ASTRONOMY AND ASTROPHYSICS. Unveiling the disk-jet system in the massive (proto)star IRAS

Transcription

1 Astron. Astrophys. 345, (1999) Unveiling the disk-jet system in the massive (proto)star IRAS ASTRONOMY AND ASTROPHYSICS R. Cesaroni 1, M. Felli 1, T. Jenness 2, R. Neri 3, L. Olmi 4, M. Robberto 5,6, L. Testi 1,7, and C.M. Walmsley 1 1 Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, I Firenze, Italy 2 Joint Astronomy Centre, 660 N. A ohōkū Place, Hilo, HI 96720, USA 3 IRAM, 300 Rue de la Piscine, Domaine Universitaire, F St. Martin d Hères Cedex, France 4 LMT Project and FCRAO, University of Massachusetts, 630 L.G.R.C., Amherst, MA 01003, USA 5 Max-Planck-Institut für Astronomie, Königstuhl 17, D Heidelberg, Germany 6 Osservatorio Astronomico di Torino, Str. Osservatorio 20, I Pino Torinese, Italy 7 Division of Mathematics, Physics and Astronomy, MS105-24, California Institute of Technology, Pasadena, CA 91125, USA Received 18 December 1998 / Accepted 9 March 1999 Abstract. We present the results of line and continuum observations towards the source IRAS , performed at 1.3 mm and 3.5 mm with the Plateau de Bure interferometer, from 350 µm to 2 mm with the James Clerk Maxwell telescope, and at 10 and 20 µm with the United Kingdom infrared telescope. The results fully confirm the findings of Cesaroni et al. (1997), namely that IRAS is a very young stellar object embedded in a dense, hot core and lying at the centre of a rotating disk. The bipolar jet imaged by Cesaroni et al. (1997) in the µm H 2 line is seen also in the SiO(2 1) transition, which allows to study the velocity field in the jet. A simple model is developed to obtain the inclination angle of the jet (and hence of the disk axis), which turns out to be almost perpendicular to the line of sight. By studying the diameter of the disk in different transitions and the corresponding line widths and peak velocities, one can demonstrate that the disk is Keplerian and collapsing, and thus compute the mass of the central object and the accretion luminosity. We show that if all the mass inducing the Keplerian rotation is concentrated in a single star, then this cannot be a ZAMS star, but more likely a massive protostar which derives its luminosity from accretion. Key words: ISM: clouds ISM: individual objects: IRAS ISM: molecules radio lines: ISM 1. Introduction The study of the formation of massive stars is directly related to the investigation of sources believed to be the signposts of embedded luminous young stellar objects (YSOs). Such signposts are mostly ultracompact (UC) Hii regions, H 2 O masers, and IRAS sources selected on the basis of their far-infrared (FIR) colours (see Kurtz et al and references therein). In recent years, the attention has gradually shifted towards regions with higher density and temperature, which should represent the en- Send offprint requests to: R. Cesaroni (cesa@arcetri.astro.it) vironment immediately facing the embedded YSO. In particular, an ever increasing number of dense, hot molecular condensations have been found close to (but not necessarily associated with) UC Hii regions and coincident with H 2 O masers: these are named hot cores. The interest in such objects is justified by the fact that very likely they represent the site where high-mass stars are born (Kurtz et al. 1999). The relevance of the hot core fine structure in the study of massive star formation is one of the reasons which led Cesaroni et al. (1997) (hereafter CFTWO) to perform high angular resolution observations of the molecular clump associated with the luminous source IRAS , situated at a distance of 1.7 kpc. By mapping the CH 3 CN and HCO + emission with the Plateau de Bure interferometer (PdBI), they were able to resolve the inner part of the bipolar outflow and detect a hot core at the geometric centre of the flow. Both the core and the outflow are centred at the nominal position of the IRAS source. Interestingly, the hot core is affected by a velocity gradient perpendicular to the axis of the flow, which is suggestive of rotation. Also, H 2 emission at µm is seen along the axis of the flow, which is likely arising from shocked gas (see Ayala et al. 1998) in a jet propagating from the embedded YSO. CFTWO conclude that IRAS is very likely a disk-outflow system around a young massive (proto)star. In this work we present a follow-up of the study of CFTWO. The main goals of our project were: to map the hot core with higher angular resolution in the CH 3 CN lines, in order to get a better picture of the velocity gradient and confirm the rotating disk hypothesis; to study structure and kinematics of the jet by means of the SiO emission, which is known to arise from shocked molecular gas; to obtain high angular resolution images of the continuum emission in IRAS , from the millimeter to the mid-infrared (MIR). The latter results are important to assess which fraction of the FIR emission measured at low angular resolution by IRAS is coming from the hot core, and hence give a more accurate estimate of its bolometric luminosity. In order to attain these goals, we have performed a whole

2 950 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Table 1. Frequency set-up for detected molecular lines with the PdBI. DSB tuning at 1.3 mm was used Line Centre frequency Bandwidth Channel spacing a (MHz) (MHz) (MHz) (km s 1 ) 3.5 mm continuum H 13 CO + (1 0) SiO(2 1) mm continuum b , , , 8.6 CH 3CN(12 11) CH 3CN& CH 13 3 CN(12 11) CH 3CN(12 11) a the effective spectral resolution is about twice the nominal channel spacing b the 1.3 mm continuum was obtained by averaging one 80 MHz and two 160 MHz bandwidths in the lower side band Table 2. Instrumental parameters for the IRAM PdBI observations Parameter Value Centre position α(1950) = 20 h 12 m 41 ṣ 00 δ(1950) = Observing mode continuum+line; single side-band at 3.5 mm, double side-band at 1.3 mm Number of antennas 5 Baseline range m Band centre GHz; GHz Number of sections in the correlator 6 Primary HPBW 55 ;23 Synthesised HPBW (P.A.) 3.5 mm: (18 ) 1.3 mm: (49 ) Primary flux density and ( Jy at 3.5 mm; Jy at 1.3 mm) bandpass calibrators 3C 273 ( Jy at 3.5 mm; Jy at 1.3 mm) 3C 345 (4.1 Jy at 3.5 mm; 1.1 Jy at 1.3 mm) Phase calibrators ( Jy at 3.5 mm; Jy at 1.3 mm) ( Jy at 3.5 mm; Jy at 1.3 mm) set of observations using a variety of instruments from the millimeter (with the PdBI), to the sub-millimeter (with the James Clerk Maxwell telescope; hereafter JCMT), to the MIR (with the United Kingdom infrared telescope; hereafter UKIRT). The technical details of these observations are given in Sect. 2, while the results are illustrated in Sect. 3 and discussed in Sect. 4. Finally, the conclusions are drawn in Sect Observations and data reduction 2.1. Plateau de Bure interferometer The interferometric observations were carried out between December 1996 and March 1997 using the IRAM 1 five element array at Plateau de Bure, France (Guilloteau et al. 1992). The five 15 m antennas were equipped with GHz and GHz SIS receivers operating simultaneously with doublesideband (DSB) temperatures of 35 K at 3 mm and 50 K at 1.3 mm. The receivers were tuned single side-band at 3.5 mm and double side-band at 1.3 mm. The facility correlator was centred at GHz in the lower side-band at 3.5 mm and 1 IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain). at GHz in the upper side-band at 1.3 mm. We show the frequency set-up of the correlator and list the detected spectral lines in Table 1. The six units in the correlator were placed in such a way that a frequency range free of lines could be used to measure the continuum flux. The spectral resolution of the correlator is about twice the channel spacings given in Table 1. Source visibilities were phase calibrated by means of interspersed observations (every 20 minutes) of nearby point sources. The bandpass calibration was carried out in the antennabased manner. The flux of the primary calibrators was bootstrapped from IRAM monitoring measurements and used to derive the absolute flux density scale. Table 2 gives a list of the main parameters for our PdBI observations. The calibration and data reduction were made using the standard IRAM/GAG software. Continuum subtraction was performed in the U V plane by using the integral over the linefree channels of the 160 MHz units. Finally, channel maps were produced for all the lines. The conversion factor from flux to brightness temperature in the synthesised beam is 54.7 K (Jy/beam) 1 at 3.5 mm and 50.2 K (Jy/beam) 1 at 1.3 mm.

3 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Table 3. Approximate JCMT beam-sizes (full-width half maximum) at the observed wavelengths. These values do not include the error beams that appear at high frequencies. The table also includes the total integration time observed in each mode and frequency (including the time spent nodding) λ HPBW Integration time (sec) (µm) (arcsec) photometry mapping James Clerk Maxwell Telescope SCUBA The submillimetre observations were carried out with the Submillimetre Common-User Bolometer Array (SCUBA) (Holland et al. 1999) on the James Clerk Maxwell Telescope on Mauna Kea, Hawaii. SCUBA consists of three millimetrewave photometry pixels and two sub-millimetre arrays: a shortwavelength array containing 91 pixels and a long-wavelength array containing 37 pixels arranged in a close-packed hexagon. Both arrays have approximately the same field-of-view (2. 3) and can be used simultaneously. The data were taken in August 1997 with fully sampled maps generated at 3 sampling at 850, 450, 750 and 350 µm using the jiggle observing mode and photometry (single-pixel) observations taken at 1.35 and 2.0 mm. Azimuth chopping of 2 was used for the map data and 1 for the photometry data. Uranus was used for calibration (Orton et al. 1986). The zenith atmospheric opacity at 850 µm ranged from 0.36 to 0.40 during the observations. Integration times and beam sizes at these wavelengths are detailed in Table 3. The data were reduced using the SCUBA User Reduction Facility (SURF) (Jenness & Lightfoot 1998) with the outer ring of bolometers being used for sky removal as detailed in Jenness et al. (1998) United Kingdom infrared telescope Mid-infrared images of IRAS have been obtained using MAX, the mid-ir imager developed by the Max-Planck- Institut für Astronomie (MPIA) for the United Kingdom Infrared Telescope (UKIRT) (Robberto & Herbst 1998). The camera is equipped with a Rockwell International SiAs BIB array optimised for high-background applications. The allreflective camera optics provide a scale of 0.27 arcsec/pixel, corresponding to a field of view of Observations were performed under excellent weather conditions on the night of 29th August 1996, taking full advantage of the newly installed top-ring and the hexapod secondary mirror mount with tip-tilt adaptive control developed at the MPIA. Tip-tilt correction allows MAX to routinely attain diffraction limited conditions (0. 54 at 10 µm on UKIRT) on hour-long intergration times. We used a N-band filter with effective wavelength λ 0 =10.2 µm, half-width bandpass λ=5.2 µm, and a Q-band filter with λ 0 =19.9 µm and λ=1.9 µm. Airmass was 1.6. Our data acquisition strategy was the usual chopping and beam switching techniques. On both filters, the integration time per single frame was msec. Chopping at 2.2 Hz, we took 20 single frames per chop position. After 100 single frames on source, the telescope was pointed to the offset beam to correct for non-uniform illumination effects introduced by chopping. This cycle was repeated 4 times in the N band and 3 times in the Q bands, providing a total on-source integration time of sec and sec, respectively. Photometric calibration was obtained observing at various airmasses a set of red giant stars included in a list of faint ISO-PHOT calibrators, namely BS 6913 (spectral type K1III), BS 7648 (sp. type K5III) and BS 7635 (sp. type M0III). No attempt was made to determine the absolute position in right ascension and declination of the images. For the morphology of the region, which will become clear in the following, we assume that the position of the 10 and 20 µm peaks coincides with that of the 1.3 mm continuum peak. 3. Results In this section we present the results of the line and continuum observations towards IRAS The former consist only of the PdBI data, whereas the latter contain also the data obtained with the JCMT and UKIRT Molecular line data from the PdBI The SiO(2 1), H 13 CO + (1 0), and CH 3 CN(12 11) lines mapped with the PdBI had been also observed with the 30- m telescope by various authors. The availability of single dish spectra makes possible a comparison with those obtained by integrating the line over the whole emitting region imaged with the interferometer. This is done in Fig. 1. The absorption features affecting the H 13 CO + profile are due to extended emission which is lost in the high resolution map; on the contrary, all the SiO and CH 3 CN emission is recovered by the interferometer, which proves that it arises from a relatively compact structure. In Fig. 2 the same comparison is shown for all the K components of CH 3 CN(12 11), reinforcing the previous conclusion. A detailed description of the structure of the emitting regions can be obtained from the study of the different lines in the PdBI maps. These are presented in the following sections Maps in the SiO(2 1) line One of the goals of this work was to use the SiO emission to trace the spatial and velocity structure of the outflow seen by CFTWO in the HCO + (1 0) line. It is thus of great interest to compare maps in these two transitions, as done in Fig. 3. In the left panels we show the HCO + maps (contours) as in Fig. 7 of CFTWO, whereas in the right panels similar maps for the SiO line are presented. Note that for both transitions the maps in

4 952 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Fig. 2. Similar to Fig. 1 for all the CH 3CN(12 11) lines. The numbers above and below the spectrum indicate respectively the position of the different K components of CH 3CN and CH 3 13 CN Fig. 1. Comparison between single dish (thin line) and interferometer (thick line) spectra: the latter have been obtained by integrating the emission over the whole region covered by the single dish HPBW. The single dish (IRAM 30-m telescope) data have been taken from CFTWO, as for CH 3CN(12 11), and Acord (priv. comm.), as for SiO(2 1) and H 13 CO + (1 0). The dotted horizontal and vertical lines correspond respectively to F ν=0 Jy and V LSR= 3.5 km s 1. The dashed vertical lines indicate the values of V LSR separating the bulk emission from the inner wings and the inner wings from the outer wings (see CFTWO for a definition of the outer and inner wings) the top panels have been obtained by integrating the emission in the outer wings of the line (i.e. from 28 to 9 km s 1 and from2to21kms 1 ), whereas the bottom panels correspond to integrals in the inner wings (i.e. from 9 to 5 km s 1 and from 2 to 2 km s 1 ). For a definition of the inner and outer wings see CFTWO. For the sake of comparison, also the µm image in the H 2 line is shown (grey scale). A few interesting remarks can be made on the basis of Fig. 3. First of all, it is clear that the SiO emission traces the H 2 emission better than HCO + : this confirms (see e.g. Bachiller & Gutiérrez 1997) that the SiO molecule is closely associated with shocked gas, whereas HCO +, although tracing the flow, seems to arise from a larger region. We stress that such a conclusion holds not only for the low velocity (bottom panels of Fig. 3) but also for the high velocity gas (top panels), where HCO + still presents a larger structure than SiO. The different distribution of these two outflow tracers is even more striking in Fig. 4, where the bulk emission (i.e. from 5 to 2 km s 1 ) in the two lines is compared. There is little doubt that HCO + peaks at the core seen by CFTWO at the centre of the flow, whereas almost no SiO emission is detected at that position. This confirms that the SiO molecule is strongly enhanced in the shocked gas. Finally, by comparing the left with the right panels of Fig. 3, one can see that the velocity structure of the outflowing gas is quite similar in the two tracers. However, an interesting difference can be pointed out: while at high velocities (top panels) the blue and red lobes of the flow are well separated in both lines, at low velocities (bottom panels) such a separation is sharper in the HCO + map. In fact, in the SiO map both blue- (full contours) and red-shifted (dashed) emission is seen in each lobe of the flow: this makes the velocity reversal noted by CFTWO for HCO + (i.e. the exchange between blue and red lobes going from low to high velocities) less striking for SiO, and suggests that the SiO emitting gas is confined into a smaller solid angle. We shall devote Sect. 4.1 to a better discussion of this point Maps in the CH 3 CN(12 11) line As already found by CFTWO for the CH 3 CN(5 4) line, also the CH 3 CN(12 11) emission originates from a compact molecular core located at the nominal position of the IRAS source, namely at the centre of the outflow. Incidentally, we note that the CH 3 CN(12 11) lines are optically thick, as derived from the ratio between these and the corresponding transitions of the CH 3 13 CN isotopomer: for a relative abundance [CH 3 CN/CH 3 13 CN]=60 (see e.g. Wilson & Rood 1994), one obtains CH 3 CN optical depths ranging from 6 to 26 depending on the K component. The angular resolution achieved with the new observations ( 0. 7) is 5 times better than that obtained by CFTWO and comparable to the estimated core diameter: this should allow to barely resolve the core and confirm the velocity gradient seen in the CH 3 CN(5 4) line. In Fig. 5a we show the map of the CH 3 CN(12 11) line emission integrated under the K=0 and 1 components. Although most of the emission is coming from the central unresolved portion of the core, a minor fraction of it arises from a more extended region 3 in size, which looks marginally elongated in the NE SW direction. Such an elongation is consistent with that found in the unresolved core (see Fig. 5b) by applying the same method used by CFTWO to investigate the structure and

5 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Fig. 3. Left panels: maps of blue- (full contours) and red-shifted (dashed) HCO + (1 0) emission (from CFTWO). The emission has been integrated under the line in the velocity intervals from 28 to 9 km s 1 and from 2 to 21 km s 1 in the top panel, and from 9 to 5 km s 1 and from 2 to 2 kms 1 in the bottom panel. Contour levels range from 50 to 260 by 30 mjy/beam. The dashed ellipse in the bottom right represents the synthesised HPBW. Right panels: same as left panels, for the SiO line. Contour levels range from 6 to 56 by 10 mjy/beam. The triangles indicate the positions of the H 2O maser spots from Tofani et al. (1995). In all panels, the grey scale represents the H 2 line emission at µm (from CFTWO) Fig. 4. Top panel: map of the HCO + (1 0) emission (from CFTWO) integrated under the line from 5 to 2 km s 1. Contour levels range from 0.4 to 3.6 by 0.4 Jy/beam. Bottom panel: same as top panel, for the SiO(2 1) line. Contour levels range from to by Jy/beam. The dashed ellipses in the bottom right represent the synthesised HPBW. The triangles indicate the positions of the H 2O maser spots from Tofani et al. (1995) velocity gradient of the CH 3 CN emission on a scale smaller than the synthesised HPBW. This consists in making channel maps in the lines and fitting a 2-D gaussian to the map in each channel: in this way one can determine the peak position of the CH 3 CN emission at each velocity, namely at the V LSR associated with each channel. The peak positions thus computed are plotted in Fig. 5b: their distribution is elongated in the NE SW direction. The position uncertainty, represented by the error bars in the figure, has been set equal to 0.45 HPBW/(S/N) (see e.g. Zhang et al. 1998a) with S/N signal-to-noise ratio. One can then plot the velocity associated with each of these peak positions against the corresponding offset measured along the line of symmetry of the distribution in Fig. 5b: this is done in Fig. 5c, which shows a nice correlation and demonstrates the existence of a steady velocity gradient ( 4300 km s 1 pc 1 ) through the core. As in CFTWO, we stress that the trend shown in Fig. 5c is not a real rotation curve, because the HPBW ( 0. 7) is larger than the angular range shown in this figure ( 0. 4). This also demonstrates that to really see the structure of the gas around the star one needs still higher angular resolution than that attained by us. The previous results fully confirm the findings of CFTWO, thus supporting the rotating disk interpretation proposed by them. In Sect. 4.2 we shall discuss this topic in some better detail. Note that in Fig. 10 of CFTWO the size of the CH 3 CN distribution looks larger ( 1 ) than in our Fig. 5b ( 0. 4): this is due to the fact that the 2-D gaussian fit tends to privilege the strong, compact (and hence unresolved) core with respect to the fainter, extended emission around it (represented by the lowest contour in Fig. 5a). However, the velocity gradient seen in the

6 954 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Fig. 5a c. a Map of the CH 3CN(12 11) emission integrated under the K=0 and 1 lines. Contour levels range from 0.03 to 0.8 by 0.11 Jy/beam. The (0,0) position corresponds to the phase centre given in Table 2. The dashed ellipse represents the synthesised HPBW. b Peak positions measured in different velocity channels for all the CH 3CN(12 11) lines. The filled square indicates the position of the 1.3 mm continuum peak. The straight line denotes the symmetry axis of the distribution. Typical error bars are shown in the top right. c Channel velocities versus angular offset along the symmetry axis drawn in the middle panel. The dashed line indicates the position of the 1.3 mm continuum peak. Typical error bars are shown in the top left core and shown in Fig. 5c is also present in the faint extended CH 3 CN emission: for example, the velocities of the spectra at positions (+1,+1 ) and ( 1, 0. 5) in Fig. 5a are respectively 5.6 and 3.1 km s Maps in the H 13 CO + (1 0) line The u v coverage obtained in our observations turned out to be lacking of sufficiently short baselines to properly map the extended emission in the H 13 CO + (1 0) transition, as witnessed Fig. 6. Bottom panel: PdBI maps of the H 13 CO + (1 0) emission integrated under the line inner wings, namely from 9 to 5 km s 1 (full contours) and from 2 to 2 km s 1 (dashed). Contour levels range from 0.03 to 0.09 by 0.01 Jy/beam. The grey scale represents the emission integrated in the bulk of the line, i.e. from 5 to 2 km s 1. Top panel: maps of the H 13 CO + line emission integrated from 5 to 3.5 km s 1 (full contours) and from 3.5 to 2 km s 1 (dashed). Contour levels range from to by Jy/beam. The grey scale represents the H 2 line emission at µm from CFTWO. Note that the redshifted emission to the top right of the figure does not belong to the compact structure at the centre of the map and is probably associated with extended nebular material in outlying regions. The dashed ellipse in the bottom right of each panel represents the synthesised HPBW. The triangles indicate the positions of the H 2O maser spots from Tofani et al. (1995) also by the line profile shown in Fig. 1: this is seriously affected by absorption features due to the existence of extended structures that cannot be imaged by the interferometer in the configuration used. However, our purpose was to map the regions on scales not larger than the size of the HCO + outflow seen by CFTWO, namely < 20 : this is comparable to the maximum imageable size and hence the data are still worth for studying the structure of the flow and disk. For the sake of comparison with HCO + and SiO, one would like to produce maps of the H 13 CO + emission by integrating under the line in the same velocity intervals used in Fig. 3: however, the H 13 CO + (1 0) line is more narrow than the same line in the main species, so that no emission is detected in the outer wings, namely below 9 km s 1 and above 2 km s 1. Therefore, in the bottom panel of Fig. 6 we can show only the outflow maps obtained integrating the H 13 CO + line emission in the inner wings, namely the equivalent of the maps in the bottom panel

7 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS of Fig. 3. The agreement between the H 13 CO + and HCO + low velocity outflows is remarkable. The grey scale in the same figure represents the bulk emission (i.e. from 5 to 2 km s 1 ) in the H 13 CO + (1 0) line: this peaks at the hot core position. Thanks to the high spectral resolution used to observe this line (see Table 1), one can separate the bulk emission red-shifted with respect to the line centre ( 3.5 km s 1 ), from that blue-shifted. This is done in the top panel of Fig. 6, where we show maps of the H 13 CO + emission integrated in the velocity intervals from 5 to 3.5 km s 1 (full contours) and from 3.5 to 2 km s 1 (dashed), overlaied with the H 2 line emission from CFTWO. In this map the H 13 CO + emission concentrates in two regions, one associated with the core at the H 2 O maser position, the other (represented by the red-shifted emission to the top right of the figure) belonging to a more extended structure which is not of interest for the present work. The latter could be related to a large scale outflow revealed in recent CO(1 0) observations (Shepherd priv. comm.): this could indicate the existence of multiple outflow episodes originating from the same region, possibly from distinct YSOs. However, although important, a study of these phenomena requires more extended maps than those presented here and hence goes beyond our purposes. Confining thus our attention only to the blue- and red-shifted H 13 CO + emission at the centre of the map, we note that, unlike the outflow axis, oriented from SE to NW, in this case the blue-shifted gas lies to the NE and the red-shifted to the SW: such a structure resembles although on a much larger scale that seen in CH 3 CN. In conclusion, it is tempting to speculate that the H 13 CO + emission arises both from the outflowing gas (on a scale of 20 and in a velocity range of 11 km s 1 ) and from a flattened rotating structure (on a scale of < 10 and in a velocity range of 3kms 1 ) roughly perpendicular to the outflow axis. Such a structure, might correspond to the outer layers of the disk seen in CH 3 CN. We shall come back to this point in Sect Continuum data An updated version of the continuum spectrum of IRAS is shown in Fig. 7. With respect to Fig. 11 of CFTWO, the new figure contains the flux measured at 7 mm (3.1±0.8 mjy) by Hofner (priv. comm.) and those obtained in the present work at 1.3 mm (with the PdBI), in the submillimeter (with the JCMT), and in the MIR (with the UKIRT). As in CFTWO, we stress that the data in the figure have been taken with different angular resolutions. The estimate of the bolometric luminosity, obtained by linear interpolation of the points of the spectrum, remains approximately the same as in CFTWO, namely 10 4 L. However, the slope of the spectrum in the millimeter looks less steep (F ν ν 2.9 ), although still consistent with optically thin dust emission. Using 0.9 as the spectral index of the dust absorption coefficient, one can fit the spectrum of Fig. 7 with a grey-body of temperature 60 K (dotted line in the figure): this is much less than T k =200 K estimated by CFTWO for the core, thus demonstrating that a substantial fraction of the continuum emission towards IRAS Fig. 7. Spectrum of the continuum emission from IRAS The point at 7 mm (star) is a VLA measurement by Hofner (priv. comm.). The filled circles, squares and pentagons indicate respectively the PdBI, JCMT, and UKIRT measurements from this work. For the meaning of the other symbols see CFTWO. The dotted line represents a grey-body fit to the data, obtained with a temperature of 60 K and a dust absorption coefficient ν 0.9 arises from cooler regions more extended than the hot core. One can reach the same conclusion just noting that the 3 mm and 1.3 mm fluxes measured with the PdBI (filled circles in Fig. 7) are 3 5 times less than the corresponding values obtained with single dish telescopes (empty circles and filled squares). A direct measure of the distribution of the dust continuum emission is represented by the maps in the sub-mm obtained with SCUBA and discussed in Sect In the following we describe the results obtained with the PdBI, JCMT, and UKIRT PdBI continuum maps With respect to CFTWO, we have obtained two major improvements: a better angular resolution at 3 mm and a map at 1.3 mm. The latter presents the advantage of a HPBW half as that at 3 mm and a signal strength ten times stronger: this allows to obtain a good picture of the core structure, which looks barely resolved. In Fig. 8 we compare the map at 3 mm with that at 1.3 mm. In order to increase the sensitivity to extended structures the 3 mm map was obtained by merging the new data with those of CFTWO, which differ in frequency by only 3 GHz. One can see that 60% of the flux at 3 mm comes from the compact core at the centre of the map, while the rest originates from a faint halo surrounding it. Such a halo very likely extends over a larger region than that covered by us and probably explains the discrepancy found between single-dish and interferometer measurements (see Fig. 7). The 1.3 mm map looks similar to the CH 3 CN map of Fig. 5a and presents an elongated shape along the NE SW direction. This is consistent with the emission originating from a flattened structure, as expected in the disk model

8 956 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Table 4. Main parameters of the continuum and methyl cyanide emission in hot core. A distance of 1.7 kpc is used to estimate the linear diameters D Tracer α(1950) δ(1950) F ν FWHP Θ S D T SB T B (Jy) (arcsec) (arcsec) (AU) (K) (K) 3 mm continuum 20:12: :04: mm continuum 20:12: :04: CH 3CN(12 11) K=0+1 20:12: :04: a a average value over the line profile, from the blue wing of the K=0 transition to the red wing of the K=1. Table 5. Flux density and size measurements from SCUBA observations. Measurements were taken using photometry (single pixel) mode and mapping mode. Peak and integrated flux densities have been derived from mapping data. The quoted 3σ errors are a combination of the standard error and estimated calibration error λ Photom. Map (peak) Map (int.) FWHP (µm) (Jy/beam) (Jy/beam) (Jy) (arcsec) ± ± ± ± ± ± ± ± ± ± ± ±90 15 Fig. 8. Top panel: map of the 3 mm continuum emission towards IRAS Contour levels range from 0.6 to 14.6 by 2 mjy/beam. Bottom panel: same as top for the 1.3 mm continuum. Contour levels correspond to 5 and 5 to 125 by 20 mjy/beam. The dashed ellipses in the bottom right represent the synthesised HPBW. The triangles indicate the positions of the H 2O maser spots from Tofani et al. (1995) discussed by CFTWO, although much better angular resolution is needed to assess this beyond any doubt. In Table 4 we give the main parameters of the continuum maps, namely the peak position, the integrated flux (F ν ), the measured full width at half power (FWHP), the deconvolved diameter (Θ S ), the linear diameter (D), the brightness temperature measured in the synthesised beam (T SB ), and the value of this after correction for beam dilution (T B ). For the sake of comparison, also the same values for the CH 3 CN map of Fig. 5a are listed. All tracers peak essentially at the same position. The diameter in the 3 mm continuum is greater than that in the methyl cyanide line and even more in the 1.3 mm continuum: this is probably due to the extended structure previously discussed, which could not be imaged at 1.3 mm. The increase of T B from 3 mm to 1.3 mm is consistent with the fact that also the dust optical depth increases with frequency. Finally, it is worth noting that the brightness temperature of the CH 3 CN(12 11) line is an order of magnitude larger than that quoted by CFTWO for the (5 4) transition (see their Table 4): this proves that we are beginning to resolve the optically thick core, thus approaching the value of the excitation temperature estimated by CFTWO ( 200 K) JCMT data Table 5 details the multi-wavelength flux measurements determined from the JCMT data. These data show that there is much more flux in extended structure than picked up by the photometry data with approximately 4 to 5 times as much flux present in the map data than seen in a single beam. Although the source size determined from the interferometric observations ( 1 at 1.3 mm) suggests that the core is approximately a point source at all SCUBA wavelengths, the photometry observations show that there is still extended structure being detected by the longwavelength measurements. This can be seen in Fig. 9 where the emission at all wavelengths extends over a region much larger than the corresponding HPBW. Equivalently, one may note that the FWHPs listed in Table 5 are significantly larger than the corresponding HPBWs. It is also worth pointing out that, despite the similar angular resolution, the FWHP at 350 µm is slightly larger than that at 450 µm: this difference, if significant, suggests that the emission is getting optically thick at short wavelengths. Such an effect can be seen better in Fig. 10: here we show a

9 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Fig. 10. Radial profiles of the flux density at 350 µm (full line) and 450 µm (dashed) averaged over circular annuli centred at the position of the 1.3 mm continuum peak Fig. 9. Maps of the sub-mm continuum emission towards IRAS , obtained with SCUBA. Contour levels range from 8 to 56 by 8 Jy/beam, from 1 to 29 by 4 Jy/beam, from 0.25 to 6.55 by 0.9 Jy/beam, and from 0.2 to 5 by 0.6 Jy/beam respectively for the 350, 450, 750, and 850 µm maps. The cross marks the position of the 1.3 mm continuum peak. The dashed ellipse in the bottom right of each panel represents the HPBW comparison between the normalised radial profiles at 350 and 450 µm, obtained by averaging the flux density over circular annuli centred at the position of the 1.3 mm continuum peak. This comparison seems to indicate that the profile is broader at 350 µm than at 450 µm. Probably, the most striking feature of the maps in Fig. 9 is the almost perfect symmetry: the emission peaks at the core position and then falls down smoothly in all directions. This proves that the energy source heating the dusty cloud traced by the sub-mm emission is located inside the inner 10 of it. This is consistent with IRAS being the heating agent of the surrounding molecular, dusty clump. It is worth noting that the 850 µm map presents a north south elongation which cannot be seen at the other wavelengths: this is likely due to the contribution of the CO(3 2) line which happens to fall in the bandwidth used to measure the continuum emission. Such a contribution is estimated to amount to 5% of the continuum flux. Indeed, the existence of a bipolar molecular outflow in the N S direction on a large scale has been recently revealed by interferometric maps in the CO(1 0) transition (Shepherd priv. comm.) UKIRT images The IRAS data available for IRAS are affected by two major problems: the poor angular resolution and the nondetection of the region at 12 µm, where such a resolution would be best. Therefore, it is of interest to study the structure of the source at these wavelengths and determine the value of the flux in the MIR. To this purpose, we have used the MAX camera on the UKIRT to perform observations towards IRAS at 10 and 20 µm. The fluxes measured at 10 and 20 µm are respectively equal to 0.32 and 30 Jy. The 20 µm image is shown in Fig. 11; the 10 µm image is essentially identical, although the S/N is poorer. In the figure, we have overlaied the MIR emission with the images obtained by CFTWO in the H 2 line and in the broad K-band (2.2 µm continuum). As explained in Sect. 2.3, no astrometry was made to fix the absolute position of the MIR images and we assumed the 10 and 20 µm peaks to be coincident with the 1.3 mm continuum peak. Such an assumption seems sensible although arbitrary and allows a direct comparison between the structure of the MIR emission and that seen at other wavelengths. The 20 µm emission extends mostly over the same region as that occupied by the core seen in the CH 3 CN lines and in the millimeter continuum. However, one sees also a faint tail of emission elongated to the NW: interestingly, this tail overlaps very well with a peak of the H 2 line emission close to the core (see Fig. 11, top), whereas only some faint 2.2 µm continuum is detected at the same position (see Fig. 11, bottom). One may wonder whether the MIR emission from the tail is thermal in origin or just scattered radiation coming from the core. In fact, in the latter case the spectral index between 10 and 20 µm should be the same both for the core and the tail, whereas in the former the tail is likely to have a lower colour temperature than the core, which contains the embedded source of energy. The ratio between the 20 and 10 µm fluxes changes from 80 (corresponding to a colour temperature of 111 K) in the core to 110 (106 K) in the tail: such a difference is not large enough to unambiguously prove the thermal origin of the MIR emission in the tail, but nevertheless indicates that some difference does exist. In particular, using Eq. (31) of Goldreich & Kwan (1974) one can estimate the dust temperature of the tail: assuming a luminosity of 10 4 L for the embedded star and a distance of 2 from the star to the tail, one obtains 50 K, to be compared with a colour temperature of 106 K. Given the uncertainties and the rough correspondence between the colour temperature and the

10 958 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Fig µm image of IRAS (contours) overlaied with grey scale representation of the H 2 line (top panel) and K-band continuum (bottom) at 2.2µm (from CFTWO). Contour levels range from 1.4 to 7.3 by 1.18 Jy arcsec 2. The triangles indicate the positions of the H 2O maser spots from Tofani et al. (1995) Fig. 12. Position velocity plot along the axis of the flow for the HCO + (1 0) (top panel) and SiO(2 1) (bottom) lines. Contour levels correspond to 0.05, 0.05 to 4.45 by 0.4 Jy/beam for HCO +, and to 0.002, to 0.05 by Jy/beam for SiO. The crosses in the bottom left of each panel indicate the angular and spectral resolution. The horizontal and vertical dotted lines correspond respectively to the bulk velocity and to the peak position of the 3 mm continuum dust temperature, the discrepancy of a factor 2 does not seem dramatic and supports the hypothesis that in the MIR we are observing thermal emission from dust. We conclude that the previous findings are consistent with the suggestion of CFTWO that the H 2 line arises in shocked, heated gas along the jet feeding the outflow (see also Ayala et al. 1998): in this case, the MIR must be thermal in origin, being coincident with the H 2 line emission. Instead, no correlation between mid- and near-infrared continuum emission is seen because the latter is just scattered radiation along the walls of the cavity created by the jet. 4. Discussion and interpretation CFTWO concluded that IRAS consists of a diskoutflow system originating from a young early type massive (proto)star. Is this scenario confirmed by the present observations? And, if so, do they improve the understanding of the jet/outflow and disk structure? In the following we shall give an answer to these questions. In particular, in Sect. 4.1 we show how the new SiO measurements are fundamental to understand the geometry and kinematics of the outflow/jet, while in Sect. 4.2 we make use of the CH 3 CN(12 11) sub-arcsecond resolution observations to improve our knowledge of the disk dynamics The outflow/jet system As already noted in Sect and discussed in CFTWO, the morphology of the outflow in IRAS is very intriguing: the blue-shifted lobe of the flow corresponding to the inner blue wing lies to the SE, whereas the same lobe lies to the NW in the maps of the outer wings; a similar statement holds for the red-shifted lobe (see Fig. 3, left). The tentative explanation proposed by CFTWO and supported by models such as that by Cabrit & Bertout (1986) is that the axis of the outflow lies very close to the plane of the sky. As we shall see, such an interpretation is confirmed by the SiO observations. In order to better illustrate the kinematics and morphology of the outflow/jet system, we show in Fig. 12 the position velocity plots along the flow axis for the HCO + (1 0) and SiO(2 1) lines. The difference between the two plots is striking: although in both transitions the lowest contour describes an S-shaped pattern, the peaks in the two maps lie on opposite sides with respect to the bulk velocity. This can be seen for example looking at the peak to the NW, which is blue-shifted in SiO, but red-shifted in HCO +. Another difference consists of the fact that most of the HCO + emission arises inside about ±8 from the core position, whereas for SiO it is strong only outside that range. Finally, a similar difference holds for the velocity range, because the SiO emission extends up to ±25 km s 1 with respect to the bulk

11 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Fig. 13. Same as Fig. 12 for the SiO(2 1) line only. Contour levels correspond to 2.5, 2.5 to 44.5 by 6 mjy/beam. The dashed lines and ellipse identify the region of the plot where line emission is expected in the best fit model described in Appendix A velocity, whereas only very faint HCO + emission is detected beyond ±5 kms 1. A possible explanation for these characteristics is that the relatively wide outflow seen in HCO + is fed by the narrow, well collimated jet traced by the SiO and H 2 line emission. In this scenario the SiO molecules would be more abundant in the shocked layer of gas compressed and accelerated by the jet, whereas HCO + would dominate in the molecular gas further away from the shock. In other words, the molecular outflow would roughly consist of two components: one (the jet) ejected with high velocity in a narrow angle around the axis of the flow; the other expanding at lower velocity and much less beamed. The SiO abundance would be strongly enhanced in the shocked region between the two components and would hence trace the high velocity gas. On the other hand, the bulk of the HCO + line would trace the low velocity component, although also some faint high velocity HCO + emission is clearly detected: this probably traces the gas closer to the SiO emitting region. A similar example is given by the flow structure seen in the CO(2 1) line towards HH 211 (Guilloteau et al. 1997). In order to confirm the previous scenario, one should elaborate a detailed physical model describing the line emission in an expanding flow with given density and excitation temperature gradients. Such a model goes beyond the purpose of this work; however, it is possible to fit the S-shaped pattern of the SiO position velocity plot using a simple-minded approach. Our assumption is that the jet traced by SiO is conical in shape and that the gas in it is uniformly accelerated up to a maximum distance from the centre, where its velocity drops abruptly to zero. The latter hypothesis is justified by the clear existence of a bow shock at the end of the SiO jet (see Figs. 3 and 4). This simple model allows to easily reproduce the maximum extent of the emission pattern in Fig. 12 bottom, namely the shape of the faintest contour level. Note that the assumption of velocity proportional to the distance from the centre ( Hubble law) is consistent with what observed in other outflows such as that in HH 211 (see Guilloteau et al. 1997). The details of the model are given in Appendix A; however, it is easy to understand that the emitting region in the position velocity plot is included between two straight lines (corresponding to the expansion velocity along the surface of the cone) and an ellipse (corresponding to the maximum velocity reached by the expanding gas). The inclination of the straight lines and the size and eccentricity of the ellipse depend on the free parameters of the model, namely the semi-aperture angle of the cone, its inclination with respect to the plane of the sky, the maximum velocity of the gas, and the angular distance from the centre at which this velocity is reached. In Appendix A we demonstrate that these parameters can be reasonably determined from the observational data and we present the best fit obtained in Fig. 13. The most important result of the model is the accuracy of the determination of crucial quantities such as the inclination angle of the jet axis. This is due to the sensitive dependence of the fit on the input parameters. Clearly the inclination angle is very small, thus confirming the hypothesis of CFTWO that the flow axis lies close to the plane of the sky. Incidentally, we note that the faint emission seen in Fig. 13 along the line of sight through the centre at V LSR 70, 50, and 90 km s 1 is very unlikely to be blue- and red-shifted SiO. We believe that this is to be attributed to lines of other species such as C 2 H 5 CN and C 2 H 5 OH arising from the core. Finally, we note that the previous treatment cannot be applied to HCO +, which presents a much more complicated morphology. For this, a more complex model than the one presented here is needed Time scale and physical parameters of the jet As done by CFTWO for the HCO + outflow, we can use the SiO emission to estimate the relevant parameters of the jet. These are given in Table 6, where the mass was obtained by integrating the SiO line emission over the whole jet, in the outer and inner wings, i.e. from 28 to 5 km s 1 andfrom2to21kms 1. Optically thin emission was assumed, with an excitation temperature of 30 K as in CFTWO. The SiO abundance relative to H 2 was chosen equal to 210 9, similar to that found by Acord et al. (1997) for the flow in the high-mass YSO G The momentum and kinetic energy of the jet were estimated by summing the contribution of each velocity channel. The time scale, t, of the jet was calculated assuming uniform acceleration of the ejected gas, as suggested by the model previously discussed: in this case t =2R 0 /v 0, where R 0 =0.099 pc is the maximum distance from the centre reached by the ejected gas and v 0 =100 km s 1 is the velocity of the gas at that distance (see Appendix A for an estimate of these values).

12 960 R. Cesaroni et al.: Unveiling the disk-jet system in IRAS Table 6. Physical parameters of jet a Parameter Value b Mass (M ) 4.4 Radius (pc) Momentum (M km s 1 ) 45 (290) Energy (ergs) ( ) Time scale (yrs) 1900 c Momentum rate (M km s 1 yrs 1 ) (0.15) Mech. lumin. (L ) 2.6 (107) a assuming T ex=30 K and [SiO/H 2]= b values in parentheses are corrected for an inclination angle φ =9 of the jet axis with respect to the plane of the sky (see Appendix A) c obtained assuming a true jet expansion velocity v 0=100 km s 1 (see Appendix A) We note that all the quantities in Table 6 computed from the observed velocity (i.e. energy, momentum, and luminosity) must be corrected for the inclination of the flow axis: this can be obtained by dividing the momentum by sin φ and the energy and luminosity by sin 2 φ, where φ is the angle between the flow axis and the plane of the sky. In Appendix A we show that a reasonable guess is φ 9 and hence the correction factor amounts to 6.4 for the momentum and momentum rate, and 41 for the energy and mechanical luminosity The rotating disk As shown in Sect , the CH 3 CN(12 11) observations confirm that the hot core in IRAS is elongated perpendicularly to the jet/outflow direction and presents a velocity gradient along its symmetry axis. Also, we note that the peak of the millimeter continuum emission lies to a good approximation at the centre of such an elongated structure (see Fig. 5b). The CFTWO interpretation is that the methyl cyanide core is indeed a rotating disk around a high-mass (proto)stellar object. A similar conclusion is reached by Zhang et al. (1998b), who detect a rotating disk on a much larger scale ( 4, i.e pc) in the ammonia (1,1) and (2,2) inversion transitions. If we are really observing such a well defined entity and not just the resulting effect of two or more distinct cores unresolved within the PdBI beam, then one should expect to see systematic trends in other parameters than velocity. In other words, if one deals with a well defined symmetric disk with a (proto)star at its centre, then it makes sense to ask a couple of more questions, namely: is there any temperature gradient in the disk, due to the embedded central stellar object? and, is it possible to obtain some better information on the velocity field in the disk? In this section we address an answer to these questions Temperature gradient If a temperature gradient does exist, this should affect the distribution of lines with different excitation energies. In order to investigate this effect, we plot in Fig. 14 three quantities mea- Fig. 14. Plot of the peak velocity (top panel), FWHM (middle), and angular diameter (bottom) measured in each K component of the CH 3CN(5 4) (from CFTWO) and (12 11) lines, versus the energy of the corresponding lower level of the transition. In the bottom panel we plot also the points for the K=2 and 4 lines of CH 13 3 CN(12 11). Note that the number of points in the bottom and middle panels differs from that in the bottom panel because for a few K lines the estimate of Θ S was affected by too large an error to be reliable, although it was possible to estimate V LSR and V 1 from the spectrum. The dashed 2 and dotted horizontal lines in the bottom panel indicate the size and the associated error of the 1.3 mm continuum emission sured for each K line of the CH 3 CN(5 4) (from CFTWO) and (12 11) transitions, as a function of the corresponding excitation energy. These quantities are the peak velocity, the full width at half maximum ( V 1 2 ), and the angular diameter (Θ S)ofthe map obtained by integrating the emission under the line. Θ S is computed from the measured FWHP with a simple gaussian deconvolution; also, the position of the emission peak is the same for the different K lines, within the uncertainty. The values of V 1 2 and V LSR have been obtained with a gaussian fit: note that the formal errors of the fit are much less than the spectral resolution and must hence be taken with some care. The errors on Θ S, instead, have been computed from the FWHP obtained after increasing and decreasing the value of the peak emission by an amount equal to the 1σ RMS noise of the corresponding map. Looking at Fig. 14, one sees that Θ S is clearly decreasing with increasing excitation energy of the transition. This demon-