molecular cloud/hii region interfaces in the star …aa.springer.de/papers/7320003/2300972.pdfto...

Astron. Astrophys. 320, 972–992 (1997) ASTRONOMYAND

ASTROPHYSICS

Molecular cloud/HII region interfacesin the star forming region NGC 6357?

F. Massi1 ??, J. Brand2, and M. Felli31 Universita degli Studi di Firenze, Dipartimento di Astronomia e Fisica dello Spazio, Largo E. Fermi 5, I-50125 Firenze, Italy2 Istituto di Radioastronomia, CNR, Via Gobetti 101, I-40129 Bologna, Italy3 Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy

Received 13 February 1996 / Accepted 8 October 1996

Abstract. We have performed 12CO(1–0), 12CO(2–1), 13CO(1–0), 13CO(2–1), C18O(1–0), C18O(2–1), HCO+(1–0) andH13CO+(1–0) observations towards two selected areas in the starforming complex NGC 6357 with angular resolutions from 21′′

to 55′′. In particular, we have mapped the molecular gas aroundthe two Hii regions G353.1+0.6 and G353.2+0.9 in the 12CO(1–0) and 13CO(1–0) transitions with a resolution of ∼ 43′′. Thisimproves on the coarser 12CO(1–0) observations previously car-ried out by others. We have also studied the physical propertiesof gas along strips through the molecular cloud/Hii region in-terfaces.

For G353.1+0.6, the observations confirm that an ionizationfront is eroding a warm, dense molecular cloud located to thenorth of the optical nebula. The molecular gas appears frag-mented on a scale size > 0.5 pc and not all the components arerelated to the Hii region. There is evidence of a density increasenear the front and indications of temperature gradients towardthe exciting stars. This is further suggested by the presence of12CO(1–0) self-absorption produced by a cooler external layer.The isotopic ratio X(13CO)/X(C18O) decreases from the innerpart of the clouds towards the front, contrary to what is theoret-ically predicted and observed in many other regions. This maybe due to excitation or optical depths effects. An evolutionaryscenario is proposed where the exciting stars form at the edgeof a molecular cloud.

The morphology of G353.2+0.9 is rather different from whatpreviously believed. Only a weak “bar” of molecular materialwas found to the south of the sharp ionization front observedoptically and in the radio-continuum, and most of the molecularemission arises from regions behind or to the north of the Hiiregion. This indicates that we are viewing a late stage “blis-ter” configuration face-on. This region is fragmented on a scalesize > 0.5 pc, and a warm, dense and compact molecular frag-

Send offprint requests to: J. Brand ([email protected])? Based on observations collected at the European Southern Obser-vatory, La Silla, Chile?? Present address: Osservatorio Astronomico di Roma, Via Osser-vatorio 2, I 00040 Monteporzio C. (Roma), Italy

ment coincides with the elephant trunk visible in Hα images.Other clouds with high (∼ 40 K) 12CO(1–0) brightness temper-atures surround the nebula to the north. Around their peaks, the12CO(2–1)/12CO(1–0) main beam temperatures ratio is signifi-cantly< 1, contrary to what is observed in G353.1+0.6. In bothG353.1+0.6 and G353.2+0.9 a large range of radial velocities isobserved close to the ionization fronts, and the molecular struc-tures interacting with the ionized gas have virial masses greaterthan masses calculated under the assumption of LTE.

These observations have allowed us to better understand themorphology of the two regions and to sketch the physical prop-erties of molecular clouds exposed to UV radiation. Dynamicalinteractions between ionized and molecular gas are used in or-der to estimate the age (<∼ a few 105 yrs) of both nebulae. Alsothe effects of UV radiation in determining the morphology ofmolecular gas are considered. The heating sources of the molec-ular clouds are the early type stars of the Hii regions, at theedges of the clouds. The 12CO(1–0) opacity, τ , seems to affect12CO(1–0) main beam temperatures and an empirical relationbetween τ and the visual extinction AV has been determined.

Key words: ISM: clouds – ISM: Hii regions – ISM: NGC 6357– ISM: molecules – radio lines: ISM

1. Introduction

The Hii region/molecular cloud complexes NGC 6334 andNGC 6357, in the Sagittarius spiral arm, are very active re-gions of star formation in the Galaxy. Of the two complexes,NGC 6334 seems to be the most active in forming massive stars,but as demonstrated by Felli et al. (1990; hereafter Fea90) mas-sive star formation is presently taking place in NGC 6357 aswell.

Optical images show that NGC 6357 (a.k.a. RCW 131,W 22, Sh-2 11) comprises several distinct Hii regions (see e.g.Fig. 2 of Fea90) in different stages of evolution. Numerous Oand B stars are located in the region, which have been stud-ied by Neckel (1978, 1984) and Lortet et al. (1984), and there

F. Massi et al.: Molecular cloud/Hii region interfaces in the star forming region NGC 6357 973

are a number of thermal radio-continuum sources. The complexhas been (rather crudely) mapped in the far IR (FIR) and COby McBreen et al. (1983), who found widespread CO emissionat Vlsr∼ −4 to −1 km s−1, often peaking in intensity near theposition of the FIR sources.

The regions around three of the radio sources were studiedin detail by Fea90 with the VLA at 6 cm. One of the three re-gions (G353.2+0.7) showed no compact Hii regions, nor earlytype stars, and will not be considered further. The other two,viz. G353.1+0.6 and G353.2+0.9, were found to contain brightstructures on scale sizes from∼3′′ to several arcminutes, and arethe subject of the present study. G353.1+0.6 is a more evolvedHii region, ionized by visible early-type stars located outsidethe radio nebulosity but inside a faint optical diffuse nebulosity.Here most of the radio-continuum emission arises from the ob-scured region north of the nebula; the properties of several of theradio structures found with the VLA are consistent with theirbeing ionized skins of extended molecular clouds. No signs ofrecent star formation were found associated with this Hii region.G353.2+0.9, on the other hand, is the brightest component inNGC 6357 at all wavelengths studied, and was found by Fea90to be younger. The Hα and radio-continuum emission both showa sharp boundary in brightness towards the south, indicating anionization front seen almost edge-on, and suggesting the pres-ence of obscuring material to the south of the Hii region. In-side the nebulosity several indicators of recently formed brightstars, such as compact Hii regions and sources with IR excessemission were found. These are probably responsible for theexcitation of the nebula (see Fig. 7a in Fea90).

In view of the morphology of the Hii regions, it is of interestto study the molecular clouds associated with G353.1+0.6 andG353.2+0.9 in greater detail. In particular, a study of the Hiiregion/molecular cloud interfaces should allow one to directlycompare the conditions immediately adjacent to the ionizationfronts of two Hii regions with different evolutionary status andstar formation activity. In this paper we analyze observations ofseveral molecular species and transitions. We adopt a distanceD = 1.7 kpc to both regions, from early type stars (Neckel1978). The kinematic near-distance (∼1 kpc), based on the ra-dial velocity of CO (McBreen et al. 1983; this work) and radiorecombination lines (Wilson et al. 1970), is rather uncertainsince the longitude of these sources is within 7◦ of the direc-tion of the galactic center. In Sect. 2 we describe the observa-tions and the data reduction. The results for G 353.1+0.6 andG 353.2+0.9 are presented and discussed in Sects. 3 and 4, andthe main conclusions are summarized in Sect. 5.

2. Observations

Molecular-line observations were carried out between 5–11May 1990 with the 15-m Swedish-ESO Submillimeter Tele-scope (SEST) at La Silla, Chile. The SEST and its instrumen-tation have been described in detail by Booth et al. (1989).The receivers employed cooled Schottky barrier diode mixers;a high-resolution Acousto-Optical Spectrometer was used asa backend, with a total bandwidth of about 100 MHz and a

resolution of 43 kHz, corresponding to velocity resolutions of∼ 0.15 km s−1 at 89 GHz, ∼ 0.12 km s−1 at 115 GHz, and∼ 0.06 km s−1 at 230 GHz. The full width at half maximum(FWHM) beam width is 55′′ at 89 GHz, 43′′ at 115 GHz, and21′′ at 230 GHz, respectively.

The observations were made employing frequency-switching with switch interval depending on the line frequency:this interval is larger than the extent of the emission to avoidoverlap, but still small enough that the emission feature ispresent both in the signal and reference cycles. The data werecalibrated using the standard chopper wheel technique as de-scribed by Ulich & Haas (1976). Typical rms noise values are0.5 K for 12CO(1–0), 0.7 K for 12CO(2–1), 0.2 K for 13CO(1–0) and C18O(2–1), 0.6 K for 13CO(2–1), 0.1 K for C18O(1–0),HCO+(1–0), and H13CO+(1–0). Pointing was checked every 3hours on the SiO(v = 1, J = 2−1) maser sources W Hya and VxSgr, and was found to be accurate to within 5′′. The temperature-scale was regularly checked on M17SW (T∗A=33.3 K) and thelocal peak position in each of the two objects; the intensity cal-ibration is accurate to within ∼ 10%.

The observed line intensities are expressed as main beambrightness temperature (Tmb=T ∗A/ηmb), with ηmb between 0.70and 0.79.

The 12CO(1–0) and 13CO(1–0) transitions were used mainlyto determine the large scale distribution of the molecular gas,and as indicators of kinetic temperatures and column densi-ties. These data are complemented with 12CO(2–1), 13CO(2–1),C18O(1–0) and C18O(2–1) observations along strips through theionization front in order to study variations in physical param-eters, such as density, kinetic temperature, and isotopic ratiosacross the photon-dominated region (PDR). HCO+(1–0) andH13CO+(1–0) emission were observed at selected positions tofind the regions with higher density. For G353.1+0.6 observedpositions are offsets (in arcseconds) in right ascension (∆α)and declination (∆δ) from a zero point α = 17h22m17.s5 δ =−34◦18′10′′ (1950.0). For G353.2+0.9 the zero point is α =17h21m27.s0 δ = −34◦08′40′′ (1950.0). In both cases the zeropoint coincides with the position of the ionization front.

Observations of 12CO(1–0) and 13CO(1–0) were performedon a grid with a spacing of 40′′, and along two strips at ∆α =−80′′ (for G353.1+0.6) or ∆α = −40′′ (for G353.2+0.9) and∆α=0′′ cutting across the ionized gas. The other molecules andtransitions were observed only along the two strips, with stepsin declination of 20′′ or integer multiples thereof.

The 12CO(1–0) spectra required some care before folding,because the emission is quite extended and there are a num-ber of features (in the signal and reference cycles) at variousvelocities. Based on the difference in velocity between thesefeatures and the bulk of the emission, they do not seem to berelated to the regions of interest, but can affect the spectrum ofthe main emission after folding. The 12CO(1–0) spectra with afrequency offset of 20 MHz have the red side of the main lineaffected by a negative feature in the reference (‘dip’) related toan emission feature at Vlsr∼ −40 km s−1. In these cases, beforefolding we have used gaussian fits to remove the large dip nearthe main line; the same procedure was followed for small peaks

974 F. Massi et al.: Molecular cloud/Hii region interfaces in the star forming region NGC 6357

Fig. 1. All 12CO(1–0) spectra taken towards G353.1+0.6. Velocity andtemperature scale are indicated in the upper right-hand corner.

elsewhere in the spectrum which could produce dips in the mainline after folding.

Finally, (1–0) and (2–1) spectra of the various CO-isotopeshave been resampled to a velocity resolution of 0.12 km s−1 and0.06 km s−1, respectively, while the HCO+ and H13CO+ spectrahave been resampled to a velocity resolution of 0.15 km s−1.

3. Results and analysis

3.1. G 353.1+0.6

3.1.1. Morphology

Fig. 1 shows all the 12CO(1–0) spectra taken towards thissource, with the emission between −20 and 20 km s−1, while

in Fig. 2 and 3 we present all the spectra taken along the stripsat ∆α=0′′ and ∆α=−80′′, respectively. The bulk of the emis-sion is between Vlsr=−10 and +6 km s−1, and shows a rathercomplex behaviour, with shape and intensity (max. Tmb∼40 K)changing significantly over separations of a single beam. Thisbroad emission feature is likely to be made up of individual, nar-rower components, with different velocities which contribute atdifferent locations. We shall later discuss this broad emission interms of individual components. In addition there is weaker andnarrower emission at about−40 km s−1, with a maximum tem-perature of a few K, which does not change appreciably withposition. This latter component is probably not related to theradio source [the Vlsr of the ionized gas is −4.1 ± 0.9 km s−1

(H109α, Wilson et al. 1970)], and it will not be considered fur-ther.

Comparison of 12CO(1–0) integrated intensity with the redplate of the ESO/SRC Atlas (see Fig. 16 of Fea90) shows thatthe lower contours in the southern part of the molecular cloudrun parallel to the ionization front, which is north of the opticallyvisible diffuse emission, at the border of a region of strong ob-scuration. The 12CO emission peaks further north in the regionof obscuration.

Based on channel maps and velocity-position contour di-agrams, we have identified several distinct emission compo-nents (clouds) within the observed field. Fig. 4 shows velocity-declination contour diagrams of 5 transitions along the two stripsat ∆α = 0′′ and ∆α = −80′′. Individual clouds are indicated byletters. Their position on the plane of the sky can be seen fromintegrated intensity contour plots over the appropriate velocityintervals (see Fig. 5). All components are incompletely mapped,and all, with the exception of A, have their peak emission out-side the observed field. Note that the strong, narrow feature at∼7 km s−1 in the HCO+-diagrams is not of interstellar origin,but is probably an interference spike.

The narrow component D at 6 km s−1 (see Fig. 4) does notseem to be related to the ionization front. In fact it is visible overthe whole field, with Tmb<∼9 K and a line width ∼1.4 km s−1,thus resembling a quiescent dark cloud. It has been detectedonly in 12CO(1–0), 12CO(2–1) and 13CO(1–0) and the ratio of12CO(2–1) and 12CO(1–0) peak temperatures is on average<∼1,implying the emission comes from a cold, extended region witha relatively low density. The signature of this cloud has also beenfound through narrow absorption lines in H2CO (Whiteoak &Gardner 1974) and OH (at 1667 MHz; Goss 1968). This, to-gether with the radial velocity, forbidden according to the galac-tic circular rotation model, suggests that component D is due toa line-of-sight cloud, closer to the Sun than NGC 6357. Con-sequently, it will be excluded from further analysis. This cloudmight be identified as dark cloud Kh 348 (Khavtassi 1955).

Figs. 4 and 5 show that the molecular cloud is composed ofseveral distinct components. The most conspicuous structure iscomponent A, which appears as an extended bright (Tmb≤ 39 K)emission region centered near position (∆α,∆δ)=(0,80) andcovering the velocity interval from−12 to +2 km s−1. Cloud Aclearly lies north of the ionization front, but its lower isophotesin the southern part overlap the radio-continuum emission and


Fig. 2. Spectra of all observed lines in G353.1+0.6 along the strip at ∆α = 0′′. Transitions and temperature scale are indicated above eachcolumn; velocity ranges are indicated below.


Fig. 3. Same as Fig. 2, for the strip at ∆α = −80′′.

tend to run parallel with it, suggesting an interaction betweenmolecular and ionized gas. Component E represents an east-west filamentary structure, blue-shifted with respect to the peakvelocity of component A, and clearly overlaps the ionizationfront. In the east, component E dominates the emission withTmb∼21 K. This suggests that cloud E lies in front of the ion-ized gas and obscures it, while cloud A bounds the Hii region in

the north. Component C consists essentially of a peak near thesouth-western edge of the VLA image, and is red-shifted with re-spect to the peak velocity of cloud A. Although our observationsonly partially cover it, it seems that molecular emission withroughly the same velocity as cloud C may be associated withthe “West Arm” of Fea90 (see Fig. 5). Fig. 4 shows componentB with Tmb∼17 K north of component A. Since the correspond-


Fig. 4. Vlsr– ∆δ contour diagrams for the two strips in G353.1+0.6. Contour intervals (Tmb) are: for 12CO(1–0) and 12CO(2–1) steps of 3 K from2 K; for 13CO(1–0) steps of 1 K from 0.5 K; for 13CO(2–1) steps of 3 K from 2 K; for HCO+(1–0) steps of 0.5 K from 0.5 K. Clouds discussedin the text are indicated by letters.

ing peaks at ∆α=0′′ and ∆α=−80′′ have different velocities,component B may be composed of two clouds, to the east andwest, but we have no observations at ∆δ outside the two strips.However, both clouds are very far from the ionization front andprobably unrelated to it. Other weaker structures (F and H) liesouth of the ionization front and have low Tmb; it is thereforedifficult to say whether or not these represent small fragmentsassociated with the Hii region.

3.1.2. Self-absorption

At Vlsr∼ −0.7 km s−1 and around ∆δ ∼0′′ the 12CO(1–0) con-tours have a ‘dip’ in the line profiles over an area larger than thebeam. This is well illustrated in the Vlsr-∆α diagram of Fig. 6.This feature might be due either to two components with slightlydifferent velocities, or to absorption by cooler foreground ma-terial.

The 13CO spectra show no dip, and can be satisfactorilyfitted by a combination of several gaussian emission compo-nents, which favours a (self-)absorption interpretation. On theother hand, the dip is not evident in the 12CO(2–1) profile at∆α=0′′, where self-absorption effects are expected to be morepronounced (Phillips et al. 1981). However, the 12CO(1–0) and12CO(2–1) line profiles at the locations where the 12CO(1–0)clearly shows a dip can be best approximated by the sum of upto four gaussians, one of which is in absorption, with roughly

the same velocities in the two transitions, while emission-onlygaussians never give good fits. The negative gaussian needed tofit the 12CO(2–1) profiles is weaker than that for the 12CO(1–0)transition.

The self-reversal is strongest at (0,0). We note that the dipis also clearly visible in the HCO+(1–0) line. Unfortunately,observations have been made with different beam widths, andin particular the HCO+ and 12CO(1–0) spectra are more beamdiluted than those of 12CO(2–1).

Given the similarity between theVlsr- ∆α diagrams of Fig. 6and the results of model computations of the emission profilesexpected from a cloud surrounded by a colder shell (cf. Fig. 4 ofPhillips et al. 1981) we believe that the dip is due to colder ma-terial surrounding component A, rather than due to a foregroundcloud. Adopting a two-layer cloud model in which an absorb-ing layer with excitation temperature TSR and optical depth τSR

surrounds a cloud with brightness temperature TB (TB > TSR),the observed main beam temperature of the dip, TD, is given by(Phillips et al. 1981)

TD = J ′(TB) exp (−τSR) + J ′(TSR)[1− exp (−τSR)], (1)

where J ′(T ) = (hν/k)[F (T ) − F (TBG)], F (T ) =[exp(hν/kT )− 1]−1, and TBG=2.7 K is the temperature of thebackground radiation.

Taking J ′(TB) equal to the peak Tmb of the 12CO(1–0) tran-sition (39 K), an upper limit to TSR can be obtained assuming


Fig. 5. 12CO(1–0) integrated intensity contour plots of G353.1+0.6.Velocity ranges are indicated at the bottom of each panel. Contour val-ues, in the format low(step)high, are: 1(2)7, 10(10)50, 70(20)150 Kkm s−1. The 10 K km s−1 contour is drawn thicker. The clouds dis-cussed in the text are indicated by letters.

τSR � 1. We find TSR <23 K at (0,0) and TSR <31 K at (0,40),two positions that we shall use as examples. But obviously τSR

cannot be � 1, otherwise we would see the absorbing com-ponent in the 13CO(1–0) transition as well. Conversely, lowerlimits to τSR can be obtained from Eq. (1), assuming an exci-tation temperature TSR. Table 1 lists τSR for TSR= 5, 10, 15 K,the velocity dispersion δv of the absorbing layer in the 12CO(1–0) transition derived from the gaussian fits, and the FWHM∆v(13CO).

As noted above, the self-reversal is less strong in the12CO(2–1) than in the 12CO(1–0) transition. This implies thateither physical conditions in the absorbing layer cause the12CO(2–1) optical depth τ21 to be smaller than τ10, or the absorp-tion is produced by many small clumps and their contributionis greater in the 12CO(1–0) beam than in the 12CO(2–1) beam.As shown by Loren et al. (1981), if the kinetic temperature is5 K, then τ10 > τ21 for any density, and if Tkin=10, 15 K, then

Fig. 6a and b. Vlsr–∆α contour diagrams at ∆δ = 0′′ for G353.1+0.6.a (left) 12CO(1–0), b (right) 13CO(1–0). Contour intervals (Tmb) are insteps of 4 K from 2 K (a) and in steps of 1 K from 1 K (b).

Table 1. 12CO(1–0) self-reversals and line widths at two positions inG353.1+0.6.

∆α ∆δ ∆v(13CO) δv τSR(5) τSR(10) τSR(15)(”) (”) (km s−1) (km s−1)0 40 7.0 1.0 0.36 0.42 0.520 0 5.8 0.9 0.60 0.76 1.05

with densities in the absorbing layer of <∼500 cm−3, when τ10<∼1 we still have τ10 > τ21.

Since we detect the dip in the HCO+(1–0) transition, wecan obtain a lower limit to τSR from Eq. (1) by assuming thatTSR(HCO+)= 3 K, roughly the background radiation tempera-ture. At (0,0) we find τSR(HCO+) ≥0.94.

In summary, a clumpy layer surrounding cloud A, at least20 K cooler than the inside cloud, and with moderate opticaldepth, could explain the observed dip.

The 13CO(1–0) line profile at the position where the 12CO(1–0) dip is stronger can be fit by 2 main gaussian compo-nents and the dip is closer to the more red-shifted one (Vlsr∼−1.2 km s−1). The difference in velocity between the dipand the gaussian component is positive and of the order of0.5 km s−1. This could be indicative of a large scale motionof the external layers towards the inner parts of the cloud.


3.1.3. Variations in molecular emission across the ionizationfront

The total molecular emission along the line-of-sight towardsG353.1+0.6 is the result of the combined emission of severalindividual components. In separating that part of the emissionwhich is actually associated with the Hii region under study, itwould help if we could deconvolve each emission profile intogaussian components. We first fitted gaussians to the opticallythin 13CO(1–0), C18O(1–0) and C18O(2–1) emission profiles;then we tried to fit gaussians at the same velocities to the opti-cally thick profiles, including a negative component where nec-essary (see previous section). In general this procedure workedquite satisfactorally, except that the optically thick profiles couldnot always be fitted with gaussians at exactly the same velocityas those for the optically thin lines, because saturation affectsthe line shapes, and since the optically thick and thin emissionoriginate in different parts of the cloud. In the majority of casesthe velocity differences between gaussian fits to optically thinand thick profiles at the same position are within 0.5 km s−1,but sometimes differences of up to 1 km s−1 are found.

Although emission is found at velocities between −12 and+4 km s−1, it seems clear that there are two main components(at −5 and −1 km s−1) which are present along the wholeextent of the strips. Other components appear at a few posi-tions at higher and lower velocities along the strips. Between∆δ = −40′′ and ∆δ=80′′ (the VLA continuum emission isfrom ∆δ ∼ −20′′ to −80′′) the first moment of the emis-sion of any line lies at Vlsr∼ −3.5 km s−1 at ∆α=0′′, and atVlsr∼ −4 km s−1 at ∆α=−80′′. Note that the ionized gas is atVlsr=−4.1±0.9 km s−1, the same velocity as, or slightly bluerthan that of the molecular gas. To the south and north the cen-troid of the molecular emission moves to higher velocities (tobecome lower again further from the ionization front), due tothe appearance of other clouds at different velocities.

In Fig. 7a,b we show the integrated 12CO(1–0) and 13CO(1–0) emission between −6 and 0 km s−1 (component A), associ-ated with the Hii region, while in Fig. 7c,d we show the inte-grated emission along the two strips in α, for all the lines withsufficient signal-to-noise. Figs. 7a,b show that there is a sharpdecrease in integrated emission at the location of the ioniza-tion front, with contour lines running more or less parallel toit. From Figs.7c,d we see that this rapid decrease occurs in allmeasured lines, although much more so at ∆α=0′′, but also thatthe behaviour differs for the various tracers. At ∆α=0′′ bothC18O and HCO+ (i.e. the high density tracers) peak closer to theionization front than do 12CO and 13CO, indicating the front isassociated with a density enhancement, as would be expectedfor a shock front preceding the ionization front. This effect isnot so clear in the optically thick lines, because the high opticaldepth may mask the density increase. The strip at ∆α = −80′′

shows a much more gradual north-south decrease of the totalemission. At this ∆α, the radio emission detected by Fea90becomes bifurcated, as relatively weak radio-continuum emis-sion has been detected at (−80,−20), while the intensity of theemission increases again at position (−80,−80). At the latter

position Hα emission has been detected (see Fea90), and theradio peak has been interpreted as a local density enhancementproduced by the interaction between the stellar UV radiationand a molecular fragment, of which the outer layer is ionized.The fact that the HCO+(1–0) emission has a slightly slowerdecrease than the other transitions is perhaps due to a similardensity enhancement when moving towards the shock front or,alternatively, to its higher beam dilution.

3.1.4. Column densities

Column densities can be determined in three ways. The firstmethod uses 12CO(1–0) and 13CO(1–0) data, and yields LTEcolumn densities Nlte (see e.g. Martin and Barrett 1978; Harjuet al. 1990). We have adopted an H2/

13CO ratio of 5 × 105

(Dickman 1978).The second method is based on the LVG model of Goldsmith

et al. (1983) and uses the (2–1) and (1–0) brightness tempera-tures of any isotopic species of CO (in our case we shall use13CO) and an assumption of its abundance relative to H2.

Finally, column densities can also be derived from 12CO(1–0) data alone, by making use of the empirical relation betweenWCO≡

∫T (12CO)dv and N (H2)

N (H2) = X ×WCO (2)

with X=(2.3±0.3)×1020 cm−2(K km s−1)−1 (Strong et al.1988).

For each component we have calculatedNlte using the gaus-sians fitted to the 13CO(1–0) line profiles. The excitation temper-atures have been determined from 12CO(1–0) by superimposingthe 13CO(1–0) gaussians on the 12CO(1–0) spectra and readingoff the 12CO(1–0) Tmb at the peak velocity of each gaussian.The same has been repeated for 12CO(2–1) in order to calcu-late 12CO(2–1)/12CO(1–0) for single components. Where pos-sible, we have checked that the column densities obtained usingTex’s from optically thin 13CO or C18O line ratios are similar.We have also calculated total column densities from 12CO(1–0)integrated emission, using Eq. (2) with the value of X givenabove, and compared them both with the sum of Nlte along thesame line of sight and with the column densities obtained fromthe LVG model. C18O(1–0) LTE column densities were trans-formed into N (H2) by using the conversion given by Frerkinget al. (1982). In all cases we found values to be the same withina factor <∼ 2–3, with greater differences at positions of low col-umn density. The mean ratio of WCO column density to totalNlte (derived from 13CO summing all contributions of singlecomponents along the line of sight) is 4± 4 but only in few (6out of 43) locations it is > 5 or even much greater. The totalWCO mass in the field, including a 1.36 correction for helium, is3000 M�, while the total LTE mass, also corrected for helium,is 1600 M�. The total LTE mass was obtained correcting alsofor locations with no 13CO(1–0) observations using a LTE-WCO

relation derived from a linear fit to all available couples of data.Fig. 8 shows Nlte of component A along the two strips,

as a function of ∆δ. Note that component A has been decon-volved in two subcomponents at ∼ −5 and ∼ −1 km s−1 (see


Fig. 7. a,b Contour plots of 12CO(1–0) and 13CO(1–0) integrated emission (−6 ≤Vlsr≤0 km s−1) from component A. Contour values are in stepsof 20 K km s−1 from 20 K km s−1 for 12CO(1–0), and in steps of 5 K km s−1 from 5 K km s−1 for 13CO(1–0). The dashed lines indicate theionization front. c and d Distribution of the normalized integrated intensity along the two declination strips for the various transitions observed.Codes are indicated below the panels. Dashed lines as in a and b.

Sect. 3.1.3). In both strips the column density peaks north of theionization front, but at ∆α = 0′′ that occurs nearer to the frontandN (H2) suddenly decreases when approaching it. Maximumvalues are of the order of 1022 cm−2 and are clearly greater at∆α = 0′′. Component E (not shown in Fig. 8) peaks at the ion-ization front with column densities of the order of 1021 cm−2

and up to 1022 cm−2 at ∆α = −80′′, where it dominates theemission. Component C (not shown) peaks south of the ion-ization front with N (H2) of the order of 1021 cm−2. Both Eand C show a limited extent in declination, so their emissionis probably somewhat beam diluted and their column densitiesunderestimated.

Lower resolution IR maps (see e.g. McBreen et al. 1983;Persi et al. 1986) are quite similar to low resolution radio-continuum maps (see e.g. Schraml & Mezger 1969); since IRemission traces dust, this suggests that the total (molecular,atomic and ionized) hydrogen column density can be quite high

also at the ionization front. PDR models (Tielens & Hollen-bach 1985; Hollenbach et al. 1991) indicate that CO forms atAV ∼ 2 mag; assuming

N (HI + HII + 2H2) ' 2× 1021AV cm−2 mag−1 (3)

(e.g. Bertoldi & McKee 1992), this impliesN (HI+HII+2H2) '4 × 1021 cm−2. Since in G353.1+0.6 the PDR is illuminatedfrom below and is viewed almost edge-on, the width of the VLAimage of Fea90 at ∆α=0′′ (∼60′′ or ∼0.5 pc) gives an upperlimit to the depth of the region where CO is dissociated. Themean density of this region is then∼ 2.6×103 cm−3; assumingan extent along the line of sight of the order of the size of A(1 pc), we find that the total hydrogen column density is quitehigh, of the order of ∼ 8× 1021 cm−2.


Fig. 8. LTE column density of the two main subcomponents of A: A1(Vlsr∼ −5 km s−1; filled squares) and A2 (Vlsr∼ −1 km s−1; opensquares), along the two strips in G353.1+0.6.

3.1.5. Masses and mean densities

To obtain the masses of the single components we have consid-ered contour plots of the area of all gaussians with roughly thesame central velocity on the plane of the sky, and have deter-mined the size of each component as 2R = D

√Θ1Θ2, where Θ1

and Θ2 are maximum and minimum angular sizes (in radians)at half maximum integrated intensity and D is the distance (inpc). The mean densities have then been calculated both from theH2 mass divided by the volume of a sphere of radiusR and fromthe peak column density (Nlte) divided by linear size, assumingthe extent along the line of sight to be equal to 2R. The twomethods give similar values (within a factor of 3 in the worstcase). The physical parameters are given in Table 2. We havealso calculated the virial masses using the relation of MacLarenet al. (1988) for a homogeneous sphere and the FWHM of thegaussians fitted to the 13CO(1–0) line profiles. They tend to besystematically larger (by a factor of 2.4–12) than LTE masses,which is not surprising if we consider that ∆vmay not representonly gravitational interaction.

3.1.6. Temperatures and relative abundances

The excitation temperature of 12CO(1–0), Tex, can be used asa first approximation of Tkin because 12CO(1–0) thermalizes atlow densities (∼ 100 cm−3). The warmest component is A, withTex∼42 K at (0,40) which rapidly decreases when approachingthe ionization front. Tex can also be obtained from the (2–1)/(1–0)-line ratios of optically thin isotopes (e.g. Levreault 1988);generally 13CO and C18O line ratios give lower values thanthose of 12CO, but while the 12CO brightness temperature of Apeaks at (0,40), its C18O excitation temperature clearly peaks at

Table 2. Physical parameters of components related to the ionizationfront of G353.1+0.6. A distance of 1.7 kpc has been assumed. All LTEmasses are corrected for He. All densities, unless explicitly stated, arederived from H2 mass divided by volume.

Component Central Size Mass Mass MeanVelocity (2R) (LTE) (Virial) Density

(km s−1) (pc) (M�) (M�) (cm−3)

A1 ∼ −5 1.0 850 2000 3× 104

A2 ∼ −1 0.8 370 1700 3× 104

E ∼ −6÷−8 > 0.7 > 70 > 750 8× 103(a)C ∼ 3 > 0.5 > 27 > 200 6× 103

a Quoted density is maximum Nlte divided by 2R.

(0,0) where the integrated emission of the high density tracers(including C18O) also peaks. At ∆α = −80′′, all Tex’s (12CO,13CO and C18O) are smaller and decrease more gently, and alsocolumn densities are lower suggesting that smaller Tex’s aredue to a weaker coupling between radiation and gas at lowerdensity. Clearly, the kinetic temperature cannot have the samebehaviour as the 12CO(1–0) excitation temperature (i.e. decreas-ing towards the edges) since at the cloud border the beam fillingfactor becomes dominant.

Abundances were derived from the LVG model of Gold-smith et al. (1983). We have adopted a (constant) kinetic tem-perature of 50 K, which is the model value that allows to re-produce the observed peak temperatures of A, C and E. From12CO, this yields densities of 103−104 cm−3. The same valueshave been obtained from 13CO and C18O, although C18O showsa density increase at (0,0), further substantiating what was al-ready suggested by the integrated emission of the high densitytracers. Using the FWHM of the gaussians as a first approxima-tion of the velocity field, we also estimated abundances of ∼ afew 10−5 for 12CO, ∼ 10−6 for 13CO, and ∼ 10−7 for C18O,which are more or less “standard values” for a molecular cloud(see e.g. Dickman 1978, Frerking et al. 1982). It is uncertainwhether abundances decrease towards the edges of the clouds,due to the different beam dilution of (2–1) and (1–0) transitions.The LVG column density is generally lower than, though withina factor of 2, the corresponding LTE values.

3.1.7. Optical depths

The degree of saturation of the CO emission can be assessedfrom line ratios (e.g. Levreault 1988). Along both strips wefind R1 ≡T[12CO(2–1)]/T[12CO(1–0)]∼ 1, with an increasetowards the edges of the field (where, however, the signal-to-noise ratio is smaller [between 1 and 5]). The mean value is1.2 ± 0.2 for the subcomponent of A at ∼ −5 km s−1, butdoes not change if we include all available data. We have alsoexamined the dependence of R1 on H2 column density or, al-ternatively, on AV . In fact, as indicated by Eq. (3), the twoquantities are directly related (see also Harjumpaa & Mattila


1996). AV was obtained from Eq. (3) using only the H2 con-tribution to N. R1 values tend to cluster between 0.9 and 1.2,especially at highAV . The highestAV ’s are at the cloud center,so they are less likely to be affected by beam dilution differ-ences in 2–1 and 1–0 transitions; the dispersion is consequentlylower. Ratios R2 ≡T[12CO(1–0)]/T[13CO(1–0)] are <∼5 alongthe strip at ∆α=0′′ and <∼10 along the strip at ∆α = −80′′.These values are typical of optically thick 12CO(1–0) emission.For 13CO we obtain R1 >∼2 along the two strips (indicating op-tically thin emission), except at ∆α=0′′ between ∆δ=40′′ and−40′′, whereR1 ' 1.2−1.4. There 13CO(2–1) is at most mod-erately optically thick. In fact these observed line ratios with theassumption of optical thin emission give Tex(13CO) much lessthan Tex(C18O) and Tex(12CO), whereas Tex(13CO) is expectedto lie between Tex(12CO) and Tex(C18O). At (0,0) the line ra-tio HCO+(1–0)/H13CO+(1–0) is ∼8, much less than the localinterstellar medium 12C/13C abundance ratio (∼70; Wilson &Matteucci 1992), suggesting that also HCO+(1–0) is opticallythick.

3.1.8. Isotopic ratios

Fig. 9 shows T[13CO(1–0)]/T[C18O(1–0)] and T[13CO(2–1)]/T[C18O(2–1)] for component A along the strip at ∆α=0′′.Along the strip at ∆α = −80′′ we observe values (and lowerlimits) of the order of 5–10 for both transitions. The subcom-ponent of A at ∼ −5 km s−1 shows a steady decrease of lineratios from the northern edge to the ionization front (from 12.5to 7 for J = 1 − 0 and from 8.4 to 3 for J = 2 − 1). Thesame behaviour is observed if ratios are determined from theintegrated emission, although then the decrease is more gentle(e.g. from 9.5 to 7.7 for J = 1 − 0). The two main sourcesof uncertainties arise from the gaussian fits and the temperaturecalibration. We estimate that errors are about 15–20% for (1–0).For the subcomponent at −1 km s−1 the situation is less clear,as the uncertainties are larger (up to 30% for J = 1−0) becausegaussian fits to this component are less accurate. If both 13COand C18O are optically thin, line ratios are equal to abundanceratios; because 13CO is moderately optically thick, a correctionfor optical depth has been applied.

Because of chemical fractionation and selectivephoto-dissociation, chemical models predict an increase ofX(13CO)/X(C18O) at the edges of a cloud (which instead oc-curs only at the northern edge of A). Even though the opticaldepth of 13CO(1–0) obtained from the ratio 13CO(1–0)/12CO(1–0) increases towards the ionization front from 0.2 to 0.3, this canaccount for at most 6% decrease of line ratios. Beam dilutioncannot explain the observed decrease, since C18O is less abun-dant and is likely to be more beam diluted than 13CO at the edge.However saturation of 13CO accounts for the lower values of (2–1) line ratios, since 13CO(2–1) is generally more optically thickthan 13CO(1–0) (and is moderately thick as noted in Sect. 3.1.7).Since at (0,0) the 13CO and C18O emission seem to come fromdifferent regions (a more diffuse one and a high density one),excitation effects can affect line ratios. Also, chemical fraction-

Fig. 9. T(13CO)/T(C18O) ratios both for (1–0) and (2–1) of the twomain subcomponents of A: A1 (Vlsr∼ −5 km s−1, filled squares) andA2 (Vlsr∼ −1 km s−1, open squares), along the strip at ∆α = 0′′ inG353.1+0.6.

ation may be inhibited near the ionization front because of hightemperature and density (Langer et al. 1984).

Assuming X(16O)/X(18O) ∼500 (Wilson & Matteucci1992), and using an optical depth correction and the ratios ofgaussian areas, we find X(12C)/X(13C) ∼ 48 − 57, more orless in agreement with the local interstellar medium value of∼70 (Wilson & Matteucci 1992).

3.2. G 353.2+0.9

3.2.1. Morphology

In Fig. 10 we present the 12CO(1–0) spectra taken towardsthis source. The weak feature at about −40 km s−1 has beendetected at virtually every position in this field as well andwill not be considered further as it is probably unrelated[Vlsr(H109α)=−3.8± 0.8 km s−1 (Wilson et al. 1970)]. Spec-tra of all species, observed along the strips at ∆α=0′′ and∆α = −40′′, are presented in Fig. 11 and 12 respectively. Thebulk of the emission occurs at velocities roughly between −14and +4 km s−1. From Fig. 10 it is evident that the velocity struc-ture of the gas observed towards this region is quite complex:line profiles change considerably from one position to the next,sometimes abruptly (as along the strip at ∆α=0′′), sometimesmore gradually (as along ∆α = −40′′). This is illustrated moreclearly by the diagrams presented in Fig. 13, which show the in-tegrated 12CO(1–0) emission in intervals of 1 km s−1; the cen-tral values of the bins are indicated in each panel. From thesediagrams we see first of all that the emission in the south-easterncorner of the map is blue-shifted with respect to rest of theemission. Comparison with the red ESO/SRC plate (see Fig. 2of Fea90) shows that this “South-Eastern Complex” (S.E.C.) is


Fig. 10. All 12CO(1–0) spectra taken towards G353.2+0.9. Velocity and temperature scale are indicated in the upper right-hand corner.

not associated with the Hii region under consideration, but co-incides with a filament of obscuration, whose edges are outlinedby diffuse emission. The emission towards the remainder of themapped region coincides with the Hii region and is a collec-tion of several different velocity components within the beam.The most relevant of these have been indicated with letters inFig. 13.

The emission integrated between −14 and +4 km s−1 iscompletely dominated by component “B”, with its peak to theNW of the optical Hii region, and the S.E.C. There is a stronggradient in the emission towards the SW quadrant of the mappedregion; the region at ∆α <0′′ and ∆δ < −40′′ is characterizedby emission at a level of<20 K km s−1. At the northern bound-ary of this region the contours run almost parallel to the ioniza-tion front (which coincides with the sharp edge in Hα emission;

see Fig. 7 of Fea90). The early-type stars of the cluster Pis24(Neckel 1984) are located just north of the cavity, in a regionof relatively little molecular emission, and one might suspectthat these stars have blown the region surrounding them clear ofmolecular material. While this may be so, the stars of Pis24 arenot the cause of the Hii region (Fea90) because the ionizationfront is located between the Pis24 cluster and the diffuse Hαemission. Rather, the Hii region is excited by sources locatedinside the nebula (Fea90). These authors concluded that the ob-scuration causing the sharp boundary in Hα must therefore beconnected to the ionization front. In fact, components E, G andH, indicated in Fig. 13, form a chain of lower Tmb cloudletsparallel to the edge of the Hα emission, and lying just southof it. Of these 3 clouds, E and G have the most intense emis-sion, and are presumably also the densest. They are bordering


Fig. 11. Spectra of all observed lines in G353.2+0.9 along the strip at ∆α=0′′. Only spectra at intervals of 40′′ are shown. Transitions andtemperature scale are indicated above each column; velocity ranges below.

the western half of the ionization front, where the Hα emissionhas its sharpest boundary. The weaker component, H, lies alongthe eastern half of the front, where the Hα emission has a morefuzzy edge.

Other features visible in Fig. 13 can account for the obscu-ration seen towards other parts of the Hα emission (see Fig. 7bin Fea90). Component C coincides with the ‘elephant trunk’-shaped obscuration cutting through the Hα emission, and which

houses two compact radio sources (called A and B by Fea90).It has been detected in all lines, including HCO+ and H13CO+,implying it must have a high density (>105 cm−3).

Components B and F cause the obscuration seen towardsthe NW and NE of the Hα emission respectively. They may belocated partly behind the Hii emission, as weak diffuse emissioncan be seen in the lower part of both components (Fig. 7 ofFea90).


Fig. 12. Same as Fig. 11, for ∆α = −40′′.

In summary, contrary to the situation in G353.1+0.6, wheremost of the molecular gas is located beyond the ionization frontas seen from the exciting stars and viewed in an edge-on geom-etry, in G353.2+0.9 most of the molecular emission is locatedeither along the line of sight or north of the optical emission,with only weaker components bordering the southern edge ofthe ionization front. That is, this star forming complex is essen-tially viewed face-on. Components B and F, which are on eitherside of the brightest part of the optical nebulosity (and conse-quently are partly behind it) peak at Vlsr∼ −2 km s−1, so thatwith respect to these clouds the ionized gas atVlsr=−3.8 km s−1

is streaming towards the observer, while component C, whichcoincides with the elephant trunk (and therefore must be in frontof the optical nebula) peaks at Vlsr∼ −5.5 km s−1. This geom-etry implies that the PDR extends over the whole of the mapped

region, i.e. lies in front of the molecular clouds, at least forcomponents B and F, so that we should find “PDR-values” forphysical parameters at all positions, rather than a north-southtrend across the ionization front.

3.2.2. The ionization front

In G353.2+0.9 the emission profile at each position is the com-bination of various components at different velocities. For po-sitions north of the ionization front (i.e. ∆δ >0′′) the dominantemission has central velocities between −6 and −1 km s−1.The emission at each position can be roughly separated intofour velocity ranges: Vlsr< −7 km s−1 (S.E.C), −7 <Vlsr<−3 km s−1 (A and C), −3 <Vlsr< 0 km s−1 (B and F),and Vlsr>0 km s−1 (E), which has been confirmed by first fit-


Fig. 13. Channel maps of 12CO(1–0) emission in G353.2+0.9. The channel width is 1 km s−1, and the central velocity is indicated in each panel.Contour levels are 2(2)20, 25(5)40 K km s−1. The 20 K km s−1 contour is drawn thick.

ting gaussians to the (relatively) optically thin 13CO(1–0) andC18O(1–0) emission and then to the optically thick 12CO(1–0)spectra. At all positions the spectra could be deconvolved in 2 to4 gaussian components with central velocities that are the samewithin 0.5 km s−1 (1 km s−1 in a few cases). The correspond-ing temperatures of the optically thick transitions were derivedby the same procedure as for G353.1+0.6 (see Sect. 3.1.4).

With temperatures determined in this way, we can assess thebehaviour of the molecular emission with respect to the ioniza-tion front along the two strips. The emission of all molecularspecies changes in the same manner (see Figs. 11 and 12) andthe intensities drop very rapidly at the location of the front.At ∆α=0′′, the emission between −7 and −3 km s−1 peaks

at two locations, namely component A (∆δ=160–200′′) and C(∆δ=40′′). Whereas the 12CO(1–0) and 13CO(1–0) lines havetheir primary peak in component A, the transitions more sen-sitive to density, such as HCO+(1–0) and H13CO+(1–0), onlyhave a secondary maximum there, and are stronger in compo-nent C, indicating that C is the denser one. C18O also peaksat C, but only this region was covered for this molecule. Be-tween ∆δ=50′′ and 150′′ components with Vlsr between −7 to−3 km s−1 (i.e. A and C) have lower level emission, while com-ponents B and F have their peaks. The (2–1) intensities showessentially the same features, although 12CO(2–1) is as strongat component C as it is at A. There are slight offsets between thelocations of the peak 12CO(2–1) and 12CO(1–0) emission: the


12CO(2–1) has the same sampling, 20′′, as the 12CO(1–0) witha resolution twice that of 12CO(1–0), and is therefore expectedto better resolve the peak. For the emission between −7 and−3 km s−1 component A seems to peak more to the north in12CO(2–1) than in 12CO(1–0).

The behaviour of the peak temperatures along the strip at∆α = −40′′ is similar to that at ∆α = 0′′, except that the formerdoes not pass through component C, and so there is only onepeak for components with Vlsr between −7 and −3 km s−1.

South of the ionization front (∆δ ∼0′′), all molecules havemuch lower intensities, so it is difficult to follow variations.Component E, which dominates at these locations, has a peakTmb[12CO(1–0)] of only ∼7 K, while along the strip at ∆α=0′′

components A, B, and C reach a peak Tmb[12CO(1–0)] of ∼39,∼31, and ∼25 K, respectively. The lower temperature of com-ponent E could be explained by its small size and therefore byits being more beam-diluted.

As both 12CO transitions are certainly optically thick, theratio R1 should be about 0.8 – 1.0 for the range of Tex wefind here (10–40 K; Levrault 1988). Because the beam size at230 GHz is smaller, we expect R1 to be larger than its realvalue, especially if the gas is clumpy on scale sizes comparableto the 12CO(2–1) beam.R1 has been calculated for each velocitycomponent, at each position where both transitions have beendetected, and lies between 0.5 and 1.4, with a mean value of0.9±0.2. R1 tends to values between 0.6 and 0.9 at high AV(where the effect of beam dilution decreases; see Sect. 3.1.7).This means that on average the cloud is not very clumpy onscales of <∼40′′. However, to explain the lower values of R1

may require a two-layer model with an optically thick innerlayer and a diffuse optically thin warmer front layer which,because of τ21 > τ10, preferentially absorbs the (2–1) emission,as discussed by Pagani et al. (1993) for the case of RCW 34.

3.2.3. Column densities

We determined column densities as for G353.1+0.6. Again, dif-ferent methods give column densities generally within a factorof 2–3, with larger differences at positions with lower emission.The total mass of the molecular gas in our observations has beendetermined using Eq. (2), from Tmb[12CO(1–0)] integrated be-tween Vlsr=−10 or −16 (including the S.E.C.) and 4 km s−1

summed over all positions. With the 1.36 correction for He, thisyields 5700 M�. Adding all Nlte values from gaussian com-ponents yields a total mass of 3500 M�, corrected for He andwith a contribution from positions not observed in 13CO(1–0)estimated as in G353.1+0.6. The mean ratio of WCO columndensity to total LTE column density (determined by adding allcontributions of single components along the line of sight) is2.5±1.8, lower than in G353.1+0.6. The difference may be dueto a different fraction of positions with low-level emission inthe two regions. However, 13CO was observed over a smallerarea than 12CO, and furthermore 13CO was not always detected.LTE masses of single components are reported in Table 3.WCO

masses are roughly twice as much.

Fig. 14. LTE column density of several velocity components inG353.2+0.9 along the two main strips. Components A and C are in-dicated by filled squares; B by open squares; E by triangles. ∆δ=0′′

coincides with the position of the ionization front.

Table 3. Physical parameters of molecular clouds in G353.2+0.9. Adistance of 1.7 kpc has been assumed. All LTE masses are correctedfor He. All densities are derived from H2 mass divided by volume

Component Central Size Mass Mass MeanVelocity (2R) (LTE) (Virial) Density

(km s−1) (pc) (M�) (M�) (cm−3)

A ∼ −6 1.8 1500 1500 7× 103

B ∼ −1 0.9 300 800 104

C ∼ −5 0.6 75 400 104

E ∼ 2 0.7 50a 500 3× 103(a)F ∼ −3 0.9 350 600 9× 103

G ∼ −4 0.9 90a 390 3× 103(a)H ∼ −2 0.8 25a 420 103(a)

a Mass and density calculated from WCO column densities.

Fig. 14 shows Nlte along the strips at ∆α=0′′ and ∆α =−40′′ for components A, B, C and E. With the exception of E,peakNlte’s are∼ 1022 cm−2, similar to those obtained for com-ponent A of G353.1+0.6, although line profiles are deconvolvedin finer detail there. Component E, which dominates south ofthe ionization front, has a peakNlte of∼ 5.5×1020 cm−2, witha relatively large uncertainty because 13CO is very weak.NWCO

is much larger, at 4 × 1021 cm−2. It is evident from Fig. 14that Nlte rapidly decreases towards the ionization front and hasa minimum value south of the ionization front, confirming thefact that there is a relative scarcity of molecular gas there.

Sizes and mean densities, reported in Table 3, were deter-mined in the same way as for G353.1+0.6. Densities are gener-


ally of the order of 103–104 cm−3; the density of componentsE, G and H has been determined from theWCO column density.For comparison with the other values, note that the maximumNlte’s of E, G and H divided by 2R are only∼ 200 cm−3. Virialmasses (assuming each cloud is a homogeneous sphere) are alsogiven in Table 3; for components C, E and H these are muchhigher than their LTE or WCO-masses, indicating that these arenot in virial equilibrium. However, since it is not clear if the linewidth is due to turbulence or to ordered motions, it is difficultto say whether the clouds may be collapsing or expanding.

3.2.4. Temperatures and relative abundances

As for G353.1+0.6, we estimated Tkin from Tex[12CO(1–0)]. Atleast along the two main strips,Tkin rapidly decreases both southand north of the peaks. Line ratios suggest that 13CO and C18Oexcitation temperatures are lower than Tex(12CO), at least in thenorthern part of the field. Components A and B have a peakTex(12CO) ∼40 K, while C reaches ∼30 K, and components Eand G have Tex(12CO) ∼10–16 K. Since C, E and G are verysmall features, their Tmb[12CO(1–0)] may be underestimatedbecause of beam dilution. Thus, as for G353.1+0.6, we haveadopted a constantTkin∼ 50 K to use as input for the LVG modelof Goldsmith et al. (1983). The lowR1 ratios and high 12CO(1–0) brightness temperatures of components A, B and F requiresdensities < 102–103 cm−3 and abundances X(12CO)> 10−4

to be explained. But 13CO and C18O yield densities ∼ 103–104 cm−3 and abundances X(13CO)∼ 10−6 and X(C18O)∼10−7, so a low density absorbing layer that reduces R1 mayexist. Component C has densities ∼ 103 cm−3 (for C18O, 104

cm−3) and abundancesX(12CO)∼ 10−5,X(13CO)∼ 10−6 andX(C18O)∼ 10−7. Both E and H have densities∼ 102 cm−3 andX(12CO)∼ 10−4–10−5; their emission in 13CO and C18O hasnot been detected. Finally, component G has a density ∼ 103

cm−3 and abundances X(12CO)∼ 10−6–10−5 and X(13CO)∼10−7. LVG column densities are generally lower (but within afactor of 2–3) than the corresponding Nlte.

3.2.5. Optical depths

As we have shown in Sect. 3.2.2, line ratios R1 are generally<1 with many values ≤0.7. This indicates either low excita-tion temperatures or low opacities (Levreault 1988). However,R2 is <∼10 almost everywhere, implying that 12CO is opticallythick. Along the two main strips, T[13CO(2–1)]/T[13CO(1–0)]is generally between 1 and 2 for A, B and C, suggesting thatat least 13CO(2–1) may be moderately thick in some locations(and consequently excitation temperatures derived from line ra-tios may be in error at those positions). C18O line ratios, whereavailable, are indicative of low opacities. Finally, T[HCO+(1–0)]/T[H13CO+(1–0)] is 5 and 8 for A and C at the locations oftheir maxima, indicating that HCO+(1–0) is optically thick. Atother locations, only lower limits are available; these range from2 to 5.

3.2.6. Isotopic ratios

Both for the (1–0) and for the (2–1) transition the ratioT(13CO)/T(C18O) has a large scatter around a mean value of12.8 for (1–0) and 10 for (2–1) and is generally <∼ 20. No cleartrend with position is evident, although there is a marginal in-dication of an increase of the line ratios at the cloud edges,as predicted by many PDR models (e.g. Minchin et al. 1995).The mean values are greater than the terrestrial isotopic ratioof 5.5 (Taylor & Dickman 1989). Correcting for 13CO opacityand under the same assumptions as used for G353.1+0.6 (seeSect. 3.1.8), we find X(12C)/X(13C) between 22 and 61, with52 near the A-peak and 61 at the C-peak. Unfortunately, all com-ponents south of the ionization front are undetected in C18O, soonly scarcely significant lower limits (T(13CO) /T(C18O)≥ 1–4)are available there.

3.2.7. Clouds at the ionization front

As mentioned in Sect. 3.2.1 the molecular clouds labeled E, G,and H (see Fig. 13) form the material into which the ioniza-tion front is proceeding. They are in the shape of a thin strip ofgas south of the ionization front. In many respects this situationis reminiscent of the bright bar in the Orion Nebula, which isbounded by a thin strip of molecular gas on the side opposite tothat of the exciting stars of the Trapezium cluster. Also in Orionthe CO emission from this strip is much weaker than that of themolecular clouds associated with the BN/KL region (see Wil-son & Mauersberger 1991). The G353.2+0.9 complex presentsmany similarities with the Orion complex, in the sense that wesee the Hii region produced by the new-born massive stars atthe edge of a large molecular complex where star formationmay just begin and with a similar orientation with respect to theobserver. Perhaps G353.2+0.9 is in an earlier stage, since therethe equivalent of the Trapezium stars are not visible becausethe surroundings have not yet been sufficiently cleared by thestellar ionizing radiation.

An indication of the properties of the gas in this strip hasbeen obtained by using the models published by Koster et al.(1994), who consider the CO spectra emerging from clumpymolecular clouds in PDR’s. They present models for two cases:clouds that are illuminated from one side (their model B) andthose that are illuminated on both sides (model A). The visualextinction through the clouds and the illumination (by far UV[FUV] radiation) are both varied, and for each case the ther-mal and chemical structure of the clouds are calculated. Theypresent their results in a series of figures, showing the change ofthe brightness temperature at the center of the emergent 12COand 13CO profiles as a function of the upper level of the ro-tational transition, for various combinations of cloud density,AV , velocity dispersion, and intensity of the incident FUV ra-diation. For G353.2+0.9 their model B is the more appropriate,with the clouds along the ionization front being illuminated bythe embedded sources responsible for the excitation of the Hiiregion. As indicated by Fea90, the ionizing sources embeddedin G353.2+0.9 are the equivalent of 5 O9V stars, which implies


that the FUV flux density (6 < hν < 13.6 eV) is <∼ 105 timesthe average interstellar field. In fact, 5 O9 V stars yield 1.2×1039

erg s−1 (Panagia 1973). Assuming the Hii region is 1.2 pc indiameter and the stars lie at its center yields a FUV flux density<∼ 29 erg cm−2 s−1 (the inequality comes from the photons withhν < 6 eV). The average interstellar FUV flux density is ∼1.610−3 erg s−1 cm2 (Tielens & Hollenbach 1985), then the localFUV flux density is <∼ 1.8× 104 times greater. However, sincethe ionizing sources are located along the southern edge of theHii region (Fea90) at a projected distance of at most 10′′ fromthe radio border, the FUV flux density provided by even a sin-gle O9 V star on a cloud surface 0.1 pc away from it may be ashigh as 105 times the average interstellar flux density. Becausethe exact geometry and the amount of FUV radiation being ab-sorbed by matter surrounding the stars are unknown, this mustbe considered an upper limit.

As can be seen in Fig. 14, component E has Nlte ∼ 3 to6×1020 cm−2, corresponding to AV ≤2 mag. We find that forcloud E, G and H the best agreement between model and ob-servations is found for FUV field G = 103 times the averageinterstellar value, density 104 cm−3, AV = 2 mag and velocitydispersion 1 km s−1. According to the model of Koster et al.(1994), the 12CO(1–0) emission has τ < 1 in this case.

The parameters that best fit the observations of componentC, the narrow “elephant trunk” of obscuration which houses acompact radio component (Fea90), are G ≥ 103, density of theorder of 105−6 cm−3, AV ∼ 10 mag, and velocity dispersion 3km s−1.

4. Discussion

4.1. Morphology

Both in G353.1+0.6 and in G353.2+0.9 the molecular line pro-files result from the superposition of the emission of differ-ent components with scale sizes from 1 to 0.5 pc, the lowervalue probably being set by the instrumental resolution. TheHII masses, as given by Fea90, are only a small fraction ofthe total molecular masses observed in the two fields, at mostof the order of few percent, although in G353.2+0.9 they arecomparable to the fraction of mass of each of the small cloudswhich are lined along the southern front. Also the HII densities(∼ 100 cm−3), derived from the high resolution VLA obser-vations of Fea90, are lower than the H2 densities (∼ 103–104

cm−3), meaning either that the ionized gas has simply expandedor/and that the Hii regions are very clumpy and the ionized gashas a volume filling factor� 1.

The morphology of G353.1+0.6 is that of an evolved starforming region viewed edge-on, with stars/Hii region/molecularcloud aligned in a south-north direction. Low resolution FIR(McBreen et al. 1983) and IRAS (Persi et al. 1986) maps closelyresemble low-resolution radio-continuum maps (e.g. Schraml& Mezger 1969), indicating that the IR peak lies south of themolecular cloud and suggesting that the source of heating forthe dust and the cloud is the radiation from the stellar cluster,of which N49 is the brightest star. The fragmentation of the

cloud is also revealed by the structured and extended form ofthe ionization front, where Fea90 have found radio emissionfrom ionized surface layers of small molecular blobs. Theseblobs must have densities of the same order as that of the parentcloud (∼ 104 cm−3). The overall morphology is reminescentof a very late stage of a blister-type configuration and indicatesthat star formation has occurred at the edge of the molecularcloud. No traces of recent star formation were found close tothe ionization front or within the molecular cloud (Fea90).

The geometry of G353.2+0.9 is rather different from whatwas expected on the basis of previously available data (see e.g.Fea90). In fact, only a weak “bar” of molecular gas was foundto the south of the sharp ionization front observed in Hα andin the radio-continuum, whereas the majority of the molecularemission comes from behind the ionized gas and to the north ofit. In this case, then, a late stage blister configuration is viewedface-on. The molecular components associated with the ionizedgas (C, E, G, H) are just a small part of a much larger molecularcomplex. The low resolution IRAS maps (Persi et al. 1986)seem to indicate that the IR peak is located near the molecularpeak, north of the ionization front. However, a high resolutionmap at 10 µm (Frogel & Persson 1974) shows an elongatedresolved structure coinciding with the ionization front. Just likeG353.1+0.6, G353.2+0.9 is a blister type Hii region, but seenface-on, and with active star formation.

4.2. Molecular gas kinematics

Velocity dispersions can be assessed for each molecular com-ponent by averaging the 13CO(1–0) gaussian widths (FWHM)over small areas around the peaks. In both G353.1+0.6 andG353.2+0.9 the average FWHM’s are large (≥ 2 km s−1), andvirial masses are systematically greater than LTE masses, atleast for the smallest components. In G353.1+0.6 the cloudswith larger dispersions are the two main subcomponents of A(FWHM∼ 4 km s−1), north of the ionization front, whereasA, B and F in G353.2+0.9 have lower average line widths(FWHM∼ 3 km s−1) and seem closer to virial equilibrium (seeTables 2 and 3). These results may be due to the different agesof the associated Hii regions and/or to the different geometrywith respect to the exciting stars. The components closer to theionization front on the plane of the sky (C and E in G353.1+0.6and C, E, G and H in G353.2+0.9) have similar dispersions (∼2–3 km s−1) and in G353.2+0.9 their virial masses are muchgreater than their LTE masses. In both regions the correspond-ing 12CO(1–0) average gaussian FWHM’s are 0.5–1 km s−1

larger than those of 13CO(1–0); for C18O, HCO+ and H13CO+

only few data are available, but their dispersions seem generallyof the order of those of 13CO(1–0).

The interaction between Hii region and molecular cloud af-fects the kinematics of the gas, as indicated by the appearanceof red-shifted and blue-shifted molecular emissions towards theionization fronts. In G353.1+0.6 cloud A is composed of twomain subcomponents at ∼ −5 and ∼ −1 km s−1 and dom-inates the field north of the radio-continuum emission. Com-ponent E (Vlsr∼ −6 to −8 km s−1) must lie in front of this


emission, and is blue-shifted with respect to A, whereas com-ponent C (Vlsr∼ 3 km s−1), which is red-shifted with respectto A, could be located behind the ionization front with respectto the observer. This suggests that the ionized gas is blowing a“bubble” of molecular matter. Rough calculations based on the“rocket effect” (see e.g. Spitzer 1978) indicate that, assuminga systemic velocity equal to Vlsr(H109α), molecular fragmentssurrounding the ionization front might be accelerated to the Vlsr

of C and E in a time <∼ 8 × 105 yrs, comparable to the main-sequence lifetime of an O star. In G353.2+0.9 the situation ismore complex. North of the Hii region emission at two dis-tinct velocities of about −6 and −2 km s−1 are clearly visible(see Fig. 10); these velocity fields must have been pre-existing,since they are located far from the radio-continuum emission,and indicate the presence of two distinct clouds. In this casethe velocity difference may originate from a larger scale pre-existing velocity pattern. However, around the ionization frontthe associated components show a large range in velocity. Infact, the molecular emission associated with the elephant trunk,C, has a Vlsr∼ −5 km s−1, whereas that located south of it, E,has a Vlsr∼ 2 km s−1. Components G and H have Vlsr∼ −4 and∼ −2 km s−1, respectively, and lie west and east of E. SinceC and E are almost symmetrical with respect to IRS 4, oneof the possible ionizing sources (Fea90), and to its associatedcompact radio source, we can figure that, also in this case, theionized gas is blowing a bubble of molecular matter, where C isfacing the observer and E is on the opposite side. ComponentsE, G and H may then be forming an expanding ring south ofIRS 4. Again, calculations based on the rocket effect indicatethat this velocity configuration might be reached in <∼ 2 × 105

yrs. Thus, the correspondence between the radial velocities of Aand C would be coincidental and star formation might have oc-curred preferentially in the molecular clouds B, C and F, ratherthan in A.

4.3. Molecular and ionized gas

We can estimate an age for the obscured ionized region inG353.1+0.6 by determining the rate of mass erosion of A, fol-lowing the model of Spitzer (1978). We assume that the ionizingsource is N49 and use the physical parameters given by Panagia(1973) for an O5 V star. Then we obtain dM/dt ∼ −4× 10−4

M�yr−1 which, using the total HII mass estimate (81 M�) ofFea90, yields 2 × 105 yrs for the time needed to ionize thatamount of gas. This time is less than the main-sequence life-time of an O star. On the other hand, the lifetime of componentA against photo-evaporation is ∼ 3 × 106 yrs, of the order ofor slightly greater than the main-sequence lifetime of an O star.Since we do not know the exact geometrical relationship be-tween exciting stars and molecular clouds in G353.2+0.9, wecannot perform the same kind of calculation, although we mightexpect similar results. Thus, a considerable fraction of moleculargas may survive the erosion from UV radiation of the associatedOB stars.

A reasonable evolutionary scenario for the Hii regionG353.1+0.6 can be constructed by trying to relate the original

molecular gas between N49 and the present molecular bound-ary. If we assume that N49 formed at the very edge of a sphericalcloud with a density of ∼ 104 cm−3 and a diameter of 0.7 pc(the distance of the star to the ionization front) the mass of thatcloud was ∼ 100 M�. The ionization of the gas might haveoriginated the present optical nebula (with an optical diameterof at least 3 pc) in < 3 × 105 yrs if the Hii gas expands at thesound velocity (∼ 10 km s−1). This time is comparable with thelifetime of the ionized gas given above. Interestingly, the densityin this gas after the expansion should be ∼ 102 cm−3, similarto the upper limit of the electron density estimated by Fea90(60 cm−3). Since the emission measure decreases with the fifthpower of the expansion factor, this diffuse gas becomes unde-tectable in the radio-continuum. For the more compact struc-ture of G353.2+0.9, its higher mean electron density and radio-continuum brightness temperature suggest that in this case theionized gas is still close to the parent clouds and has not yetexpanded significantly.

4.4. Effects of UV radiation

If we assume a distance between N49 and cloud A surface of 1pc and use the physical parameters given by Panagia (1973) foran O5 V star, we get a FUV flux density which is a factor 7×103

greater than the average interstellar flux density. The 100/60 µmcolour temperature has been evaluated for both regions by Persiet al. (1986) and is 43 K towards G353.1+0.6 and 45 K towardsG353.2+0.9. Hollenbach et al. (1991) have studied a PDR modelwhich gives the 100/60 µm colour temperature versus the FUVflux density (see their Fig. 14): a colour temperature of 45 Kcorresponds to a FUV flux density of ∼ 3 × 103 times the av-erage interstellar flux density, which agrees with our estimatein G353.1+0.6. For G353.2+0.9 (see Sect. 3.2.7) we have ob-tained an upper limit of 105 times the average interstellar fluxdensity on the surface of the small clouds located south of theionization front. The IRAS maps of the latter region by Persiet al. (1986) suggest that the far-infrared peak is close to themolecular peak, well north of the ionization front, where theFUV flux density might well be ∼ 103 times the average inter-stellar flux density, in agreement with the model of Hollenbachet al. (1991). The extended 10 µm structure found by Frogel& Persson (1974), however, indicates the presence of hot dustclose to the ionization front, but obviously this feature has beenlost in lower resolution FIR observations. Thus, since in bothregions 100/60 µm colour temperatures, 12CO(1–0) peak tem-peratures and mean densities of the molecular gas are similar,the 12CO(1–0) brightness temperature of the bulk of moleculargas seems unaffected by the age of the bordering Hii region andis related only to the FUV flux density and to the gas density.

4.5. Heating of the molecular clouds

There are other indications that the heating model assumed forboth regions, i.e. UV radiation from sources at the edge of thecloud, is correct. In G353.1+0.6 the excitation temperature ofcloud A derived from C18O line ratios peaks at (0,0) and then


suddenly decreases when moving northward, away from theionization front. Although rather uncertain, the excitation tem-perature seems to change (at least for the subcomponent at∼ −5km s−1) from > 40 K at the peak to ∼ 20 K towards the pro-jected cloud center. Along the strip at ∆α = 0′′, the 12CO(1–0)excitation temperature peaks at ∆α = 40′′ and, with the ex-ception of the (0,0) position, is some 5–10 K higher than theC18O excitation temperature. This clearly indicates that cloudA is colder inside and that the surface layer bordering the Hiiregion is warmer than the more distant side. Then, the heatingsource is external and located south of the cloud. The same, al-though less clearly, is suggested by the excitation temperaturesfor clouds A and B in G353.2+0.9, that also seem warmer to-wards the Hii region. The existence of a temperature gradientis also indicated in Fig. 15, which shows the Tex[12CO(1–0)]and the Tmb[13CO(1–0)] versus the extinction AV (determinedfrom Eq. 3) in both regions. All available data can be approxi-mately fitted by a T [1−exp(−cAV )] relation but with differentT . This implies that each line samples an almost constant ex-citation temperature region, but the 12CO(1–0) line is sensitiveto a warmer envelope while the 13CO(1–0) line traces an inner,cooler part of the molecular cloud. The dependence of bright-ness temperature on AV would then be mainly determined bythe opacity rather than by Tex differences. The 12CO(1–0) datacan be fitted by T ∼ 40 and c ∼ 0.1, whereas, assuming LTE,Tex= 40 K, a line width = 3 km s−1, theN (H2)/N (13CO) ratio ofDickman (1978) and an isotopic ratio X(12CO)/X(13CO)= 60,a theoretical value c ∼ 0.5 is obtained. This may be due to aresidual 12CO(1–0) beam dilution which lowers Tmb at cloudedges (i.e. at low to moderate AV locations), thus deformingthe curve before it saturates. Instead, the 13CO(1–0) points canbe fitted by T ∼ 8 K (G353.1+0.6) and ∼ 10 K (G353.2+0.9),and c = 0.08, which agrees quite well with the theoretical valuefor Tex= 15 K and a line width = 2 km s−1. If there is a Tex

gradient in both regions, then column densities calculated withTex’s derived fromTmb[12CO(1–0)] must be revised. But we findthe error is at most∼ 50 %, and this certainly would not changethe results of the previous sections.

4.6. Line ratios

The clearest difference between the two regions is in the R1

ratios, which are > 1 in G353.1+0.6 (Sect. 3.1.7) and < 1 inG353.2+0.9 (Sect. 3.2.2). Because the latter values pose a fewproblems for their interpretation, e.g. by an LVG model, wehave suggested the existence of a warmer and low density in-terclump medium affecting the 12CO ratios in G353.9+0.2, asstudied by Pagani et al. (1993) for the case of RCW 34. How-ever, it is not clear whether the difference in line ratios betweenthe two regions is due to geometrical effects (face-on versusedge-on structure) or to the different ages of the nebulae (e.g.,after some time the interclump medium may be photoevapo-rated by the UV radiation), or both. Towards the S.E.C. lineratios are ∼ 1, i.e. “normal”. On the other hand, 12CO(1–0)self-absorption by cooler clumps with small velocity dispersionhas been found in G353.1+0.6; this may be due to the exciting

Fig. 15. a 12CO(1–0) excitation temperature versus extinction inG353.2+0.9. The solid line shows T [1 − exp(−cAV )], with T = 40K and c = 0.1. b Same as in a, for G353.1+0.6. c 13CO(1–0) mainbeam temperature versus extinction in G353.2+0.9; solid line as in a,but with T = 10 K and c = 0.08. d Same as in c, for G353.1+0.6, butT = 8 K and c = 0.08.

stars/cloud geometry, since molecular clumps not facing the Hiiregion directly but still in front of the observer may not be effi-ciently heated by UV radiation. Being cooler, the line of sightcomponents will appear as absorption dips.

5. Conclusions

We have presented and analyzed line observations of molec-ular clouds associated with the two Hii regions G353.1+0.6and G353.2+0.9, which are part of the star forming complexNGC 6357. The main results can be summarized as follows:

1. The molecular emission towards G353.1+0.6 arises northof the optical nebula, where an ionization front is eroding amolecular cloud of∼ 1 pc of diameter which we have calledcomponent A. An elongated structure, component E, runsalong the ionization front and is blue-shifted with respectto A. This suggests that this is placed between the observerand the ionization front and causes its obscuration. Towardsthe ionization front there is also red-shifted molecular ma-terial (Vlsr∼ 3 km s−1), which could be behind it. We haveinterpreted these features as an expanding bubble poweredby the OB stars in the optical nebula.

2. A cooler, clumpy molecular layer external to componentA may account for a 12CO(1-0) self-absorption feature de-tected at ∼ −0.7 km s−1.

3. Gaussian fits to the line profiles show that component A iscomposed of at least two main subcomponents at∼ −5 and∼ −1 km s−1. Typical H2 densities, derived from Nlte andan LVG model, are 103–104 cm−3, and there is evidence ofa density increase near the ionization front. A mass estimatefor component A is ∼ 1300 M�, and its maximum Nlte is∼ 1022 cm−2.

4. A kinetic temperature of ∼ 40–50 K is obtained from ourobservations, although there are clear indications of the ex-istence of temperature gradients in component A. This cloudis not only colder in the inner parts, but the C18O excitation


temperature peaks near the ionization front, confirming thatthe heating sources (the cluster of OB stars) are external andlocated south of the cloud.

5. We estimated that the values for CO abundances aretypical (although 12CO seems slightly underabundant) inG353.1+0.6. A decrease of the 13CO/C18O isotopic ratio to-wards the ionization front was detected across componentA.

6. The morphology of G353.2+0.9 is rather different fromwhat was expected from the previously available indica-tions (Fea90). In fact, only a weak “bar” of molecular gaswas found to the south of the sharp ionization front, whereasthe majority of the molecular emission comes from a regionbehind or to the north of the Hii region. This is composed oftwo main features with Vlsr∼ −6 and∼ −2 km s−1, whoseorigins should be investigated by observations on a largerscale. The newly-born stars within the nebula, which arealso its exciting sources, are more probably related to thecloud at −2 km s−1; if so, the Vlsr(H109α), ∼ −4 km s−1,suggests that the ionized gas is slowly expanding towardsthe observer. Thus, in this case we are clearly viewing anionization front face-on.

7. The molecular fragments directly associated to G353.2+0.9are component C, which coincides with the elephant trunkvisible in the Hα image (Fea90), and components E, G andH, small clouds that border the southern edge of the neb-ula. These have very different Vlsr’s, in the range from −5to 2 km s−1, suggesting they are located on the surface ofan expanding ionized bubble powered by IRS 4, which isprobably the main source of energy in the nebula (Fea90).Since component C is visible in all observed lines, includingHCO+(1–0) and H13CO+(1–0), it must be very dense (> 105

cm−3). It houses two compact radio sources and IRS 4 islocated to the south, at the apex of the elephant trunk. Onthe other hand, components E, G and H are relatively weak,withTmb[12CO(1–0)] in the range 5–13 K, and their sizes are< 1 pc. Density estimates vary between 102 and 104 cm−3,depending on the adopted model, and masses are≤ 90 M�.

8. The other components surrounding G353.2+0.9, i.e. A, Band F, have densities ∼ 103–104 cm−3, Tmb[12CO(1–0)]> 30 K and LTE masses between 300 and 1500 M�. In thiscase, too, a Tkin∼ 40–50 K seems appropriate, and there aresome indications of colder inner regions. Maximum Nlte’sare∼ 1022 cm−2 and we found roughly the same abundancesfor optically thin CO isotopes as in G353.1+0.6, while 12COabundances range from 10−6 to > 10−4.

9. As shown by Fea90, the Pis24 cluster appears unrelatedto G353.2+0.9, and it is located inside a large cavity withrelatively little molecular gas. We can speculate that theinteraction between the cluster and the gas originated thiscavity much before the formation of the Hii region and hasnot left other traces beside the large molecular hole.

10. G353.1+0.6 is an evolved Hii region with an age of∼ 10−5

yrs, powered by a cluster of OB stars within the opticalnebula. The progenitor molecular cloud which originatedthese stars was photoevaporated by the UV radiation and ex-

panded to form the low density sphere of ionized gas aroundthem. G353.2+0.9 appears instead as a younger, more com-pact Hii region.

11. The main difference between the two regions is in the R1

ratios, which are> 1 in G353.1+0.6 and< 1 in G353.2+0.9.In the latter case, we have proposed the existence of a diffusewarmer and low density interclump gas which absorbs the12CO(2–1) emission more than the 12CO(1–0) emission.

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