6. Discussion

In the previous section, a detailed description of the observations and analysis for each individual species has been given. We discuss here the results in general, both for the physical structure and the chemical state of the three sources.

6.1. Physical structure

In Paper I, the physical parameters were derived primarily from the observed excitation under the assumption that a single physical component is present within the 15'' beam (0.16 pc). We have analyzed here the excitation of many more species and can compare the results with those found from . At the same time, it provides information on the question of multiple components within the beam. In this respect, it is useful to compare the observed scale in W 3 with that of structures found in the more nearby Orion giant molecular cloud, which is at 450 rather than 2300 pc. From observations of the extended Orion ridge component, it appears that variations in physical and chemical parameters occur on scales of a few arcmin, corresponding to (Ungerechts et al. 1995; Schilke et al. 1992). Thus, for the surrounding, more quiescent molecular cloud material, the linear scales probed in W 3 are comparable and the assumption of a single component may be reasonable. However, in the Orion/KL region around the massive young stellar object IRc2, significant changes occur on scales of < 10'', corresponding to . Structures of similar linear size in the W 3 region will be strongly diluted in the 15'' beam.

W 3 IRS4

For IRS4, the deviations from the picture of a single extended, quiescent warm cloud are the smallest. The cloud seems to be well described by a density and , as derived from . The excitation of other species such as CS, CN, CCH and HNC results in similar parameters. However, there is some evidence that warmer material is present in the beam. Weak evidence comes from the fact that the CO main beam brightness temperatures and the line ratios indicate temperatures of order . The strongest case is formed by the rotational temperature . Sulfur-bearing molecules have also been found in other regions to be associated with elevated temperatures and shocks (Ziurys 1991).

There are several ways to create elevated temperatures in the IRS4 region. The most likely are through shocks or through heating by ultraviolet radiation. In star-forming regions shocks are often associated with outflowing material, but IRS4 does not show any clear sign of an outflow in its line profiles (see Sect. 5.1). Recently, ``hyper-compact'' objects (e.g. Claussen et al. 1994) have been discovered in giant molecular clouds. These are small (<1'') clumps of ionized material seen in VLA measurements. The ionization presumably originates from embedded B-stars, but is kept to small scales due to very high densities close to the star. Therefore it will only have an impact on very small scales within the molecular cloud. In Paper I, we postulated the picture that IRS4 is located on the far side of the molecular cloud and has already broken free on one side, illuminating the cloud from behind. In this scenario, the ultraviolet radiation from the newly formed star can heat the gas and set up a photon-dominated region (PDR) on the back side. In addition, it will drive an ionization front into the dense material. At the interface between the expanding ionization front and the molecular cloud, a shock will be present, which can also heat the gas and return grain mantle material. The shocked gas can be quite extended, since the H II region itself already has a size of several arcsec (Colley 1980). The PDR will extend over similar scales or larger. This interpretation can be distinguished from the hyper-compacts scenario through interferometer observations of species like .

W 3 IRS5

In the case of IRS5, more than one physical component is certainly present in the beam, since it is well known that the source has a strong outflow (see HMMT 1994 for the most recent results). In addition, the extremely deep self-absorption in the CO profiles shows that there must be a temperature and density gradient or stratification. The depth of the absorption in the lower-J lines indicates that low temperature () and/or low density gas is present on the outside. The fact that the 6-5 line still shows self-absorption (HMMT 1994), but not the 9-8 line (Boreiko & Betz 1991), suggests that a range of conditions is present. The latter authors give an upper limit on the temperature of the absorbing gas of 80 K, whereas HMMT find an excitation temperature of 26 K. From the smaller self-absorption in the line and its absence in the lines, the column density of the absorbing gas is limited to . Since this is a factor of 6 smaller than the column density of the emitting gas and since our submillimeter observations are biased toward the warmer and denser gas, our results are not expected to refer to the cold component.

Although gradients are likely to be present, the excitation analysis shows that most molecules (including CS, , HCN) are located in fairly extended gas with and . Part of this agreement stems from the fact that most lines in the 230 and 345 GHz windows have critical densities of order a few , so that it is no surprise that most detected species give a similar result. Observations of higher excitation lines are needed to determine whether the density is fairly homogeneous, or whether a gradient or power law distribution is present, as would be expected for a collapsing cloud core.

A physically distinct region appears to be probed by the sulfur-bearing molecules, at least by those containing oxygen. In Paper I, it was already found that the emission comes from a warmer region, which was tentatively connected with the outflow. The rotational temperature has been better constrained in this work, but the same conclusion remains (see Sect. 5.5.3). In addition, it was found from the strengths of the optically thick lines of SO and that the source size is . Just as in IRS4, the chemical composition of this material seems disjunct from that of the bulk of the gas. In Orion IRc2, turned out to be a major coolant of the gas in this frequency range (Groesbeck 1994). This is not the case for IRS5. Not only is the column density much lower than that toward IRc2, but the line width is smaller as well. A crude estimate shows that the total emission in the 345 GHz range is less than half of that of CO, putting IRS5 somewhere in between IRc2 and the Orion-S source in this respect (Groesbeck 1994). may be formed from the large amount of in the same gas.

The region probed by SiO appears to be even denser with . Similarly high densities and even higher temperatures have been found for SiO toward Orion IRc2, which has brightness temperatures in interferometer observations of (Wright et al. 1995). This suggests disruption of grains by strong shocks and rapid formation of SiO in the warm post-shock gas. Note, however, that in contrast with Orion/KL, the SiO and lines are not noticeably broader toward IRS5 than those of other species. Also, the SiO is not kinematically displaced in contrast with L 1448 (Guilloteau et al. 1992) and NGC 1333 IRAS4 (Blake et al. 1995). Interferometric observations of both SiO and are needed to properly investigate these hot and dense regions.

The fact that no rotational lines within the first vibrational levels of HCN, CS and/or CCH have been detected toward IRS5 implies that the hot and dense gas does not contain large quantities of these molecules.

W 3()

At least two different components are observed toward W 3() within the 15'' JCMT beam. For many species, a significant fraction of the emission comes from the warm and dense () core (>20'') in which both W 3() and W 3(OH) are located.

The second component is the more compact ``hot core'' type region, which is warmer () and denser () than the surrounding envelope. The saturated organic molecules such as (Wink et al. 1994), part of the and , and probably most of the dimethyl ether and methyl formate are found here as well, based on their rotational temperatures. The unresolved continuum source of <1 '' (Wilner et al. 1995) is too diluted in our beam to result in detectable emission. In addition, at submillimeter wavelengths the continuum is optically thick, so that the line emission must originate from a larger region. This strong continuum radiation can, however, affect the excitation of molecules such as HNCO and HDO (Helmich et al. 1996), and lead to large populations in vibrationally excited levels of molecules such as HCN and CS.

The , SO and SiO lines again indicate high temperature and densities. Most likely, these species are located in the outflow traced by the masers (Alcolea et al. 1992). This region probably contains the found by Phillips et al. (1992).

6.2. Chemical characteristics

Surveys like the one presented here are ideal tools to establish the beam-averaged composition of the gas, since the excitation effects can largely be taken out. As discussed in Sect. 4.3, it is more difficult to make the final step toward abundances, which can be directly compared with models. The beam-averaged column densities are obtained from , but this refers to the cloud as a whole whereas some molecules may be more localized. The excitation analysis only gives an estimate of the local number density, not of its column density. For some species such as in IRS5 or in W 3(), the local abundances may be more than two orders of magnitude higher than the beam averaged values listed in Table 5 (click here) (see Sect. 5.4.3). These abundances should therefore be interpreted with care.

 

W 3 W 3 W 3 Extended Compact

IRS4 IRS5

0 . 5 0 . 5 0 . 5 0 . 5 0 . 5

CO 1 . 35(-4) 1 . 35(-4) 1 . 35(-4) 1 . 35(-4) 1 . 35(-4)

CN 2 . 4(-9) 6 . 9(-10) 5 . 7(-10) 1 . 0(-8) 1 . 9(-9)

HCN 6 . 0(-9) 2 . 0(-9) 6 . 3(-9)

HNC 1 . 4(-9) 7 . 3(-10) 4 . 7(-10)

<1 . 3(-10) <3 . 8(-11) 7 . 8(-11) 1 . 5(-8)

HNCO <2 . 0(-10) 1 . 7(-10) 2 . 5(-9)

CCH 1 . 4(-9) 4 . 6(-10) 6 . 8(-10)

5 . 5(-9) <3 . 9(-10) 3 . 2(-9) 2 . 3(-8) 3 . 1(-9)

5 . 7(-10) 3 . 1(-10) 2 . 1(-9) <2 . 6(-8) 9 . 1(-8)

3 . 0(-9) 6 . 2(-10) 4 . 6(-8) 1 . 9(-8) 9 . 4(-7)

<1 . 4(-10) <7 . 2(-11) 1 . 4(-10) 1 . 2(-8)

- - 1 . 0(-8) <5 . 6(-9) 4 . 7(-8)

- - 3 . 5(-9) <1 . 1(-8) 7 . 8(-8)

CS 2 . 9(-9) 7 . 7(-10) 5 . 2(-9) 3 . 3(-8) 2 . 5(-8)

SO 1 . 0(-9) 1 . 9(-8) 6 . 8(-9) 1 . 1(-8) 6 . 3(-7)

1 . 1(-9) 2 . 0(-8) 5 . 2(-9) 4 . 6(-9) 4 . 1(-7)

3 . 1(-10) 5 . 0(-11) 7 . 8(-10) 3 . 4(-9) 3 . 1(-9)

7 . 9(-10) 7 . 7(-10) 5 . 2(-9)

OCS 3 . 6(-10) 3 . 4(-10) 1 . 2(-9) <1 . 0(-8) 6 . 9(-9)

SiO 2 . 6(-11) 2 . 0(-11) 3 . 1(-10) <1 . 5(-9) 1 . 9(-8)

2 . 7(-9) 7 . 3(-10) 1 . 2(-9) 1 . 1(-8) 2 . 4(-9)

<1 . 4(-10) 2 . 4(-12) . 9(-11) 5 . 9(-10) 1 . 0(-10)

- <7 . 7(-12) -

- <5 . 8(-11) -

<3 . 3(-10) 4 . 6(-10) 5 . 2(-10)

. . . . .

Abundances are calculated using ; a(b) denotes .
Column densities taken from Sutton et al. (1995). Their CO column density was translated into by taking as for the three W 3 sources.  

The resulting abundances are listed in Table 5 (click here). They reinforce the main conclusion from Paper I that there are large chemical differences between the three sources which cannot be ascribed to excitation effects. Toward IRS4, mostly simple molecules are found. The abundances of these species are low toward IRS5, except those of sulfur- and silicon-containing species. Toward W 3(), the abundances of complex organic species such as , , dimethyl ether and methyl formate are at least 1-2 orders of magnitude larger than those found toward the other two sources. For comparison, the abundances found by Sutton et al. (1995) for the Orion extended ridge and the compact ridge scaled to the same CO abundances are included. In the following, some aspects of the chemistry in the three sources will be discussed in more detail.

Toward IRS4 the abundances of most molecules listed in Table 5 (click here) are expected to be reliable, since their emission fills the 230 and 345 GHz beams to first order. Moreover, there seems to be no indication of a density gradient for these species. Only the SO, and perhaps SiO pose some problems, if they arise in a more localized, hotter (shocked) region. Their listed abundances should be regarded as lower limits.

The abundances of unsaturated radicals and molecules like CN, and are large toward IRS4, suggesting that these molecules are present mostly in the quiescent gas and avoid the shock or ``hot core'' type conditions. This is consistent with observations of the low mass young stellar object IRAS 16293-2422, where CN and were also found to trace the colder, more quiescent outer part of the circumbinary envelope (van Dishoeck et al. 1995). CN and are known to be prominent in the outer zones of photon-dominated regions (PDR) such as the Orion Bar (Jansen et al. 1995; Sternberg & Dalgarno 1995; Simon et al. 1995; Fuente et al. 1995). Indeed, the observed abundances of simple molecules in IRS4 are very similar to those found in the Orion Bar (see Helmich et al. 1997). The extended ridge in Orion also shows enhancements in these species. The ultraviolet radiation from IRS4 (and any other stars in the W 3 region) could create a PDR at the back of the cloud, resulting in large abundances of these species.

The abundances toward IRS5 are more difficult to interpret since there certainly is some stratification of material in the beam. Nevertheless, there is no doubt that the (beam-averaged) abundances of and SiO in this source are larger by at least an order of magnitude compared with IRS4. Surprisingly, however, the abundances of sulfur-bearing molecules containing carbon (CS, ) appear lower than those found toward IRS4. Their line widths are also narrower, suggesting that these species are located in the outer, more quiescent part of the envelope. The proposed connection of and SiO with the outflow in IRS5 needs to be tested further through interferometer observations of these species. If the column density derived by Mitchell et al. (1991) for the outflow is used, the abundances are increased by up to two orders of magnitude.

Although IRS5 has a powerful outflow and emits copious infrared photons, apparently only a small fraction of the envelope has been affected, since the abundances of many species, including HCN and , are lower compared with IRS4 and W 3(). In Paper I, we suggested that at least some of these molecules are still present in the ice mantles on the grains. This case is particularly strong for , where the amount of solid as measured by Allamandola et al. (1992) is 4 orders of magnitude larger than that of gas phase . The amount of depletion for many other molecules is still unknown, however, although most of the CO is certainly located in the gas phase (Tielens et al. 1991). Observations with the Infrared Space Observatory (ISO) will provide a complete inventory of the ice-mantle composition, which together with this work can be used to determine the relative amounts present in the gas and on the grains (cf. van Dishoeck et al. 1996).

Toward W 3() the largest complication is formed by the fact that the relative abundances of the molecules in the more extended core and in the compact clump are poorly determined. Nevertheless, the high temperature, the high methanol abundance, and the large column densities of methyl formate, dimethyl ether and OCS are clear signs of a ``hot core'' chemistry within the beam, in which molecules such as are released from the grains and drive a rapid gas-phase chemistry leading to complex organic molecules. This chemistry has been modelled e.g. by Charnley et al. (1992), and from their work it is obvious that at least must have passed since the onset of methanol release. The amount of methanol relative to methyl formate and dimethyl ether is small, an order of magnitude smaller than found in e.g. the Orion compact ridge (see Table 5 (click here)), and therefore difficult to account for in chemical models. The situation is worse in G34.3 (Millar private communication) where the methanol column density equals the dimethyl ether and methyl formate column densities. High optical depths can influence the methanol abundance severely, but as shown above it is not very likely to be significant here. More detailed chemical modeling will be presented in a later paper (Helmich et al. 1997).

The HCN over HNC ratio has been used to determine time scales for the chemical evolution (Goldsmith et al. 1986; Irvine et al. 1987; Schilke et al. 1992). It is beyond the scope of this work to perform a similar analysis here for the W 3 sources, but it is worthwhile to note that the values for the three W 3 sources lie somewhere between those of Orion IRc2 and the Ridge.

Only few deuterated species have been found in this work. The simplest molecules like , DCN and DNC have been detected and these show indeed signs of some deuterium fractionation suggesting that the cloud had lower temperatures in its past. On the other hand, the observed amount of deuteration is less than the values of 0.01-0.1 found in cold clouds such as TMC-1, and more comparable to that found in other warm clouds (see e.g. Table 2 (click here) of van Dishoeck et al. 1993a for an overview). It is also much lower than found in the low-mass star-forming region IRAS 16293-2422 (van Dishoeck et al. 1995). Together with the relatively low deuteration for HDO (Helmich et al. 1996) this can be interpreted as an indication that the cloud never went through a very cold () phase. The limits on most other species such as HDCO and are not very stringent.

 

DNC/HNC < 7 . < 4 . 4 .

- 2 . 4 .

< 9 . < 0 . 14 < 8 .

< 3 . < 5 . < 2 .

< 7 . < 0 . 25 < 1 .

- -

. . .

Helmich et al. (1996).