observed in the inner regions and in the outer Galaxy.
|IR 01045||H42||-69.3||13.8||7.0||1||spherical||3||3||1, 12, 21||4.11|
|IR 01420||CO||-81.9||16.3||9.5||2, 3||spherical||21||2||21||2.86|
|IR 02219||H66||-36.5||11.0||3.3||1||shell||35||0||1, 22, 23||52.9|
|IR 02383||CO||-71.5||15.8||8.6||2, 3||cometary||28||15||21||2.67|
|IR 04025||CO||-60.2||18.5||10.7||2, 3||-||-||-||-||1.27|
|IR 06158||CO||+31.3||18.0||9.7||2, 3, 5||complex||60||11||25||5.29|
|IR 17160||H109(110)||-69.0||3.0||5.7||(11.0)||6||complex||10||19||26, 27||35.6||(133)|
|IR 17279||H109(110)||-33.0||3.4||5.1||(11.8)||6||complex||8||6||26, 27||10.4||(55.9)|
|IR 17455||H70||-19.6||0.2||8.3||(8.7)||8||core-halo||14||0||27, 28, 29||49.2||(54.1)|
|IR 17591||CS||+15.4||5.5||3.0||(13.8)||4||complex||8||0||1, 27, 28||2.04||(43.1)|
|IR 18032||CS||+4.6||7.6||1.0||(15.8)||4, 9||cometary||15||0||27, 29, 30||0.99||(249)|
|IR 18116||H92||+51.9||4.3||4.5||(12.0)||10||cometary||20||0||27, 28||18.6||(132)|
|IR 18162||CO||+11.9||6.6||1.9||(14.7)||11, 12, 13||core-halo||5||2||12||2.43||(146)|
|IR 18434||H76||+95.3||4.6||5.7||(9.0)||14||cometary||7||0||27, 30, 31||70.9||(177)|
|IR 18469||CS||+87.2||4.8||5.3||(9.2)||4, 9||core-halo||13||0||27||11.1||(33.4)|
|IR 18479||H66||+15.3||7.5||1.2||(13.1)||1||complex||7||0||1, 27, 28||1.26||(150)|
|IR 18502||H76||+100.0||5.7||7.1||14||core-halo||6||0||27, 28||19.6|
|IR 19442||CS||+22.5||7.6||2.5||(5.8)||4, 9, 16||spherical||1||0||33||5.30||(28.5)|
|IR 19598||H110||-27.5||9.8||8.5||17||arc-like||4||0||1, 17, 34||201|
|IR 21270||CO||-93.4||14.8||11.3||3, 20||complex||120||25||21, 25||7.60|
|IR 21306||CS||-71.1||12.6||8.3||4||complex||18||36, 37||12.1|
|IR 23133||CS||-56.3||11.7||5.5||4||spherical||15||4||37, 38||18.5|
a Approximate source size in arcsec. b Approximate offset in arcsec between the source peak and the
ISO pointing direction. c The luminosity corresponding with the far distance
is given in brackets.
2 CO components found. No other molecular
Complex source: velocity corresponds to the
averaged value of the components (velocity differences of 2.5 kms-1).
Averaged velocity from the different
Velocity forbidden in the sense of Galactic
rotation. A distance of 0.2 kpc is assumed
(Mehringer et al. 1998; Lis 1991).
No IRAS fluxes available.
REFERENCES: (1) Afflerbach et al. (1996), (2) (1994), (3) Wouterloot & Brand (1989), (4) Bronfman et al. (1996), (5) Fich & Blitz (1984), (6) Caswell & Haynes (1987), (7) Liszt & Spiker (1995), (8) Liszt (1992), (9) Plume et al. (1992), (10) Garay et al. (1994), (11) de Vries et al. (1984), (12) McCutcheon et al. (1991), (13) Shepherd & Churchwell (1996), (14) Wood & Churchwell (1989), (15) Reifenstein et al. (1970), (16) Anglada et al. (1996), (17) Roelfsema et al. (1988), (18) Roelfsema et al. (1989), (19) Walmsley et al. (1981), (20) Blitz et al. (1982), (21) Rudolph et al. (1996), (22) Roelfsema & Goss (1991), (23) Tieftrunk et al. (1997), (24) Carpenter et al. (1990), (25) Fich (1993), (26) Martín-Hernández (2001), in prep., (27) Becker et al. (1994), (28) Garay et al. (1993), (29) Mehringer et al. (1998), (30) Cesaroni et al. (1994), (31) Fey et al. (1995), (32) Mehringer (1994), (33) Barsony (1989), (34) De Pree et al. (1994), (35) Isaacman (1984), (36) Ho et al. (1981), (37) Ho et al. (1986), (38) Kurtz et al. (1999)
Despite the selection criterion that program sources should be point-like, many sources were found to be extended during follow-up studies. In part, this is due to the fact that the selection was often only based on IRAS data which has a limited spatial resolution. Recent radio observations show that the morphology of the selected Hii regions is more complex than originally suspected and that some of the sources show structure at the size scale of the SWS apertures. The fact that some of the nebulae are extended and structured makes the interpretation of the ISO spectral data less straightforward than in the case of a point-like source. The main reason is that in the combined SWS-LWS spectrum the aperture size changes over the entire wavelength range from for SWS band 1A to diameter for LWS detector SW3. As a result the region of the extended nebula seen by ISO will be a function of wavelength. In addition, the morphology, size of the sources, the offset from the source peak position and the pointing error of the satellite are important for the observed flux, as mentioned in Sects. 3.1.3 and 3.2.3. Table 4 gives an overview of these characteristics for each source in order to estimate whether these effects may be important. The source morphology and size derived from radio data are given in Cols. 7 and 8, respectively. The morphology classification is based on work by Wood & Churchwell (1989). Column 10 gives the approximate offset between the centre of the source and the ISO pointing direction. For those sources not mentioned in the table, no radio information is available in the literature.
IRAS 02575: Kurtz et al. (1994) found 2 components: A (3.8 , core-halo) and B (1.1 , unresolved). The ISO observations are centred on source A. The ISO apertures contain both a compact HII region and a YSO. IRAS 12073: weaker components are found in both SWS and LWS beams (Martín-Hernández et al. 2001, in preparation). IRAS 15502: weak, unresolved companions are present in the beam (Martín-Hernández et al. 2001 in preparation). IRAS 16128 & IRAS 17160: the ISO pointing is at the edge of the source, excluding large regions in the smaller apertures (Martín-Hernández et al. 2001 in preparation). IRAS 17279: the ISO pointing is at the edge of the source, excluding large regions in the smaller apertures. Weak, unresolved companions are present in the beam (Martín-Hernández et al. 2001 in preparation). IRAS 17591: the ISO pointing is centred on component C (Garay et al. 1993; Afflerbach et al. 1996). IRAS 18032: ISO is centred on component B (Garay et al. 1993; Cesaroni et al. 1994). IRAS 18162: is a star forming region. Contains beside a compact Hii region also a YSO. IRAS 18434: Fey et al. (1995) found extended emission (22 ) at 1.3 cm. Observations at 21 cm by Kim & Koo (2001) found two strong components embedded in an extended emission (6). IRAS 18469: Kurtz et al. (1994) found two unresolved components. Extended emission of 60 is found by Kurtz et al. (1999). IRAS 18479: Kurtz et al. (1994) found 4 components. Extended emission (20 ) is found by Kurtz et al. (1999). IRAS 19442: extended emission of 45 is found (Barsony 1989; Kurtz et al. 1999). IRAS 19598: the object is not isolated. De Pree et al. (1994) found extended emission (17 ). IRAS 21270: the ISO pointing is in between two sources. The given size correspond to the complete complex. The given offset is the offset of the ISO pointing from the closest component. IRAS 21306: the ISO pointing is in between two sources. The given offset is the offset of the ISO pointing from the closest component.
The FIR luminosities are determined by fitting a blackbody through the IRAS fluxes at 25, 60 and 100 taken from the Point Source Catalogue. The applied IRAS colour corrections are those for a source radiating like a blackbody. The derived luminosities are given in Table 4.
We choose to derive the luminosities in this way for the following reasons. First, one could consider integrating the full ISO SWS-LWS spectrum. However, due to the combination of the different aperture sizes within SWS/LWS resp. and between SWS and LWS on one hand and the extend of our sources on the other hand, different regions are looked at within different aperture sizes. Hence, only for point-like sources, the full ISO SWS-LWS spectrum reflects the SED of the source over the total wavelength range, i.e. from 2.3 to 196 . For extended sources, one can not compare straightforward the obtained spectra in different aperture sizes. When integrating the LWS spectrum only, the short wavelength contribution (<45 ), which is a significant fraction of the total luminosity, is ignored. By fitting a (modified) blackbody to the LWS spectrum, the peak of the blackbody is not well defined causing large uncertainties in the derived luminosities. Finally, the IRAS flux at 12 is ignored since many features beside the dust continuum attribute to this filter.
Depending on which IRAS ratio (i.e. (25)/ (60), (25)/ (100) or (60)/ (100)) is taken, a different temperature is derived for this blackbody and hence different colour corrections. This influences the derived luminosity by less than 4%. In general, the FIR spectra of HII regions are not perfect blackbodies but rather a combination of (modified) blackbodies at different temperatures. Derivation of the luminosity assuming that the source radiates like a single temperature modified blackbody gives a difference of less then 7% with those derived by assuming a blackbody. The error on the derived luminosities is dominated by the absolute flux calibration of IRAS, i.e. 30%.
The IRAS fluxes (60) for IRAS 21190 and (100) for IRAS 11143 and IRAS 18116 are upper limits. Hence, the derived luminosities for those sources are also upper limits.
|Figure 8: The combined SWS-LWS spectra. The fine-structure lines, the hydrogen recombination lines and the main PAH features are identified. The wavelengths at which an aperture jump occurs within the SWS wavelength region and the wavelength between the two spectrometers SWS and LWS, are indicated by flags. The diameter of the LWS beam smoothly changes, ranging from 67 -85 , indicated by the dotted line.|
|Open with DEXTER|
Detailed discussions of the atomic fine-structure from the ionised gas and the photo-dissociation regions (PDRs) are left to later publications (PaperII; Morisset et al. 2001; Damour et al. in prep.).
First studies on the dust properties in compact Hii regions have been given in a series of papers (Roelfsema et al. 1996; Cox & Roelfsema 1999; Peeters et al. 1999; Van Kerckhoven et al. 2000; Hony et al. 2001) and a complete study will be presented in Peeters et al. (in prep.). The study of the dust continuum will be presented by Jones et al. (in prep.). The ice content in compact Hii regions is discussed in Boogert (1999), Boogert et al. (in prep.). Hereafter we briefly describe the content of the infrared spectra of compact Hii regions and outline the potential use of the present catalogue in future studies.
In the spectra presented in Fig. 8, the main atoms (O, N, Ne, Ar, S, C and Si) are detected via their fine-structure lines. The lines of [Cii ], [Oi] and [Siii ], which originate in the PDRs associated with the Hii regions, are seen in almost every source. They provide an important diagnostic of the physical conditions of the dense warm gas located just outside the ionised nebula. A detailed analysis will be given in Damour et al. (in prep.) where the present data will be compared to PDR model predictions. The strongest lines in the combined SWS-LWS spectra are emitted by atoms in the ionised nebula which, in rough order of importance, are the fine-structure lines of O++, S++ and S+++, Ne+ and Ne++, Ar+ and Ar++ and N++and N+ (see Fig. 8 and Tables 5 and 6). Except for oxygen, the species are present in different stages of ionisation and the range of ionisation potential sampled by these lines extends to 41 eV. This information allows us to describe the ionisation state of each Hii region and to derive the properties of the exciting star(s). The Hii regions studied in this paper form a prime sample to examine the variations of the nebular properties with galactocentric distance and to investigate the distribution of the relative and absolute element abundances in the Milky Way. The study of the ionisation state and the element abundance based on the present sample of compact Hii regions is the major aim of the UCHii ISO program and the subject of an accompanying paper (Martín-Hernández et al. 2002, i.e. PaperII).
In addition to the fine-structure lines, HI recombination lines are present in the spectra (see Fig. 8). These HI recombination lines are used to derive the extinction law at these wavelengths by comparing the observed line strengths to the predictions of recombination theory (see PaperII).
Some sources exhibit unusual lines. IRAS 12073 has an emission line at 4.3 ([KIV]?), IRAS 12331 at 21.8 and 27.1 (HI or [ArIII]; [SV]) and K3-50 A at 4.295, 4.377 and 4.617 (?, HI, [KIII]?). The SWS AOT06 of IRAS 18434 shows also the 4.617 [KIII] line.
In a few sources, molecular lines of are detected. In one case, IRAS 21190+5140 (M1-78), six emission lines of molecular hydrogen are present in the SWS spectrum, i.e. 1-0 O(3) to O(5) and 0-0 S(3) to S(5). IRAS 12063 shows the 0-0 S(5) line and K3-50 A exhibits the 1-0 O(3) and 0-0 S(1) lines.
The molecular absorption line of OH at 119.2
is present in the following
sources: IRAS 17221, IRAS 17279, IRAS 17455, IRAS 17591, IRAS 18032,
IRAS 18317, IRAS 18434, IRAS 18469, IRAS 18479, IRAS 18502 and Sgr C. The
latter sources also shows an OH absorption line at 79.3
W3 A, IRAS 19442 and
DR 21 exhibit CO emission lines (162.9, 173.6 and 185.9
emission line at 179.7
is clearly present in Sgr C.
TDT = 64600609; TDT = 78800709; TDT = 17701258; TDT = 56101082; see Sect. 3.1.4.
a The listed errors are the square root of the quadratic sum of the difference between the up and down scan measurements and the statistical error of the line fitting (see Sect. 4). Hence, no calibration uncertainties are included! To obtain the final error, see Sect. 4.2. b At these wavelengths the calibration uncertainty is 10%. c Note that these lines are seen through different apertures.
TDT = 17701257; TDT = 56101081; see Sect. 3.1.4.
The infrared spectra of compact Hii regions are dominated by a strong continuum due to the thermal emission of dust. As mentioned in Sect. 6.3, it is not straightforward to compare the obtained spectrum of an extended source over different aperture sizes. So, caution should be taken when studying the SED of a specific source when this source is not a point source. For most of the sources, this dust continuum seems to peak in the range 40-60 when plotted in versus W/cm/ corresponding to dust temperatures of 60-70 K. In some cases (IRAS 11143, IRAS 12073, IRAS 21190 and perhaps IRAS 17455), the spectrum peaks at shorter wavelengths 30 , indicating that warmer dust is present in these nebulae. The Hii regions in this sample appear to be substantially hotter than UC Hii regions, with dust temperatures of 50 K. Probably, the nebulae in this sample are somewhat more evolved and emerging from their dust envelopes. In the case of IRAS 02575, a strong continuum is present shortwards of 5 , indicating temperatures in excess of 800 K. This unusual temperature for an Hii region can be explained by the fact that, for this source, the SWS aperture encloses both the compact Hii region and an Herbig AeBe star (see Sect. 6.2). Finally, many Hii regions show continuum emission at 12 , indicative of a high colour temperature. Yet, it cannot be due to small molecules (50 C-atoms) since those would give rise only to a weak quasi-continuum in this wavelength region (Allamandola et al. 1989). This continuum could be due to a small fraction of dust in the Hii region heated to high temperatures by resonantly scattered Lyman-alpha radiation. The complete wavelength coverage from 2.3 to 196 provides in principle important information on the dust spectrum at mid-infrared wavelengths dominated by the PAH emission bands and the far-infrared continuum which is dominated by the bigger dust grains. Improvement in the SWS calibration will be crucial to further study this part of the spectrum.
The IR emission features at 3.3, 6.2, "7.7'', 8.6 and 11.2 - commonly attributed to PAHs - are present in most of the spectra of the Hii regions together with many weaker features e.g. the 12.7 band (see Table 9). The shape and strength of these dust bands exhibit considerable variations from source to source. At the spectral resolution of the present data, the bands are found to form complex emission patterns with a variety of sub-peaks and spectral detail. For instance, the "7.7'' band consists of two bands peaking at 7.6 and 7.8 . The variations in the relative intensities of the bands are thought to be primarily due to variations in the physical conditions of the emitting region (radiation field, density, etc.) that determines the physical (ionisation) and/or chemical processing of the interstellar PAH family (e.g. Roelfsema et al. 1996; Cox & Roelfsema 1999; Peeters et al. 1999). A complete analysis of the near- and mid-infrared dust bands in the spectra of compact Hii regions, together with a study of the variation of their shapes and relative intensities as a function of the ionisation state (as measured by the atomic lines), will be presented in Peeters et al. (in prep.). In addition to these emission bands, a large number of the Hii regions show in their infrared spectra the broad absorption band at 9.7 due to the stretching mode of amorphous silicate. The corresponding bending mode at 18 might be present in the spectra of some sources but is not always clearly visible. A summary of the dust content in the sample of Hii regions of this catalogue is given in Table 9, where the presence of the dust emission bands and of the silicate absorption bands is indicated for each source. Sources which are not included in Table 9 do not show any emission/absorption features. Given the low continuum of some sources, it is not possible to assess the presence of absorption.
|Transition||(m)||1 a||Line fluxesb|
For three of the sources presented in this
paper, broad absorption bands of amorphous HO ice are clearly
detected at 3 and 6
(the stretching and bending modes, respectively).
These bands might be present in some other sources but the low
flux density levels (<5 Jy) in the 3
results in noisy spectra from which firm conclusions cannot be drawn.
In addition, the strong 3.3 and 6.2
dust emission bands
are located in the red wing of both HO bands resp. making
the situation difficult for sources with moderate amounts of extinction.
|Source||H2O Ices||CO2 Ices||Ice?|
At longer wavelengths, the continuum becomes stronger making the situation more favourable to detect ice bands in absorption. The CO bending mode at 15 is detected in 6 sources (listed in Table 10). Small variations in the substructure of the absorption band show that the ices have only been partially thermally processed ( 45 < T < 90 K) by the central heating source, in contrast to some highly processed regions around hot cores such as Sh 140 (Gerakines et al. 1999). Furthermore, extensive laboratory studies of the band profile show that interstellar CO must be mixed with H2O and CH3OH ices (see Boogert 1999; Boogert et al. in prep.). The intrinsically stronger CO stretching mode (4.27 ) is present in the spectra of 3 sources (Table 10), but the data are of much lower quality due to the lower level of the continuum at these wavelength. Furthermore, it is located in AOT Band 2A, which tends to be very noisy. It is thus not excluded that the CO stretching mode could be present in other compact Hii regions than those listed in Table 10. Finally, K3-50A is the only source in the present sample that shows an absorption band around 6.85 (Table 10). The analysis of this band is difficult since it is located in a complex spectral region which includes the 6.0 HO ice absorption band, the 6.2 PAH emission feature, the 6.8 absorption feature and the emission plateau between the 6.2 and "7.7'' PAH bands. The derived profile is similar in appearance to absorption features seen toward dense molecular clouds exhibiting signs of warmer ice (Keane et al. 2001).
The authors wish to thank the anonymous referee for usefull comments. EP thanks especially J. Cami, R. Shipman, D. Kester, F. Lahuis, S. Hony, A. Boogert, L. Decin and B. Vandenbussche for the support concerning the SWS data reduction. NLMH acknowledges the support of S. Sidher for the LWS data reduction. IA3 is a joint development of the SWS consortium. Contributing institutes are SRON, MPE, KUL and the ESA Astrophysics Division. The SWS work was supported by the Dutch ISO Data Analysis Centre (DIDAC) at the Space Research Organisation Netherlands (SRON) in Groningen, The Netherlands. The ISO Spectral Analysis Package (ISAP) is a joint development by the LWS and SWS Instrument Teams and Data Centers. Contributing institutes are CESR, IAS, IPAC, MPE, RAL and SRON. LIA is a joint development of the LWS consortium. Contributing institutes are CESR, DRAL, IPAC and the ESA Astrophysics Division. EP acknowledges the support from an NWO program subsidy (grant number 783-70-000). NLMH acknowledges the support from an Ubbo Emmius grant for graduate students at the Rijksuniversiteit Groningen.