SpringerLink
Forum Springer Astron. Astrophys.
Forum Whats New Search Orders


Astron. Astrophys. 325, 282-294 (1997)

Previous Section Next Section Title Page Table of Contents

4. Results and discussion

4.1. Continuum maps

Water masers may reasonably be associated with early type stars: evidence for this can be found, for example, in Palagi et al. (1993), who show (see their Fig. 10b) that more than 90% of the H2 O masers in star forming regions have luminosities greater than [FORMULA], namely that of a B0.5 star (see Panagia 1973). Such early type stars are expected to develop UC HII regions. What we wish to verify is that H2 O masers show up in a very early stage of the evolution of an early type star. It is thus important to study the continuum emission possibly associated with the H2 O maser spots: if the associated HII region is pointlike or absent at all, then this indicates the embedded star(s) to be very young. Our results confirm this scenario. In fact, continuum emission is detected only towards G12.68 and G24.78: in the former, the HII region looks extended and not associated with the H2 O or OH maser spots (see Fig. 1); in the latter instead (see Fig. 2), the emission comes partly from a pointlike HII region (component A) coincident with the H2 O and OH masers, and partly from a resolved HII region (component B), offset by [FORMULA] from the H2 O and OH masers.

[FIGURE] Fig. 1. Contour plot of the radio continuum emission in G12.68-0.18. The filled triangles mark the position of the H2 O and the open squares that of the OH maser spots detected by Forster & Caswell (1989). The ellipse in the top right corner shows the HPBW. The contours are drawn at 7, 10, and 13 mJy beam-1
[FIGURE] Fig. 2. Same as Fig. 1, for G24.78+0.08. The two detected components (A and B) are labelled. The contours are drawn from 4 to 60 mJy beam-1 each 7 mJy beam-1. The HPBW is shown in the bottom left corner

In Table 3 we report the main parameters of the continuum emission in the two sources detected, i.e. the position of the peak, the flux measured at this position, [FORMULA], the corresponding synthesised beam brightness temperature, [FORMULA], the observed angular diameter at half power of the emitting region, [FORMULA], the deconvolved angular diameter, [FORMULA], and the integrated flux density over the whole emitting region, [FORMULA]. For the non-detected sources we give in Table 4 the upper limits corresponding to [FORMULA] RMS values, for the peak flux and corresponding brightness temperature.


[TABLE]

Table 3. Parameters of the continuum emission



[TABLE]

Table 4. Upper limits for the continuum emission at maser position


In Table 5 the main physical quantities of the related HII regions are given, namely: the physical dimension of the clumps, D ; the excitation parameter, U, the electron density, [FORMULA], and the Lyman continuum luminosity, [FORMULA], computed using the formulae of Schraml & Metzger (1969), under the assumption of optically thin emission at 1.3 cm; and the luminosity, [FORMULA], and spectral type of the star, derived from U using the tables of Panagia (1973). We point out that the values for G12.68 and G24.78 B are given only for the sake of completeness, since these HII regions are not associated with the maser environment.


[TABLE]

Table 5. Derived physical quantities for the HII regions


In the following, we discuss the constraints set by the previous results on the HII region possibly associated with the maser clusters. We treat the detected and non-detected sources separately, stressing that for "detected" we mean sources with continuum emission positionally coincident with the masers: according to such definition G12.68 belongs to the non-detected sample.

4.1.1. Continuum detected sources: G24.78+0.08 A

From Table 5 one sees that G24.78 A is very compact and dense, as expected for a very young HII region, whereas G24.78 B looks more extended. This is consistent with the VLA observations at 6 cm of Becker et al. (1994): in a [FORMULA] beam they detect a 32.2 mJy point source at [FORMULA] and [FORMULA]. Such position is almost coincident with that of G24.78 B, which proves that most of the emission at 6 cm originates from component B. However, since the angular resolution is comparable to the separation between G24.78 A and B, it is hard to estimate the contribution of component A to the measured 6 cm flux: if one assumes G24.78 A to be an order of magnitude fainter than B, i.e. [FORMULA]  mJy, then the spectral index between 6 cm and 1.3 cm turns out to be [FORMULA], which proves the 1.3 cm free-free continuum emission of G24.78 A to be optically thick. In turn, this indicates that the G24.78 A HII region is very small and compact. Analogously, one can demonstrate that G24.78 B is instead optically thin at 1.3 cm, with a turn-off wavelength between the thick and thin regimes of [FORMULA]  cm.

We conclude that the values of Table 5 (derived under the optically thin approximation) are only lower limits in the case of G24.78 A; in particular, the spectral type of the ionising star is likely to be much earlier than O9.5. This is in agreement with the luminosity of the IRAS point source associated with G24.78 (IRAS 18355-0713A; see Palagi et al. 1993), which amounts to [FORMULA] and corresponds to an O6 star.

4.1.2. Continuum non-detected sources

For all sources of our sample but G24.78, no continuum emission has been detected at the position of the masers. How can one use upper limits to set sensible constraints on the parameters of an undetected HII region? If as discussed above (see Sect.  4.1) H2 O masers are associated with embedded early type stars, then the lack of free-free continuum emission can be explained in three manners: (i) the embedded star(s) is so young that no HII region could yet develop; (ii) the HII region is too small (i.e. young) to be detected; (iii) the HII region is too extended and faint to be detected. Both (i) and (ii) are consistent with the scenario that we wish to prove; case (iii) instead would indicate that the HII region is old: we want to demonstrate that this is not likely for our sources.

For a spherical, homogeneous, isothermal HII region, with given electron temperature, the peak brightness temperature at a given frequency can be expressed as a function only of its Strömgren radius, [FORMULA], and of [FORMULA]. Once the distance to the source and the HPBW of the telescope are known, it is possible to compute the brightness temperature which would be measured with the synthesised beam, [FORMULA] (calc), as a function of [FORMULA] and [FORMULA], and compare it with the upper limits in Table 4. In Fig. 3 we plot curves of constant [FORMULA] corresponding to these upper limits. Only points of the [FORMULA] - [FORMULA] plane falling below such curves (i.e. satisfying the condition [FORMULA]) are acceptable. We note that each curve can be divided in three parts: the first, at low [FORMULA], is parallel to the y-axis and corresponds to the case of an optically thick and unresolved HII region; the second, at intermediate values of [FORMULA], is parallel to the x-axis and indicates that the HII region is still unresolved but optically thin; and the third, at high [FORMULA], satisfies the approximate relation [FORMULA], typical of a resolved optically thin HII region.

[FIGURE] Fig. 3. Curves corresponding to constant values of peak brightness temperature in the synthesised beam, [FORMULA], of a spherical, homogeneous, isothermal HII region. We have assumed an electron temperature of 9000 K and a gaussian beam with HPW of 1:005. [FORMULA] and [FORMULA] are respectively the Strömgren radius of the HII region and the Lyman continuum of the ionising star; note that [FORMULA] is an increasing function of [FORMULA]. On the right hand side axis the values of [FORMULA] corresponding to the spectral types as of Panagia (1973) are indicated. The contours correspond to [FORMULA] upper limits in units of [FORMULA], and precisely: 2.2 K for G12.68, 0.64 K for G16.59, 0.66 K for G23.01, and 0.61 K for G28.87

When looking at Fig. 3, one must take into account that we are searching for UC HII regions, namely for [FORMULA]  pc. The figure clearly shows that in all cases if any UC HII region exists at the position of the masers, then could not be detected if either too small ([FORMULA]  pc) and optically thick, or too faint ([FORMULA]  sec-1) and optically thin. In order to discriminate between these two possibilities, one can estimate the bolometric luminosity, which can then be converted into the corresponding [FORMULA] by means of the tables of Panagia (1973). Beside G24.78 (see Sect.  4.1.1), two more objects of our sample are associated with a source of the IRAS PSC (Palagi et al. 1993): G16.59 (IRAS 18182-1433) and G28.87 (IRAS 18411-0338). Incidentally, we note that the IRAS colours of these objects satisfy the criteria set by Wood & Churchwell (1989b) for identifying IRAS sources associated with UC HII regions: this strengthens the idea that we are dealing with embedded early type stars. The bolometric luminosity, [FORMULA], estimated by integrating the four IRAS fluxes, are [FORMULA] for G16.59, and [FORMULA] for G28.87. These imply [FORMULA]  sec-1 for G16.59 and [FORMULA]  sec-1 for G28.87. By assuming these two values for [FORMULA] in Fig. 3, one derives [FORMULA]  pc for G16.59 and [FORMULA]  pc for G28.87.

A word of caution must be spent on using the IRAS fluxes to compute [FORMULA]: in fact, it is well known (see e.g. Wood & Churchwell  1989aand Churchwell et al.  1990) that IRAS luminosities are usually higher by a factor [FORMULA] than the luminosities derived from the number of ionising photons required to explain the radio continuum fluxes of known UC HII regions. This means that the effective value of [FORMULA] may be smaller by factors of 10 and 100 respectively for G28.87 and G16.59: such a reduction is not large enough to invalidate our method for setting an upper limit on the HII region size.

Unfortunately, the same method cannot be applied to G12.68 and G23.01, since they do not have an IRAS counterpart. However, it seems plausible that also in these cases the non-detection of compact continuum emission might indicate that the HII region is still too compact or not even formed. In fact, VLA observations of G12.68 at 2 cm and 3.6 cm in the B-array configuration (Kurtz, priv. comm.) did not detect any continuum emission at the position of the masers, thus confirming our hypothesis.

We conclude that very likely no or a not-yet-developed UC H II region is associated with the environment of the H2 O masers.

4.2. Line spectra

With the Medicina antenna, NH3 (1,1) and (2,2) line emission has been detected towards all objects of our sample. Moreover, NH3 (3,3) line emission has been searched for, and detected, towards all the sources but G28.87, for which the (2,2) line turns out to be very faint.

Using the VLA, we have searched for NH3 (2,2) and (3,3) towards all five sources of Table 1, but G12.68, for which only the (2,2) line was observed with a shorter integration time than the other sources, due to time restrictions. In Fig. 4 the Medicina and VLA spectra are compared. The Medicina NH3 (1,1) spectra are shown at full resolution ([FORMULA]), whereas the NH3 (2,2) and (3,3) spectra have been smoothed to the VLA spectral resolution ([FORMULA] and [FORMULA], respectively) to allow a direct comparison. The VLA spectra have been obtained by integrating over the whole area where NH3 emission is detected, with the sole exception of G12.68 and G23.01: in these sources no emission is visible in the VLA channel maps of the NH3 (2,2) line and hence we have integrated the NH3 (2,2) emission over a large area ([FORMULA]  arcsec2) around the H2 O maser spots. The spectra obtained in this manner are also shown in Fig. 4: a faint feature is seen at the expected [FORMULA], so that we consider this a positive detection.

[FIGURE] Fig. 4. Comparison between spectra observed at Medicina and at the VLA. The Medicina NH3 (1,1) spectra are shown at full spectral resolution, whereas the NH3 (2,2) and (3,3) spectra have been smoothed to the VLA resolution. The VLA spectra have been obtained by averaging the emission over a suitable area (see text). The vertical marks indicate the positions of the hyperfine satellites

From the comparison between the Medicina and the VLA spectra we can immediately deduce that: (i) the hyperfine satellites (indicated by the vertical marks in Fig. 4) do not show up in the NH3 (2,2) and (3,3) spectra of Medicina, whereas they are detected in the VLA spectra of G23.01, G24.78 and G28.87; and (ii) the main line intensity ratios between the Medicina and the VLA spectra are [FORMULA] 1 for G12.68, G23.01, and G24.78, and [FORMULA] 1 for G16.59 and G28.87 (see Table 6). These findings suggest that a compact optically thick core is present in most sources, surrounded in some cases by an optically thin envelope, which in G23.01 and G24.78 contributes substantially to the NH3 emission seen by the single dish telescope.


[TABLE]

Table 6. Main line intensity ratios between the Medicina and VLA data


Tables 7 and  8 contain the observed ammonia spectra parameters: the name of the source, the line transition, the main beam brightness temperature ([FORMULA]), the LSR velocity ([FORMULA]) and the full width at half maximum ([FORMULA]) of each hyperfine component. The last column gives the total optical depth over all components of a transition ([FORMULA] ; see Ungerechts et al.  1986), calculated for those spectra where the satellites have been detected. The parameters of Tables 7 and  8 have been obtained using the TAUFIT program which takes into account the hyperfine structure of the lines, assuming gaussian profile for each component and relative intensities consistent with the local thermodynamic equilibrium (LTE) approximation. The data regarding the VLA observations have been obtained from the spectra at the ammonia peak position (see Sect.  4.3). It is worth noting that in Table 8 also the fit parameters relative to a second molecular component of G24.78 (hereafter indicated as G24.78+0.08 M) are presented: this corresponds to the ammonia emission associated with the water maser spots in the G24.78 region at [FORMULA] and [FORMULA] (see Sect.  5). Since no NH3 peak is seen at this position, we have used the NH3 spectra obtained by averaging the NH3 (2,2) and (3,3) emission within a radius of [FORMULA] centred on the H2 O maser position: this has been arbitrarily chosen equal to the angular radius of the ammonia core seen towards G24.78 A.


[TABLE]

Table 7. Results of fits to the spectra observed at Medicina



[TABLE]

Table 8. Results of fits to the spectra taken with the VLA. In all cases but G24.78 M the spectrum at the peak position in the map was used


From these results we conclude that we detect NH3 emission with the VLA towards all of the sources observed, although in one case (G12.68) the NH3 (2,2) line is detected only after averaging over a region a few seconds of arc in size around the position of the masers. In the next section we shall investigate the distribution of the ammonia emission in detail.

4.3. Line maps

The contour plots of the main line integrated intensity for all the detected sources are shown in Fig. 5. The positions of all the maser spots detected by Forster & Caswell (1989) are indicated in the contour maps by filled triangles (H2 O masers) and open squares (OH masers). Also shown are the positions of the [FORMULA] m sources with strong NIR excess detected by Testi et al. (1994b; 1995). The association between the H2 O masers and the NIR source has been stressed and discussed at length by Testi et al. (1994b; 1995); what Fig. 5 shows is that an even closer association does exist between masers and ammonia emission, indicating that maser emission is connected with the presence of compact molecular gas. A special case is represented by G24.78, for which one of the two groups of H2 O spots lies relatively far ([FORMULA]) from the NH3 emission peak, but still inside the ammonia clump. The occurrence of absorption and emission makes G24.78 quite an interesting case: this will be analysed in detail in Sect.  5. Only towards G12.68 was no NH3 core detected, but one must take into account that: (i) G12.68 is the most distant source in our sample; (ii) only the NH3 (2,2) line has been observed; and (iii) the integration time is [FORMULA] of that used for the other sources. However, since the integrated spectrum discussed in the previous section shows a faint NH3 (2,2) line, it is very likely that also in the case of G12.68 an ammonia clump would show up in a better S/N map, at the maser position.

[FIGURE] Fig. 5a and b. Contour plots of the integrated main line intensity for G16.59-0.06 and G23.01-041 for the lines detected. The filled triangles and open squares mark respectively the positions of the H2 O and OH masers from Forster & Caswell (1989). The cross represents the position of the [FORMULA] m source with strong NIR excess detected by Testi et al. (1994b; 1995); the size of the cross indicates the [FORMULA] positional uncertainty. The thick ellipse in the lower right corner of each image represents the HPBW. Contour levels are: -8, 8 to 32 by 4 mJy beam-1 for G16.59-0.05 NH3 (2,2) and (3,3); and -8, 8 to 40 by 4 mJy beam-1 for G23.01-0.41 NH3 (3,3)

[FIGURE] Fig. 5c and d. Same as Fig. 5a but for G24.78+0.08 and G28.87+0.07. Contour levels are: -12, -8, 12, 18 to 46 by 4 mJy [FORMULA] beam for G24.78+0.08 NH3 (2,2); -7, 13, 25 to 67 by 6 mJy beam-1 for G24.78+0.08 NH3 (3,3); and -8, 8, 10 to 18 by 4 mJy beam-1 for G28.87+0.07 NH3 (2,2) and (3,3)

For each ammonia transition, Figs. 6 and  7 show the maps of G24.78 and G28.87 obtained by averaging the channels in the main line (upper panels) and those in the four satellites (lower panels). One can see that the emission due to the satellites is detected from the region where the main line is strongest, thus confirming the existence of an optically thick core coincident with the H2 O masers.

[FIGURE] Fig. 6. Contour maps of integrated NH3 (2,2) and (3,3) main line (upper panels) and satellite (lower panels) emission towards G24.78+0.08. The grey scale represents the 1.3 cm continuum. Contour levels correspond to -6, 8 to 50 by 7 mJy beam-1 for the NH3 (2,2) main line and satellites, and to -7, 7 to 67 by 6 mJy beam-1 for the NH3 (3,3) main line and satellites. White contours represent negative levels

[FIGURE] Fig. 7. Maps of the integrated NH3 (2,2) and (3,3) main line (upper panels) and satellite (lower panels) emission towards G28.87+0.07. Contour levels correspond to -8, 8 to 18 by 3 mJy beam-1 for the NH3 (2,2) main line and satellites and to -8, 8 to 20 by 4 mJy beam-1 for the NH3 (3,3) main line and satellites

In Fig. 6 also the 1.3 cm continuum map is displayed (grey scale), showing the coincidence between the absorption in the NH3 lines and the continuum of the G24.78 A UC HII region. In Fig. 8 we plot the spectra of the NH3 (2,2) and (3,3) lines corresponding to the absorption peak: the high ratio between satellites and main line clearly indicates that the UC HII region is surrounded by optically thick material.

[FIGURE] Fig. 8. Spectra of the ammonia lines at the position of the absorption deep (i.e. that of the G24.78 A UC HII region). The dashed lines indicate the zero level. The vertical ticks mark the positions of the main line and hyperfine satellites

4.4. Derived parameters of the ammonia cores

From Boltzmann plots for the Medicina data, assuming that NH3 is populated according to LTE, one can estimate the rotational temperature and total column density performing a linear fit to the data (see e.g. Cesaroni et al. 1992). The results are presented in Table 9. One can see that the single dish measurements are sensitive to cool molecular gas with temperature of 10-20 K, which is very likely distributed over more extended regions than those seen in the VLA observations, as already noted in Sect.  4.2. This applies in particular to the NH3 (1,1) transition. In fact, the ratio between the NH3 (1,1) and (2,2) column densities indicates a lower rotational temperature ([FORMULA]) than the ratio between NH3 (2,2) and (3,3): this suggests that higher excitation transitions sample hotter regions, in agreement with the findings of other authors (see e.g. Cesaroni et al.  1992).


[TABLE]

Table 9. Results of linear fits to Boltzmann plots for the Medicina data


The previous conclusion is confirmed by the VLA data, as shown in Table 10 where we summarise the derived physical parameters of the ammonia clumps: observed ([FORMULA]) and deconvolved ([FORMULA]) angular diameter, physical diameter (D), kinetic temperature ([FORMULA]), ammonia column density ([FORMULA]), ammonia mass ([FORMULA]), virial mass ([FORMULA]), H2 volume density ([FORMULA]), and ammonia abundance with respect to H2 ([FORMULA]).


[TABLE]

Table 10. Parameters of the ammonia cores derived from the VLA data


The observed diameter of the molecular clump corresponds to the full width at half power (FWHP) of the maps in the main line. The true diameters have been estimated by deconvolution of a gaussian beam from a gaussian source. For G16.59 the average of the diameters measured in the (2,2) and (3,3) maps has been used, whereas for the other sources we give the diameter in the inversion transition for which the signal to noise ratio is higher: NH3 (3,3) for G23.01 and G28.87, and (2,2) for G24.78.

In order to derive the NH3 column density at the peak position ([FORMULA]) in each molecular core, we have used the method described, e.g., by Ungerechts et al. (1986) and the NH3 partition function. From [FORMULA] and the core diameter, D, it is possible to calculate the total number of NH3 molecules in the core (N), assuming the core to be spherical and homogeneous:

[EQUATION]

The ammonia mass ([FORMULA]) is then easily estimated from N. Note that, since G23.01 has not been detected in the (2,2) line, we have derived the total column density for this object from the (3,3) column density adopting LTE at temperature [FORMULA].

Assuming that the cores are virialised, one can estimate their total masses from (MacLaren et al.  1988):

[EQUATION]

where [FORMULA] is the ammonia linewidth derived from the VLA observations (Table 8). The NH3 abundance is then given by [FORMULA], while the hydrogen density ([FORMULA]) has been calculated from the virial mass and the clump diameter.

The kinetic temperature ([FORMULA]) reported in Table 10 is the VLA peak [FORMULA] of the NH3 line corrected for the beam filling factor, i.e. divided by [FORMULA]. This is in good agreement with the rotational temperature derived from the ratio between the NH3 (2,2) and (3,3) lines, as seen with the VLA. Note that, for all sources where both NH3 (2,2) and (3,3) have been detected, the peak values of [FORMULA] in these two lines are very similar (with a maximum difference of 15%): this indicates that the two transitions are optically thick and confirms that their brightness temperature is indeed a good estimate of the line excitation temperature, or, equivalently, of [FORMULA].

From Table 10, one sees that the ammonia column densities are [FORMULA] - [FORMULA]  cm-2, and the NH3 abundances are [FORMULA] - [FORMULA]. Such values are about an order of magnitude smaller than those derived by Cesaroni et al. (1994) for a sample of hot cores close to UC HII regions. However, these authors mapped NH3 (4,4), a higher excitation transition that very likely arises from denser hotter regions than the (2,2) and (3,3) lines, which can partly explain the discrepancy between their and our results. This seems confirmed also by the slightly lower [FORMULA] of our objects as compared to those of Cesaroni et al. (1994), who find [FORMULA] ranging from 50 to 165 K.

The fundamental conclusion we wish to stress is that our VLA observations confirm the existence of compact ([FORMULA] 0.1 pc) cores positionally coincident with the H2 O and OH masers: such cores turn out to be hot (40-90 K) and massive (30-900  [FORMULA]), thus suggesting that high mass star formation is going on inside them.

Previous Section Next Section Title Page Table of Contents

© European Southern Observatory (ESO) 1997

Online publication: May 5, 1998

helpdesk.link@springer.de