ASTRONOMY AND ASTROPHYSICS. High-resolution imaging of ultracompact H II regions. II. G revisited

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1 Astron. Astrophys. 36, 3 59 (1999) High-resolution imaging of ultracompact H II regions II. G revisited ASTRONOMY AND ASTROPHYSICS M. Feldt 1, B. Stecklum, Th. Henning 1, R. Launhardt 3, and T.L. Hayward 1 Astrophysikalisches Institut und Universitäts-Sternwarte (AIU), Schillergässchen 3, D-775 Jena, Germany Thüringer Landessternwarte Tautenburg, Sternwarte 5, D-7778 Tautenburg, Germany 3 Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, D-5311 Bonn, Germany Center for Radiophysics and Space Research, Cornell University, Ithaca, NY 1853, USA Received 3 December 1998 / Accepted 3 February 1999 Abstract. We present the results of an extensive imaging campaign on the ultracompact Hii region G at near- and mid-infrared wavelengths. High-resolution data were taken in the H, K, and L bands using ESO s adaptive optics system ADONIS. They are complemented by conventional narrowband images in Brγ and H (1 )S1 as well as by mid-infrared broad- and narrow-band images. We also mapped the 1.3 mm continuum emission from the source using the SEST. We use our data to consistently explain the morphological appearance of G at all observed wavelengths as well as its spectral energy distribution. The complete model of the source consists of a spherical shell of dust with an inner, dustfree cavity of cm radius surrounding a star of spectral type O6 ZAMS. Two outflows escape this shell in opposite directions. Half of the whole configuration is evidently obscured by a very massive cloud of cold dust. Comparisons with earlier models and other ultracompact Hii regions are drawn to put G in the context of massive star formation. Key words: stars: formation stars: early-type ISM: H ii regions methods: observational 1. Introduction This is the second paper in our series on high-resolution infrared observations of ultracompact Hii regions (UCHiis). The first paper (Feldt et al. 1998, hereafter Paper I) concentrated on the object G This time G (hereafter denoted as G5.89) was chosen for two reasons. Firstly, it met our selection criterion which demands that a nearby bright star exists to serve as a wavefront sensor for the adaptive optics system. Secondly, it is one of the best-studied and yet still disputed UCHiis. G5.89 was classified as a shell-like UCHii by Wood & Churchwell (1989, hereafter WC89). Its kinematic distance of Send offprint requests to: M. Feldt (mfeldt@mpia-hd.mpg.de) Based on observations collected at the European Southern Observatory, La Silla, Chile.6 kpc from the sun was determined by Downes et al. (198) and little disputed since. Although recent estimates by Acord et al. (1998) put this distance close to the possible upper limit, we will use it throughout the paper. Apart from the VLA observations by WC89 which show the shell-like morphology, Gomez et al. (1991) observed the source with the same instrument in D configuration at 1.3 cm. To them, the source appears essentially unresolved, with a weak extension towards the southeast. Mid-infrared (MIR) observations by Ball et al. (199) showed that the MIR emission closely resembles the radio structure, apart from the southern half of the shell, where no MIR emission was detected. They regarded this phenomenon to be due to an off-centre location of the ionizing star. Models of the source by Churchwell et al. (199, hereafter CWW) and recently by Faison et al. (1998) showed that the spectral energy distribution (SED) of the source can be fitted by emission from a spherically symmetric cocoon of dust surrounding a central zero-age main-sequence (ZAMS) star of spectral type O6. To fit near-infrared (NIR) part of the spectrum, a dust-free cavity inside the cocoon was required. The radial intensity profiles of their models already predicted the presence of a ring of emission at all wavelengths. A similar model was proposed by Harvey et al. (199). However, the latter authors regard their own spherical model as highly unlikely due to the obvious non-spherical morphology of the source. G5.89 is also known to drive one of the most energetic outflows in the Galaxy. A lot of discussion has been going on about the orientation of this outflow and whether there are multiple outflows and driving sources (Harvey & Forveille 1988; Zijlstra et al. 199; Cesaroni et al. 1991; Acord et al. 1997). In this paper, we are going to present the result of a multiwavelength imaging campaign of the source. High-resolution observations at 1.6 µm,. µm, and 3.5 µm taken with ESO s adaptive optics (AO) system ADONIS are complemented by MIR data with arcsecond resolution taken at 1.6 µm, 11.7 µm, 1.8 µm, and 1 µm. These observations show that the shelllike structure of the source is preserved throughout this whole wavelength range. Additionally, large-scale narrow-band images were obtained in the H (1 )S1 and Brγ lines. These

2 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II Table 1. Summary of observations Date λ Tel./Instr. FOV a PSF Ref. Star b Limiting FWHM Mag. c 1995 Aug H/K ESO 3.6 (ADONIS/SHARP) 13. 5/. d Y338 (5.68/5.) 1.1/ Jul.17/.16 µm (Brγ/H (1 )S1) ESO. m/iracb HD (8.71) May L ESO 3.6 (ADONIS/COMIC) May 1.6 µm/ 1.8 µm ESO 3.6 m (TIMMI) ηsgr( 1.65/ 1.7).9 / Aug 11.7 µm Hale/SpectroCam αlyr (.) Oct Q ESO. (MANIAC) ηsgr (-1.7) Mar 1.3 mm ESO SEST (Bolo) 3 Uranus 1 mjy a Field of view for single frames b Brightness given in magnitudes for the observed wavelength c Derived from background noise for point sources with given PSF (1σ detection) d Seeing was 1. during the observations, the resolution is improved by the AO correction data were used to determine the extinction towards the source as well as to get a large-scale overview of the region. They indicate the presence of a molecular cloud in front of G5.89 which is obscuring about half the source from view even at infrared wavelengths. The presence of this cloud is confirmed by our continuum map taken at 1.3 mm using the SEST. We have collected infrared data of arcsecond resolution over such a wide range in wavelengths on this source for the first time. The NIR images even are of a resolution of. (full-width halfmaximum (FWHM) of the point-spread function (PSF) in K ) thanks to the adaptive optics. We use these data to build a consistent model of the source. We will argue that the spherical dust shell model is essentially correct and present a detailed fit not only to the SED, but also to the intensity distribution and various properties of the source. This model is consistent with all our observations and explains G5.89 as an almost classical Strömgren sphere. We interpret the visible break in spherical symmetry as channel openings, through which an outflow is escaping in north-south direction. This outflow is detected in our data as H (1 )S1 emission features. From data on a larger scale, we can show not only the presence of a large cloud of cold dust at the rim of which G5.89 is located, but also identify several stellar sources which belong to the complex. Thus, like in Paper I, we are probably looking at a young cluster in formation. Due to the smaller distance compared to the 6.3 kpc of G5.5-.6, we are now able to look at the single source G5.89 inside this cluster and describe it in detail. This is another hint that classical UCHiis with spherical configurations exist, but that they are usually part of young stellar clusters which will eventually form a larger ionized region.. Observations, data reduction, and calibration.1. Adaptive optics NIR imaging Near-infrared imaging was performed in August 1995 using ESO s AO system ADONIS (Beuzit et al. 199) on the 3.6 m telescope at La Silla/Chile. High-resolution images were obtained in H and K. In each band, a mosaic of three frames was obtained resulting in a total integration time of to 6 s, depending on the location in the image. During the observations the seeing was 1. The high-order AO correction improved the FWHM of the PSF to. ink and. 5inH. All frames were subject to standard bad-pixel removal, flat fielding, and dark-frame subtraction processes before being combined in the resulting images. For calibration purposes, images were taken of the UKIRT standard Y NIR narrow-band imaging Narrow-band images were taken at ESO s. m telescope on La Silla/Chile in June IRACb served as camera. Four filters were used with central wavelengths.111 µm (BP),.16 µm (BP5, used as H (1 )S1 filter),.15 µm (BP7), and.17 µm (BP8, used as Brγ filter). The central wavelengths of the filters were computed from the filter transmission curves provided by ESO. They are not the same as stated in the IRACb manual. The use of lens C resulted in a pixel scale of. 57 per pixel. A mosaic of five frames was taken with the object once in the centre of the detector and once in the centre of each quadrant. The total integration time on source added up to minutes. The same procedure was done for the standard star HD with 1 minutes of total integration time. The rather complex procedure of calibrating the narrowband images and doing the continuum subtraction is described in Appendix A..3. MIR imaging.3.1. L-band imaging The L band image was obtained in May 1996 using the COMIC camera in combination with ADONIS. 3 frames of 3 s integration time each were combined in the resulting image. As the original goal was to do polarimetry, these frames were obtained at four positions of the polariser. Because the signal-to-noise ratio proved to be too low for polarimetry, all frames were combined in the resulting image. The mean Strehl ratio given by the ADONIS software was.75, the PSF FWHM in the image

3 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II 5 is. 5 while the seeing monitor reported a seeing of 1.. No photometric calibration was done for this image µm observations The 11.7 µm image was obtained using SpectroCam-1 (Hayward et al. 1993) at the -inch Hale Telescope of the Palomar Observatory 1. The effective wavelength of the filter was 11.7 µm with FWHM = 1 µm. A 5 frame mosaic was combined into the image. The average on-source integration time at each pixel is s. The star αlyr served as a standard for flux calibration. This image will be referred to as the 1 µm image in the discussion Q-band observations The Q-band data were taken in October 1996 at ESO s. m telescope on La Silla/Chile. The mid-infrared camera MA- NIAC (Böker et al. 1997) was used to obtain images using its 18 3 µm (Q) filter. The total integration time on the object summed up to 1 s including combination of the two chopping beams. Photometric calibration was obtained by observing the standard η Sgr..3.. MIR narrow-band observations Narrow-band observations in the mid-infrared were carried out at ESO s 3.6 m telescope with the TIMMI camera (Käufl et al. 199). The filters [Siv](λ c =1.6 µm, FWHM =. µm) and [Neii] (λ c =1.8 µm, FWHM =. µm) were used. Chopping and nodding was done to obtain 6 images of 37.7 ms integration time in each of the two filters. For photometric calibration, the standard star η Sgr was observed. To derive its brightness at the desired wavelengths, we used its IRAS LRS spectrum mm continuum measurements The continuum radiation of G was mapped in March 1996 using the 15 m SEST telescope at La Silla, Chile, together with the 3 He-cooled single-channel bolometer system (Kreysa 199). The equivalent bandwidth of the bolometer is 5 GHz centred on a frequency of ν = 36 GHz (λ =1.7 mm). The effective beam size at this wavelength is θ b =3. The source was mapped four times with the double beam technique described first by Emerson et al. (1979). To generate the dual beams, a focal plane chopper with a chopping frequency of 6 Hz was used. Chopping was done in azimuth. The chopper throw was 67. The map rows were generated by moving the telescope continuously along the direction of the beam separation (i.e. in azimuth) with a scanning velocity of 8 /sec and an elevation separation between adjacent scans of 8. 1 Observations at the Palomar Observatory were made as part of a continuing collaborative agreement between the California Institute of Technology, Cornell University, and the Jet Propulsion Laboratory. Calibration maps of the planet Uranus (adopted brightness temperature 96 K; Griffin & Orton 1993) were obtained with the same technique and parameters as used for G The atmospheric transmission was measured by sky dips. Telescope pointing was found to be repeatable within ±5. Data reduction was performed with the SEST standard software and with the software package MOPSI (written by R. Zylka) which use the NOD and GAG libraries. 3. Results 3.1. NIR images Figs. 1 and show the results of our adaptive optics near-infrared imaging. The images presented in the figures were subject to filtering with the multi-scale maximum entropy method described by Pantin & Starck (1995). To enhance the details in the images as they are shown in the figures while at the same time preserving extended structures, we present an additive combination of the image with a deconvolved version. The deconvolution was achieved by applying 1 iterations of a maximum-likelihood algorithm as described in Lucy (197). The combination was done by adding.3 times the deconvolved image and.7 times the original one. The reference position given in the figures denotes the centre of the approximately spherical shell visible in the radio contours. This position will be used extensively in forthcoming sections. Superimposed on both images are the VLA maps obtained by WC89. In the lower left corner of the images, a bright star can be seen which served as wavefront reference for the AO system. A band of extended emission and unresolved sources stretches towards the north-east at a position angle of approximately 5. Several point-like sources are embedded in the extended emission. Both extended and unresolved emission are stronger in K than in H (note the different cut levels!). Apart from the bright star, the most prominent structure in K follows the northern half of the radio shell. This is what we will later discuss as the visible half sphere of the dust shell surrounding a hot O6 ZAMS star. In H, only weak remains of these structures are visible, except two supposedly stellar sources at positions (+. 3,-1. 5) and (+. 5,-. 5) in the southern half of the radio shell. Later, we will argue that these sources are probably foreground objects. 3.. MIR images The results of our broad band imaging in the mid-infrared are presented in Figs. 3 and. The Q-band image was subject to noise removal by the maximum entropy filter described above. Both images show the characteristic arc-like structure already visible in K : A banana-shaped bow opening to the south and following the contours of the northern half of the radio shell. The contours of the L band image confirm that this general structure is present also at that wavelength. However, especially from Fig. it becomes clear that extended emission exists beyond the radio shell. The extraordinary fact that the general shape of the object remains constant from the near- to the mid-infrared indicates that the optical depth does not change significantly

4 6 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II G K + cm VLA Contours G N + L Contours REFERENCE POS. R.A. 17 h 57 m 6 ṣ 76 DEC - o 3 56."7 (195) Fig. 1. K image of G5.89. The logarithmic gray scale ranges from.39 mjy/ to 31 mjy/. To enhance the visibility of details, a combination of.3 times the deconvolved and.7 times the original image is shown. The image was subject to a maximum entropy filtering algorithm (see text). The contours are from the cm VLA map by WC89. The levels are,, 6, 8, and 1 times the 1σ level of.3 Jy/beam REFERENCE POS. R.A. 17 h 57 m 6 ṣ 76 DEC - o 3 56."7 (195) Fig µm image of G5.89. The linear gray scale ranges from 15 mjy/ to 15 Jy/. Contours are from our L band image. Arbitrary levels are 3, 6, 9, 1, and 18 times the 1σ-level in the image. G H + 6cm VLA Contours G Q + cm VLA Contours REFERENCE POS. R.A. 17 h 57 m 6 ṣ 76 DEC - o 3 56."7 (195) Fig.. H image of G5.89. The logarithmic gray scale ranges from. mjy/ to 31 mjy/. The same filtering and deconvolution procedure as for Fig. 1 was applied. The contours are from the 6 cm VLA map by WC89, levels are,, 6, and 8 times the 1σ level of 5 mjy/beam REFERENCE POS. R.A. 17 h 57 m 6 ṣ 76 DEC - o 3 56."7 (195) Fig.. Q-band image of G5.89. The linear gray scale ranges from 6.9 Jy/ to Jy/. The contours are the same as in Fig. 1.

5 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II 7 over these wavelengths. On the other hand, it is noteworthy that no infrared radiation at all is detected in the southern part of the radio shell. This is extensively discussed in Sect.. We note that the maximum of the emission shifts to the west along the arc by. 6inN (11.7 µm) and by 1. inq compared to K mm continuum map Fig. 5a shows the result of our 1.3 mm continuum mapping. The 1σ rms level measured in source-free regions in the map is 115 mjy/beam. Note that the continuum emission at 1.3 mm is the sum of thermal dust emission and free-free emission. The contour map shows that the source appears slightly extended in the 3 SEST beam. Deconvolution yields an intrinsic FWHM of 16 ±1 1 ±. The total flux of the source is Jy with 17 Jy in the compact core and 5 Jy in the envelope. The extended feature seen in the map might partially be the result of the beam side lobes. However, the signal denoted by the lowest contour line contains far more power than the error beam. Thus the extension is indeed real. This becomes especially obvious from part B of the figure, where profiles of the source and, for comparison, of Uranus are shown. To access the distribution of the cold dust, we subtracted the free-free contribution from the 1.3 mm map. This was achieved by multiplying the cm map with a factor of (/.13).1 (assuming optically thin free-free emission) and convolving the resulting map with the SEST beam before subtracting it from the 1.3 mm map. The result of this operation is shown as contours in Fig. 6. Here it is obvious that the orientation and position of the thermal dust emission is consistent with that of a presumed foreground cloud (see Sect..1.1) extincting the southern half of G5.89. At the eastern edge of the mapped region (the scanning direction was north south), a second source is visible with a peak flux of 1.5 Jy/beam (approximate peak position R.A.= 17 h 57 m 35 ṣ 8 and Dec.= - 9 ). However, this source is not completely covered by our map. This source is obviously the Hii region G5.9-.3, visible in the radio continuum maps of Zijlstra et al. (199) and denoted as source B in their paper. 3.. Narrow-band data Images The results of our NIR narrow-band imaging are presented in Fig. 6. The image is colour-coded, red representing flux measured in the H (1 )S1 line, green the corresponding continuum filter flux and blue the flux measured in the Brγ filter. This image gives a large-scale overview of the region. In the lower right fewer stars are visible than in the upper left. This suggests that large clouds are obscuring our view towards that region. Close to the H source labelled A, brownish, arc-shaped structures are visible which might be regarded as rims of clouds reflecting light from nearby stars (similar to the fingers in M16; see Pound 1998, Hester et al. 1996). The impression is that G5.89 (situated at the reference position) is sitting exactly at the rim of such a cloud with only its northern half visible. Table. Line fluxes Line Location Ap. ø Flux Brγ Radio Shell (±.) 1 16 Wm H (1 )S1 A. 5 1.(±.1) 1 16 Wm H (1 )S1 B. 5.53(±.5) 1 16 Wm H (1 )S1 C. 5 3.(±.9) 1 16 W a (m ) [Neii] 1.8 µm Radio Shell ± 1 Jy b 1.6 µm Radio Shell 5..6 ± 1.8 Jy b a Measured as surface flux only b No continuum subtracted, total flux density For the mid-infrared data, no subtraction of continuum flux from the narrow-band images was attempted. The main contribution to the flux density in the image taken at 1.8 µm comes from the [Neii] line. Thus, bright regions in that image should be the highly ionized regions in G5.89. From Fig. 7, we learn that indeed these regions correlate very well with the highest radio contours. As no [Neii] emission is seen in the southern half of the radio shell, we infer that the optical depth there is still high at 1.8 µm. The image taken at 1.6 µm could contain flux from the [Siv] line. From Faison et al. (1998) we learn that this line is virtually not present in G5.89. Thus, we use the resulting image to determine the optical depth of the silicate feature by comparing its flux to that in our 11.7 µm image. The result of this procedure is presented in Fig. 8. The natural logarithm of the ratio of the two fluxes is shown as gray scale, superimposed are contours denoting the cm radio contours. We note that the logarithm of the flux ratio can only serve as a measure for the optical depth (The flux is not even measured at 9.7 µm after all!), but does not yield its absolute value. From Fig. 8, we learn that the silicate optical depth is highest at the locations of the largest radio flux. This behaviour is somewhat peculiar and not fully understood. We will come back to this feature in Sect...1 when discussing the properties of our model of G Line fluxes Line flux densities were measured using the calibration procedure described in Appendix A and then performing aperture photometry on the calibrated images. The total line fluxes were then obtained via the filter widths and are presented in Table. We note that our measured Brγ flux is about 3.8 times higher than that measured by Moorwood & Salinari (1983). Also, the flux density in the [Neii] filter measured by us differs from that measured by Faison et al. (1998) by around 8 Jy. The latter authors measured a peak flux density of 5 Jy. Of course different spectral resolutions and responses might account for this difference, because of the averaging over adjacent continuum flux. On the other hand, our measured flux of.6 Jy at 1.6 µm seems to agree with their data as seen in their Fig. 1. The neon flux is particularly interesting because it allows to determine the ratio of Ne + /H +. Unfortunately, we have no chance of subtracting the underlying continuum flux density

6 8 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II Fig. 5. a 1.3 mm map of G5.89, taken with the SEST. The contour levels are.35 (3 σ level),.7, 1.5,.5 to 1 by 1.5 Jy/beam. The thick line marks the 5% contour (5.5 Jy/beam). The beam size is indicated in the lower left. b Flux profiles of the source. The solid-line profile was extracted along the solid line in the contour map, the dashed line correspondingly. The dotted line shows a strip through a map of Uranus for comparison. from the total flux density to get the line flux. However, we can estimate the Ne + abundance by judging the line contribution to be roughly % from Faison et al. (1998). If we convert the measured total flux density of 37 Jy to a flux of Wcm using the filter width of. µm and the % mentioned above, we can use Eq. (3) of Watarai et al. (1998) to determine the ratio of the number abundances. Inserting the appropriate values for temperature (1 K) and electron number density ( cm 3, see Sect. 3.5) as well as the size of the source ( sr) and the emission measure (6 1 7 pc cm 6, see Sect. 3.5), we derive Ne + /H + = This value seems reasonable compared to the found by Watarai et al. (1998) for G and the found in the Orion Nebula (Rubin et al., 1991). The value for Orion gives the ratio of the neutral atom number abundances. Since we cannot tell anything about the Ne ++ abundance, except that it should be high because of the early type central star of G5.89, the difference might be easy to understand in this case The extinction towards G5.89 To derive the NIR-extinction towards an ionized source, one can compare the radio flux to the flux in recombination lines such as Brγ. In this work we used essentially the same procedure as in Paper I which in turn followed that described by Watson et al. (1998). One main difference to Paper I should be noted: This time, we have two radio images, taken at and 6 cm. However, the spectral energy distribution of G5.89 (see, e.g. WC89) shows that the emission at 6 cm is already optically thick. Thus, it is possible to use the beam temperature of the 6 cm map as electron temperature T e directly. The temperatures then range from.8 1 Kto1.6 1 K with a mean value of K. The temperature distribution is of course the same as the that of the 6 cm emission. With this knowledge, the emission measure and the expected Brγ flux were computed from the optically thin cm emission. Missing flux in the VLA maps should not be a problem because large-scale maps by Gomez et al. (1991) taken in VLA-D configuration show that no large-scale emission exists, except in their lowest contours. The computed emission measure peaks at pc cm 6 with an average value of pc cm 6. This deviates from the results of WC89 and those of Zijlstra et al. (199), because they used uniform electron temperatures of 1 K and 8 K, respectively. The mean electron densities implied by these values and the mean path length of.5 pc is cm 3. The resulting Brγ-extinction map is shown in Fig. 9. The Brγ-extinction rises constantly from northeast to southwest. We will later argue (Sect..1.1) that this is also due to the foreground cloud, not due to internal extinction in the dust shell The hot dust in G5.89 From our 11.7 µm and Q-band images, we can derive the dust distribution inside G5.89 as well as the temperature of that distribution starting from a simple model. Assuming the emission to be optically thin at both wavelengths, the ratio of the two flux densities yields the temperature, while the total flux density in combination with the derived temperature yields the mass of dust. The flux density emitted from a dust mass M dust at frequency ν is given by: F ν = M dusthν 3 κ ν d c 1 e hν kt 1 In this expression, the dust mass absorption coefficient κ ν is taken from Ossenkopf & Henning (199) for the appropriate wavelengths assuming a gas density of 1 5 cm 3 and a size distribution after MRN (Mathis et al., 1977) without ice mantles. The quantity T is the dust temperature, h the Planck constant, k the Boltzmann constant, and d the distance towards the source. After adjusting the PSFs of both images by convolving the 11.7 µm image with an appropriate Gaussian, we were able to derive both the temperature and the mass of the dust. This procedure was applied in image areas where both signals were above their corresponding 3σ levels. The result of this procedure can be seen in Fig. 1. The grey scale represents the mass distribution of the dust. Superimposed (1)

7 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II 9 Fig. 6. Colour-coded image of G5.89 taken in three narrow-band filters. Red represents light emitted in the H (1 )S1 line, green is narrow-band continuum emission and blue Brγ. The contours denote the 5% and the 9% level of the free-free subtracted 1.3 mm continuum emission. The marks A, B, and C denote locations where H flux was measured (see text). The arrow points to the southeastern beginning of the brownish cloud rim stretching across G5.89. G5.89 itself is located at the reference position. are contours denoting the temperature distribution. The lowest contour line is at 1 K, the spacing is 5 K. We note that our result concerning the temperature qualitatively agrees with that of Ball et al. (199). However, we are measuring a much lower absolute temperature with a peak value of only 13 K (Ball et al. derived a peak value of roughly 5 K) and an average temperature of 11 K. Several explanations apply. Ball et al. (199) discuss in depth the requirements for deriving a physical temperature from an MIR flux ratio and the drawbacks suffered by deriving a simple colour temperature. However, we believe that our estimates are considerably closer to the actual physical temperature because of the following reasons: 1. We included the dust opacities from Ossenkopf & Henning (199) for the two wavelengths.. A comparison with the computed model in Sect...1 shows that this temperature range is indeed reasonable. The corrections introduced by the finite optical depth in that model are only of the order of %. 3. Ball et al. (199) derived their colour temperature from the flux density distributions at 8.5 µm and 1.5 µm. Their measured 1.5 µm flux density of G5.89 is only 7.1 Jy,

8 5 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II G5.89 NeII G5.89 Optical Depth REFERENCE POS. R.A. 17 h 57 m 6 ṣ 76 DEC - o 3 56."7 (195) Fig. 7. Image of G5.89 taken in the [Neii] filter at 1.8 µm. The linear gray scale ranges from to 36 Jy/. Contour lines are the same as in Fig REFERENCE POS. R.A. 17 h 57 m 6 ṣ 76 DEC - o 3 56."7 (195) Fig. 8. Optical depth of G5.89 in the 9.7 µm silicate feature. The gray scale is indicative for the feature depth, dark means high optical depth. Contour lines are the same as in Fig. 1 whereas we measured 37 Jy at 1.8 µm (the [Neii] line should be included in Ball s filter width of 1%) and Faison et al. (1998) give 16 Jy at 1.5 µm. This underestimation of the 1.5 µm flux by Ball et al. (199) leads to a spectral index between 8.5 µm and 1.5 µm of virtually zero and thus to a serious overestimation of temperature. The total mass we derive by this procedure is M. Simply correcting for optical depths from Sect...1, we get M. When assuming the geometry of a half-sphere with radius. 5, we derive a dust density of gcm 3 ( gcm 3 in the optical depth corrected case.) The total flux at 11.7 µm used is 17 Jy, the total Q-Band flux 59 Jy. Of course we are aware that taking the κ ν from Ossenkopf & Henning (199) introduces some arbitrariness. The model of Ossenkopf & Henning (199) assumes a gas density of 1 5 cm 3 and coagulation for 1 5 years. To estimate this uncertainty, we calculated the dust mass for the initial MRN-distribution of dust, where no coagulation has taken place. Such a scenario might apply if the dust aggregates have been destroyed by the heat and the opacity is similar to that of the interstellar medium. Using the appropriate opacities leads to a total mass of M ( M when correcting for the optical depth). An additional complication is the possible impact of very small grains, which might be subject to quantum heating. Such grains can be heated stochastically by the impact of single photons and are not in thermodynamic equilibrium. However, we have no way to tell whether such grains are present in G5.89 and while modelling the source with radiative transfer (see Sect...1), it Fig. 9. K image of G5.89 as shown in Fig. 1. Contours denote the measured extinction in Brγ. The darker, dashed contour line marks the region where the Brγ flux is below the detection level, thus extinctions are only lower limits southwest of this line. The numbers give the mean Brγ extinctions in the corresponding regions.

9 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II 51 G Mass Distribution + Temperature REFERENCE POS. R.A. 17 h 57 m 6 ṣ 76 DEC - o 3 56."7 (195) Fig. 1. Distribution of the hot dust inside G5.89 (gray scale) and temperature distribution (solid contours) derived from the 11.7 µm and Q- band images. The logarithmic gray scale ranges from gcm (white) to.6 1 gcm (black). Contour levels are 15, 11, 115, and 1 K. No information on temperature or mass is available outside the grey area. For orientation, the contours of the cm radio map are also plotted as dotted lines. The result is only shown where both signals were above their corresponding 3σ levels G Photometries REF. POSITION: R.A. 17 h 57 m 6 ṣ 76 DEC - o 3 56."7 (195) Fig. 11. Locations of the apertures used for photometry. Results of the photometric measurements can be seen in Fig. 1 1 Fig. 1. Colour-magnitude diagram of the regions shown in Fig. 11. The arrows are pointing to the de-reddened locations of sources inside the area of our extinction map (see Fig. 9) in the diagram. The solid line to left of the diagram is the zero-age main-sequence, plotted for a distance of.6 kpc. The arrow pointing downwards from the top of this sequence gives 1 times the interstellar reddening vector towards G5.89 (Neckel & Klare, 198). Sources,, and 8 are not detected in the H-band. Error bars show 1σ errors of our photometry turned out that such an effect has very little impact on the fluxes at 11.7 and 1 µm in the case of G5.89. In conclusion, we can say that the total hot dust mass of M represents a lower limit with an uncertainty of about a factor of to Photometry Twenty-one areas were selected in the high-resolution H and K images for aperture photometry. Some of them contain point sources. All areas show remarkable deviations from their surroundings in either H or K. The locations and sizes of the photometric apertures are denoted in Fig. 11, while the results are presented in a colour-magnitude diagram shown in Fig. 1. The plotted magnitudes are derived from the integrated flux densities inside the apertures, no sky subtraction was done. Where the measured sources were inside our extinction map (Fig. 9), de-reddening was applied for a reddening vector after Rieke & Lebofsky (1985). At first glance, most of the sources scatter along a band from the lower left to the upper right of the diagram. The lower border of this band is given by the detection limits. No apparent systematics can be recognised and we hesitate to assign the label star to most of the measured sources for reasons given in the next section.. Discussion In this section, we are going to build a model of G5.89 and put it into the context of the formation of massive stars and UCHiis.

10 5 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II First we will deal with the foreground cloud of dust which is thought to be responsible for the appearance of G5.89. The presence of this foreground cloud together with the half-spherical appearance of G5.89 at almost all observed wavelengths leads us to build a model of the source which assumes a spherical shell of dust that is half blocked from view. We then argue that we also found evidence for the long-known outflow from G5.89, but not for a circumstellar disk which could serve as a driving mechanism for this outflow. Combining these findings with a look at the surroundings of G5.89, we argue that the object contains one single young massive star while at the same time it is obviously a part of a cluster of such stars..1. The surroundings of G Clouds of dust the foreground extinction The measured flux at 1.3 mm indicates that substantial amounts of cold dust are present at the location of G5.89. In Fig. 6, the 9% contour line of the dust contribution to the 1.3 mm flux shows that this emission originates clearly offset from G5.89 proper. After subtracting the free-free contribution, the remaining flux of 8.5 Jy should entirely arise from cold dust. For the following estimates, we assume that the dust emission is optically thin at 1.3 mm. Using the mass absorption coefficient at 1.3 mm given by Ossenkopf & Henning (199) and assuming a temperature of 3 K, we derive a total mass of 7.5 M. Assuming that this mass is concentrated in a core of 16 diameter, the corresponding dust opacity from Ossenkopf & Henning for 1 µm yields an extinction of about 58 mag at that wavelength. Thus, even if only a fraction of this measured column density is situated in front of G5.89, it explains easily why even at 1 µm we cannot see through the foreground dust. The similarity of the shape of G5.89 over the wavelength range between 1.6 and 1 µm means that the rim of this cloud has to be rather sharp. The column density needs to jump exactly across G5.89. The beam size of the 1.3 mm measurement does not allow us to measure the spatial distribution of the cold foreground dust accurately. However, Fig. 6 clearly shows the dark cloud south and west of G5.89 and its brownish rim stretching across the source at least at. µm. Furthermore, the extinction towards G5.89 was derived with high spatial resolution in Sect The result in Fig. 9 shows clearly that the extinction is rising in the expected direction until the NIR emission drops off completely. The cloud is also obvious by its blacking out of background stars, as seen in Fig. 6. We note that the orientation of the cloud s rim is consistent with the orientation of the source in the 1.3 mm map. The rims of this cloud appear relatively bright in all three narrow-band filters, but interestingly are weakest in Brγ. Although the widths of the filters and the calibration uncertainties of the images do not allow a quantitative analysis of such weak structures, this seems to indicate that the UV field or shocks by winds from nearby stars are sufficient to excite the H (1-)S1 line, but not to ionize the cloud material to a large extent. If the jump in density is as large as speculated above, UV emission cannot penetrate as deeply into the cloud as a shock wave triggered by winds might. Thus, a larger emitting region would be provided in H (1 )S1 than in Brγ and the former line would consequently appear brighter..1.. Star formation in G5.89 s neighbourhood In Paper I, we argued that G is a young cluster of massive stars still in its formation phase. From a comparison with the Orion BN/KL region (named after Kleinmann & Low 1967 and Becklin & Neugebauer 1967) we inferred that formation in clusters might be the rule rather than the exception for massive stars. The high-resolution images of G5.89 again reveal a number of unresolved sources embedded in extended emission. However, the situation is different: G5.89 is only.6 kpc away compared to 6.3 kpc for G Thus, we reach.5 times the linear resolution for G5.89 (1 AU in K ) compared to G If the two objects are indeed similar, as we will argue below and in Sect...1, G5.89 can for example correspond to one of the embedded radio sources in G In Paper I, we identified the stars by means of two-colour photometry. This method was again applied to G5.89 and the result can be seen in Fig. 1. Some sources are evidently stars such as source 1 (moderately reddened) and 1. From a comparison of Fig. 11 with Fig. 6, it becomes clear that the chain of sources belongs to the suspected rim of the foreground cloud. Thus, outstanding sources among them may just represent irregularities in the surface of this cloud. This is probably the case for sources 7 and 9 which have the same colour as 1 and are thought to be reflected light from that star. On the other hand, sources 6 and get quite close to the ZAMS when a de-reddening is applied as was derived for the rim of the cloud across the radio shell (see the typical de-reddening vectors in the figure). The latter two sources are thus probably stars. From their locations in the map and in the colour-magnitude diagram, sources 3 and 5 are candidates for high-mass stars. Some sources could be de-reddened directly because they are inside the area where Brγ and radio data exist. These sources have de-reddening vectors attached to them in the figure. It appears that sources 15, 16, and are non-stellar objects, while source 19 (and probably also 18) is obviously a foreground star. Thus, we have indeed a few candidate young stars close to G5.89, which are summarized in Table 3. More facts can be extracted from Fig. 6: An area of extended Brγ emission is seen around position (+17,- ), which is identical with the area of weak emission in Gomez et al. (1991). The distance of that region to G5.89 is roughly 5 so the whole complex has about the same linear size of.33 pc as G In Sect..3, we will explicitly compare G5.89 to the embedded sources in G The presence of the more evolved Hii Region G in a projected distance of 1 or 1.5 pc also represents a similarity to the situation for G In Paper I, we followed speculations by Wilner et al. (1996) that star formation spread from G to its two neighbouring UCHIIs at distances of about 5.8 pc. In that latter case however, G is thought

11 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II 53 Table 3. Presumed candidates for young stellar objects close to G5.89 Source No. Pos. Pos. Reason (see Fig. 11) R.A. De Close to B ZAMS-type stars with typical de-reddening vector, isolated source Close to B ZAMS-type stars with typical de-reddening vector Close to B ZAMS-type stars with typical de-reddening vector, colour renders reflection from source 1 unlikely Close to B ZAMS-type stars with typical de-reddening vector Close to early B ZAMS-type stars with typical de-reddening vector, fairly isolated. Table. Parameters used in the modelling of the source Central Object: Luminosity: Temperature: Distance: Dust Distribution: L = L T = K D = 6 pc Density distr.: ρ(r) r.6 Outer radius: Inner radius: Rout = 1 5 AU R in = 55 AU Dust Properties: Size distribution: a =.5.5 µm N(a) a 3.5 (MRN) Silicate / Carbon ratio: Si:C =.9 Optical data: Draine & Lee (198) Optical depth: τ 55nm = 1 Optical depth: τ. µm =8. Optical depth: τ 11.7 µm =.9 Optical depth: τ 1. µm =1.8 Total dust mass: 38 M to be the originator of star formation in the region, while in the case of G5.89, its neighbour is obviously more evolved and thus G5.89 itself would be the object where triggered formation of stars is taking place... The morphology of G The hollow sphere CWW modelled the emission from the dust around G5.89 to fit the observed spectrum of the source. Their model 5 became the de facto standard model for comparison with observational data. It was essentially recalculated in Faison et al. (1998), who claim their model differs from CWW in density and outer radius. However, the numbers given in their paper are the same as those in CWW. The model consists of a spherically symmetric dust distribution around a central star of spectral type O6 ZAMS. A fairly large cavity at the sphere s centre is depleted of all dust in order to fit the near-infrared part of the spectrum. Now that we have high-resolution information, we can compare not only spectral but also spatial information with these model calculations. The code we used for the radiative transfer (RT) calculations was developed by Manske et al. (1998) and is based on a method given by and described at length in Men shchikov & Henning (1997). The code was used in one dimensional mode as the spherically symmetric dust distribution requires no further refinements. To compare the results of the calculations to the observational data, we calculated simulated maps of the source at 1.6 µm,. µm, 3.5 µm, 1.6 µm, 11.7 µm, 1.8 µm and 1. µm. One half of the output maps were arbitrarily set to zero to account for the assumed foreground extinction which also covers the central object. The last step was to convolve the resulting maps with a PSF of 1 FWHM as a representative observational beam size. As an example, the calculated map at 11.7 µm is shown in Fig. 13a. To be able to compare the results to those of CWW, we used the same dust properties as were used for their model 5 and started off from the same geometrical assumptions. It turned out that we mainly needed to change the total mass of dust involved and the density distribution as well (slightly) the radius of the inner dust-free cavity. The parameters used for the best fit are summarised in Table. Fig. 13 shows the main results of the model calculations. We will now briefly discuss this figure and the main consequences for the very simple model of a spherically symmetric dust shell. Fig. 13b shows the resulting spectral energy distribution (SED). The solid line gives the resulting SED for the parameters listed in Table. Although these parameters are significantly different from those of CWW, the result looks almost identical to those of CWW s model 5. The diamonds show observed fluxes of the source within a mask that is shown in Fig. 13a (except for the 1.3 mm flux which is the total measured flux after subtraction of the free-free contribution). The crosses denote fluxes measured in the same aperture in the simulated maps of the model. The representation of the observed SED by the model SED is very good, except in the near-infrared. Fig. 13c compares the measured mass column densities of hot dust as derived in Sect The solid line gives the profile as measured in the map shown in Fig. 1. The profile was extracted in the following way: The source was assumed to be radially symmetric and profiles of 6 length were laid from the centre at R.A. 17 h 57 m 6 ṣ 76, DE in 19 directions from -5 to 5 around north (19 profiles in 5 steps). The arrangement of these profiles is sketched in part a of the figure, the centre was determined such that all extracted profiles do not shift with respect to each other by more than. 17 (one pixel). The profile shown is the mean of the 19 extracted profiles. The same procedure as described in Sect. 3.6 was applied to the

12 5 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II Fig. 13a f. Comparison between model calculations and observational data. a Simulated map at 11.7 µm. The lines show the extraction profiles that were used on the observed maps to extract the profiles shown in C F. Also shown is the aperture used for photometry on both the real and simulated maps. b SED. The solid line is the result of our calculations, The diamonds denote our observational data, the crosses are observations on the simulated maps. The squares are IRAS fluxes. The crosses denote fluxes measured on the simulated maps. All our own fluxes were measured inside the aperture marked in part A, except the 1.3 mm flux, which is the total measured dust contribution. c Radial distribution of the observed mass column density. The solid line is derived from the observations as described in Sect The dashed line again marks observations on the simulated maps. The dotted line denotes the real column density as integrated from the input dust distribution for the RT model. d Radial temperature distribution. The solid line represents an extraction from the temperature distribution shown in Fig. 1. The dashed line was extracted from a temperature distribution derived in the same way as described in Sect. 3.6, but from the simulated maps. The dotted lines denote temperatures of the different sorts of dust grains computed from the RT model. e Radial intensity distribution at 11.7 µm. The solid line is from our observation, the dashed line from the simulated map. f Radial distribution of the silicate optical depth as described in Sect. 3.. Again, the solid line represents extractions from observational data while the dashe line denotes an extraction from the simulated maps. The arrow in parts C F marks the radius of the inner rim of the dust shell in the model calculations. simulated 11.7 and 1. µm maps and the resulting extracted profile is denoted by the dashed line. The dotted line shows the real model mass column density in the line of sight, integrated from the input dust configuration inside the aperture region of Fig. 13a where hot dust is visible. The temperature distribution which results from the same procedure is compared in part d of Fig. 13. The solid line is again extracted from the contour map shown in Fig. 1, the dashed line from the result of the procedure applied to the simulated maps. The dotted lines show the temperatures of the different sorts of dust grains in the model. We note that the deviation of these grain temperatures from the colour temperatures derived from the maps at small radii is not an effect of optical depth, but is due to averaging the observed

13 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II 55 temperature along the line of sight. This effect is thoroughly discussed, e.g. in Schreyer et al. (1996). Fig. 13e compares the radial profile of the intensity distribution at 11.7 µm. The solid line gives the profile as measured in the map shown in Fig. 3, the dashed line represents the profile as extracted from the corresponding simulated map. Part f of Fig. 13 compares our pseudo silicate optical depth by showing the extracted profile from the map in Fig. 8 as solid line and the result of the same operation on the simulated 1.6 and 11.7 µm maps as dashed line. We will now briefly discuss the main aspects of the model: The main goal of the modelling was to show that this simple geometry can reproduce the radial intensity distribution. This is clearly shown in Fig 13a and e. The shape and the position of the maximum flux density is governed by the location of the inner radius of the dust shell and the exponent of the power law which determines the decrease in density towards the outer regions of the dust shell. The SED is practically identical to that in CWW s model 5. It fits the MIR observations nicely, as is especially visible by the marks for simulated and real observational data. Note that the solid line represents the total flux of the spherical dust shell, while the simulated observations rely on only half the flux measured through a virtual telescope with a 1 beam. The resemblance to CWW s model 5 is somewhat astonishing as our calculation relies on only about 3% of the mass compared to that of CWW. In addition, we used a density gradient exponent of.6 compared to CWW s constant density distribution. Note that when using CWW s original parameters for model 5 with our code, we get a much deeper silicate feature and an earlier and steeper falloff towards the NIR. Harvey et al. (199) pointed out that the overall mass doesn t play a decisive part in the modelling, which is true as long as only the SED is concerned. For their own models, they used density gradients of r.5 and r 3. With such steep gradients, we cannot fit our observed radial profiles. The radial profiles of density and temperature show basically that the simulated observations reproduce the real ones within a factor of two. As far as the mass is concerned, the radial dependence also follows nicely the model input quantities. When comparing the real temperature to the observed ones, one has to be careful that we never see only the inner edge but always an average along the line of sight (at best). The effects of the finite optical depth given in Table seem not to be serious. The optical depths in the table are valid for the line of sight towards the central source, whereas we are only dealing with the inner % of the dust shell as far as the hot dust is concerned. The radial dependence of the pseudo depth of the silicate feature is completely different in the model and the observations. While the observational data show that silicate optical depth is high in the same places were emission is generally high at all wavelengths, the τ computed from the RT model rises only outwards of 3 (78 AU). This behaviour of the model is easily understood in terms of line-of-sight effects. Contrary to this simple explanation, the observed high silicate depth at the inner edge of the dust shell indicates the non-applicability of spherical symmetry in this case: Starlight escaping through the channel of the outflow visible as a breach in symmetry in the radio maps (see below) and heating dust grains at the channel walls can generate MIR-emission at locations, where less and less forground dust exists to produce an absorption feature (The channel opening is assumed to be pointing north, exactly the region where our profiles are extracted). Thus the falloff of optical depth towards the outside of the shell is understandable only in terms of non-spherically symmetric dust-distribution. The low temperature of only K at the inner boundary of the dust shell implies that the dust-free cavity must be made and maintained by some mechanism other than the destruction of dust particles through mere heat. Surely the simple model of a spherical dust shell is far from perfect. The observed maps show that in the end the source is not spherically symmetric, even when assuming to see only half the sphere due to foreground dust. The radio maps show a channel aligned roughly in north-south direction and the outflow(s) from the object need openings in the shell to get out (see next section). Interestingly, most non-symmetric features in the near- and mid-infrared maps are located close to the northern end of the channel. This might indicate processes of scattering and reflection of light leaking out of the sphere through an opening. Figs. 1 and show that in the NIR the maxima in flux density are on both sides of the channel s opening, while in the MIR maps the maximum shifts towards the opening itself with longer wavelengths. The SED shows that in the NIR the predicted fluxes are too low by a factor of around 1 or more. This can also be explained by direct starlight leaking out of the shell through this channel. Because the northern end of the channel is pointing away from the observer as infered from the direction of the outflow (see nect section), this light could still not reach the observer directly. It would have to undergo one or more scattering processes, a fact which might also help to explain the extremely red colour of source 16 in Fig. 11. Albeit, we have seen that the spherical model can explain most of the observed properties and we now know that the basic asymmetry that only half the sphere is visible at infrared wavelengths is due to extremely massive foreground dust concentrations. This is exactly the scenario which was argued to be unlikely by Harvey et al. (199).... The outflow(s) of G5.89 G5.89 is thought to be the origin of one of the most massive outflows within the Galaxy. However, there is an ongoing discussion as to what the actual orientation of this outflow is and whether there is only a single one or if multiple outflows are responsible for the puzzling observational results. Among the first groups to examine this phenomenon were Harvey & Forveille (1988) who found a strong outflow in 1 CO(1-). They concen-

14 56 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II trated on a red wing towards the east of the source and a much stronger blue wing towards the southwest. However, their original data show two red wings, one of which is coincident with the source. The blue wing shows up about slightly west of south from the source. This indication for an outflow slightly off the N-S direction agrees much more with later findings than the originally anticipated E-W direction. Further interesting findings come from Zijlstra et al. (199) and Cesaroni et al. (1991). Zijlstra et al. (199) found indications for a N-S outflow from OH maser velocity fields. They derive the northern cone to be directed towards us, which is opposite to most other measurements where the blue shifted emission is detected south of the source. Cesaroni et al. (1991) mapped the outflow in C 3 S. They find the orientation of the flow to be slightly west of north with the blue wing south of the source. Their red wing centre also almost coincides with the source position. Acord et al. (1997) contributed high-resolution SiO data obtained with the VLA. They find an overall orientation of about 5 east of north with the blue-shifted emission just southwest of the source position and the red-shifted emission about to the northeast. They also speculate on the cause for the blueshifted emission to be much stronger and closer to the source than the red shifted counterpart and conclude that this is due to higher density of the surrounding material to the south. This results from their assumption that SiO is produced in a shock when the outflow hits the surrounding material. Our findings concerning the outflow can be drawn from Fig. 6: We assume that the H (1-)S1 line is due to shocks when the outflow hits the surrounding material. Three features can be seen in this line, indicated by red emission in Fig. 6. The first and strongest, marked A is in position exactly coincident with the centre of the blue shifted C 3 S emission found by Cesaroni et al. (1991). The counterpart ( C ) is about 5 north of Cesaroni s red shifted position. Feature C is quite weak and extended. With the assumption that H is shock excited, this supports the view that the density of surrounding material is much higher to the south than to the north. The feature A is about closer to the source s centre than C. The higher density towards the south is also consistent with our findings in Sect..1.1 and our assumptions on the general structure of G5.89. Another H feature, marked B originates from the star at position (-1,+ ). We do not have any colour indication on that star, but the feature might indicate that it is also a young object...3. No disk in G5.89? Yorke & Welz (1996) introduced an outflow driving mechanism that relies on the photo-evaporation of circumstellar disks in combination with a stellar wind. Their disk sizes are of the order of 3 AU (they are considering B stars!). We have no direct evidence for the presence of such a disk inside G5.89 and the necessary resolution of. 11 is inaccessible to us. However, the broken symmetry of the radio shell and the presence of the channel opening might point to a disk-like or toroid structure, as was first pointed out by Zijlstra et al. (199). This structure may well continue into a small-scale disk inside the otherwise dust-free cavity. A rough simulation of a configuration with an O6 ZAMS star surrounded by a disk of 6 M with a radius of 1 AU and an inclination of 3, shows that the disk contributes only between 1 µm and 6 µm to the spectral energy distribution significantly. However, since the flux never reaches the 1% level of the measured flux, it would not alter the measured (and modelled) SED, so we cannot detect such a disk inside G5.89 s dust shell just from the spectrum. At cm wavelengths, a disk comparable to the candidate disk of G might be visible in G5.89. Taking G339 s 3.5 cm flux of 1 mjy (Ellingsen et al., 1996) and accounting for the smaller distance and brighter central star, a comparable structure in G5.89 should emit around 19 mjy at 3.5 cm and thus be visible in the WC89 cm map. However, two arguments militate against the presence of such a disk in G5.89: The orientation of the disk plane should be east-west, as suggested by the outflow. An inclined disk should thus, according to simulations by Kessel et al. (1998) appear as a north-south elongated structure. Nothing like this is detected in the maps. A disk in east-west direction should absorb all UV photons emitted in its plane. Thus it would prevent ionization of the spherical shell to the east and the west of the central object. This is however exactly the direction where the ionization appears to be strongest. This strong ionization towards the east and west is on the other hand exactly what one would expect if an east-west oriented disk is being evaporated by UV radiation and disrupted by strong stellar winds (Hollenbach et al., 199). The pronounced structure of the ionization might thus indicate that the disk is either in a late stage of its destruction by the central star or has already vanished completely. As we practically rule out the presence of a disk inside G5.89, the question of a driving mechanism for the outflow (and indeed for the formation of the central star) remains. The possible answer is that a disk was present inside G5.89 and served as a driving mechanism for the outflow. It might be completely evaporated by now, while the outflow is still escaping the shell due to the high pressure of the ionized gas. The channel opening can serve as a collimation mechanism. Little is known about disks around massive stars, let alone their lifetimes. However, given the extreme youth of G5.89 the dynamical ages of the shell and the outflow are 6 yr (Acord et al., 1998) and 3 yr (Acord et al., 1997) either renders the lifetime of such a disk extremely short or questions the dynamical age as a reliable age estimator..3. G5.89 and the models of ultracompact H ii regions.3.1. Sketch We will first put all the discussed facts together to form a complete picture of the source. The spatial distribution of all the The Object G will be discussed in a forthcoming paper

15 M. Feldt et al.: High-resolution imaging of ultracompact H ii regions. II 57 Fig. 1. Sketch of G5.89 s configuration. The large figure is a side view where the observer is located to the right. The inset on the upper right gives the observers view of such a model. Only the ends of the outflow cones are visible here, as is expected for shocked H (1 )S1 emission. The dashed southern part of the ionized shell should be visible only at radio wavelengths. features discussed so far is shown in Fig. 1. It shows G5.89 as a spherical shell of dust surrounding the hot star with a dustfree inner cavity and two emerging outflow cones. The object is located at the rim of a large cloud of dense cold dust, which obscures its southern half from the view of the observer who is located to the right. One additional source is shown: This is an example for one of the surrounding stars, e.g. the wavefront reference star. It is illuminating the rim of the cloud and its wind presumably shocks H molecules there..3.. A strömgren sphere? The viewpoint summarised in Fig. 1 states that G5.89 is an ideal UCHii. An O6 ZAMS star has formed an ionized shell around it of roughly. 5 or.5 pc diameter. From WC89 we learn that this is a typical initial size of a Strömgren sphere. Intriguingly, the inner radius of our modelled dust shell is 55 AU or.5 pc. Thus, the ionized sphere is obviously identical with the dust-free inner cavity. We can easily see that this must be the case: The model predicts a dust density of gcm 3 at the shell s inner edge. Estimating the UV optical depth by τ UV =.8 κ K ρ dust L, () we obtain unity after a distance L =73AU or. 3. This uses a dust mass absorption coefficient κ K for. µm from Ossenkopf & Henning (1991), the factor of.8 converts the K optical depth to the desired UV depth and is taken from Mathis (199). 99% of all UV photons are absorbed 59 AU or. beyond the inner boundary of the dust shell. Thus, only the very inner edge of the dust shell can be ionized. On the other hand, we have learned that the ring structure at infrared wavelengths is due to a combination of the dust and temperature distribution and a line-of-sight effect through the hollow sphere. Since exactly the same structure is visible at radio wavelengths, it is unlikely that the dust-free, inner sphere is uniformly filled with ionized gas. Obviously, the gas (plasma) density is highest close to the inner edge of the dust shell. One more hint, that G5.89 is not a classical Strömgren sphere after all, is its expansion rate. It was measured by Acord et al. (1998) to be ±1 mas yr 1. Estimating the pressure contrast between the ionized gas (n 1 5 cm 3 from the electron density, T 1 K) and the innermost part of the dust shell (n 1 5 cm 3, T 3 K, both from the RT model), we derive a pressure contrast of about 3:1, i.e. the region should be freely expanding via an R-type ionization front. Using Eq. 9 from WC89, we derive an expected expansion rate of.8 mas yr 1 for an UCHiiof.5 pc initial radius. Thus, the expansion cannot be driven by pressure differences between ionized and non-ionized material alone. The stellar wind that disrupted the circumstellar disk may serve as an explanation here: From Puls et al. (1996) we learn that an O-star of K has typical mass-loss rates of M per year. This mass has a speed at infinity from the star of approximately 5 kms 1. Thus, the influx of kinetic energy into the hollow sphere is of the order of 1 37 Ws per year. The energy thermally stored in the ionized gas is of the same order of magnitude, so the wind may well play a decisive role in accelerating the expansion... The dust-free cavity A possible explanation for the rapid expansion would be if the inner surface of the dust shell is being photo-evaporated by the UV radiation from the central star as proposed by Hollenbach et al. (199) for disks. The inside of the cavity would in this