The spectra obtained in the line survey are presented in order of increasing frequency in Fig. 4 (click here). Tables 7 (click here) to 14 (click here) summarize the resulting fits to the line profiles, whereas Table 3 (click here) contains the beam-averaged column densities. In the following, the results will be discussed per molecule, describing first the data for IRS4, then for IRS5 and finally for W 3().
The most recent CO maps of W 3 include those by Hasegawa et al. (1994) in and and 6-5 lines and by Oldham et al. (1994) in . The data obtained here on additional isotopic species at just a few positions form a complement to this study, since mapping is beyond the scope of this work.
The CO 3-2 line toward W 3 IRS4 is strong with , but the CO profiles decrease sharply on the red side, which was already noted by Mitchell et al. (1991) and HMMT. In some of the spectra shown here the profile is further affected by emission at the off-position, seen as an extra ``absorption''. No other molecule or CO isotopomer shows a similar line shape. The profile can be affected by absorption by cold foreground material, but it also bears some resemblance to shocked profiles like those found in IC 443 (van Dishoeck et al. 1993). In the case of IRS4, such a shock could be generated at the interface region where the ionization front runs into the cloud. Because IRS4 is at the back, the shock would be coming toward us, consistent with the blue wing. HMMT also find evidence for moving gas associated with IRS4 from pedestal features in the spectra, which they interpret either as outflow or due to the expansion of the HII region. The main beam temperature of is close to the kinetic temperature of 55 K found from the formaldehyde lines. Because the line is heavily optically thick, it suggests that the warm gas fills most of the beam, and that the kinetic temperature is likely somewhat higher than 55 K.
The CO 3-2 spectrum toward W 3 IRS5 shows strong wings, caused by the massive outflow known from other studies (e.g., Mitchell et al. 1991, 1992; Choi et al. 1993, HMMT). Another interesting feature is the strong self-reversal of the 2-1 and 3-2 profiles toward IRS5. Even in the 3-2 line, the absorption is very strong, implying a large amount of cold material in front of IRS5. Our absorption is deeper than that found in HMMT, which may be partly due to lower spectral resolution in the HMMT data. Also, a slight pointing offset from the IRS5 position results in a less strong absorption and a less symmetric profile, as is well illustrated by Figs. 1 (click here) and 2 (click here) of HMMT. HMMT estimate an excitation temperature of 26 K for the foreground gas at the IRS5 position, and find that this temperature is rather constant over the cloud. Note that this value is similar to that found from the ro-vibrational absorption line spectroscopy toward IRS5 (HMMT; Mitchell et al. 1990). The depth of the absorption in our data suggests that even colder or lower excitation material is present along the line of sight.
At least three Gaussian components are needed to fit the central line, its wings and the self-absorption, making the fit somewhat uncertain. It is, however, difficult to obtain any fit for the central emission with smaller than 100 K, which again is close to the kinetic temperature found from the lines (see Paper I), indicating that this gas fills a significant fraction of the beam. Additional evidence for these high temperatures stems from high-J CO 9-8, 12-11, 14-13 and 16-15 measurements (Boreiko & Betz 1991; Betz & Boreiko 1995). On the other hand, no good fit can be obtained for larger than 200 K. Thus, the hotter gas with visible in the ro-vibrational CO absorption lines (Mitchell et al. 1990); HMMT) must originate in a volume which is smaller than the 15 \ beam.
The self-reversal of the CO 3-2 line toward W 3() is less strong than that toward IRS5. Several absorption components appear to be present, although their appearance is again affected by the small beam-switch of +180. The 3-2 line profile also shows strong wings. Although no map of the red and blue components has yet been made, the movement of the water masers (Alcolea et al. 1992; Reid et al. 1995) suggests that an outflow is present at the position. However, the close proximity of the compact HII region W 3(OH) and its unknown contribution to the line profiles can complicate the picture. The fit to the central component is even more difficult and no reliable upper limit can be given. A lower limit to the main-beam temperature is again 100 K, which is probably close to the kinetic temperature of the warm core surrounding both W 3(OH) and W 3(). Higher temperatures () may be present, but are not observed, because the emission from the core is already optically thick.
The data on the CO isotopomers for the three sources are much easier to analyze, since they do not show self-absorption and are almost Gaussian in shape. For all three sources, even the lines are still optically thick. Only the emission is surely optically thin. CO column densities have been derived using the physical parameters found in Paper I and the isotope ratios of Sect. 4.4. The column densities are subsequently determined using , as found for the warm star-forming region NGC 2024 by Lacy et al. (1994). The resulting beam-averaged total CO and column densities are summarized in Table 2 (click here). Note that the column densities are a factor of 3 lower than those listed in Paper I because of the higher adopted CO abundance.
As discussed in Sect. 3, the line has been mapped interferometrically in the direction of W 3(OH)/W 3(). The emission is concentrated in a clump toward W 3(). However, most of the single dish emission () is resolved out, suggesting that indeed the emission observed at the JCMT largely fills the beam.
|(K)||(km s)||(K km s)||(K)||(cm)||(cm)||(cm)||(cm)|
W 3 IRS4
|W 3 IRS5||4.47||4.1||19.5||100|
The molecules in this section are all linear nitrogen-bearing molecules. Due to the non-zero nuclear spin of the nitrogen atom, these molecules show hyperfine splitting of the lower rotational lines, but in the higher lines the splitting is generally smaller than the line widths, so that these blended components are treated as a single line.
CN is one of the few molecules which is more prominent in the direction of W 3 IRS4 than in the other two sources, as can be seen in several spectra around 337 and 340 GHz. Some satellite lines well separated from the main lines can be detected, but do not allow an accurate determination of the optical depth. A complication in determining the CN rotational excitation is that the lines in the 230 and 345 GHz windows stem essentially from just two energy upper levels, both lower than 40 K. The column density derived from the rotation diagram method is therefore very uncertain. Excitation calculations have been performed for this molecule, using a more extended set of collision rates of Bergman (1995, private communication). The / ratio has been used to restrict the gas density to toward IRS4, assuming . Note that the 345 GHz main lines become slightly optically thick ().
For IRS5 the observed line ratio is the same as that found for IRS4, implying a similar density of , even for the slightly higher kinetic temperature of 100 K. The beam-averaged column density is a factor 2 lower than in IRS4. For W 3(), the column density is , similar to IRS5.
Early interferometer measurements by Wright et al. (1995) and single-dish observations of Dickel et al. (1980) and Hayashi et al. (1989) have shown that there is a large concentration of HCN south of IRS4, and a lack toward IRS5. With the current improved receiver technology, not only HCN but also the and isotopomers have been detected toward all three sources, implying high optical depths in the lines of the main isotope. The best candidates for statistical equilibrium calculations are therefore the and 3-2 lines, but unfortunately the 4-3 line is severely blended with , even in the case of IRS4 where little present. Therefore the physical parameters derived in Paper I were used to fit the measured strength of the line.
The corresponding beam-averaged (20'') column density for IRS4 is , whereas the isotopomer ) gives , resulting in and , respectively. The difference between the two numbers is well within the error bars. The excitation temperature inferred from the statistical equilibrium calculations is 26 K. The fact that the optically thick main isotopomer line (J=4-3) has a main beam temperature of 5.6 K implies a source size of more than 9.5.
For IRS5 the same procedure was followed giving and , which correspond to and respectively. Again, the strength of the optically thick line suggests that the emission is beam filling. The failure of Wright et al. (1995) to detect HCN emission in this direction may thus also stem from the fact that the emission is not very concentrated and therefore filtered out by the interferometer.
For W 3() the beam-averaged column densities are and , resulting in and . The large difference suggests that even emission is still somewhat optically thick. If the emission comes from a small source, the inferred column densities are much higher and only the lines stay optically thin. A small HCN source () is consistent with the interferometer observations of Turner & Welch (1984). However, some of the emission must also come from the more extended surrounding envelope, since the excitation temperature combined with the observed main beam temperature of the optically thick HCN indicates a source of .
HNC is of interest for comparison with HCN since the HNC/HCN abundance ratio has been found to differ significantly between cold dark clouds, where it is close to unity, and warm clouds, where HCN can be more abundant by up to two orders of magnitude (Irvine et al. 1987; Goldsmith et al. 1986; Schilke et al. 1992). In W 3, HCN is also observed to be more prominent than HNC. Lines of HNC and were detected toward all three sources, while there is a detection of one line toward W 3(). The lines of the main isotopomer are most likely optically thick, thereby necessitating the use of the isotopomer for statistical equilibrium calculations. The density can in principle be derived from the 4-3/3-2 ratio, but the lines are weak and the signal-to-noise ratio is poor. Also, the 3-2 line can be affected by a blend with , especially toward IRS5. Nevertheless, the line ratios of order unity indicate high densities, consistent with or perhaps even somewhat larger than those derived in Paper I from .
If the physical parameters of Paper I are used to fit the 4-3 line, toward IRS4 and thus . The source size is found to be . For IRS5, ) is found to be and thus the beam-averaged HNC column density is . The (uncertain) source size is , smaller than that found for HCN. Toward W 3(), and thus . Using these parameters the inferred source size is 9. The corresponding HCN/HNC column density ratios for the three sources are , , and , respectively.
Although is a common molecule in the Orion and Taurus clouds, it is still undetected in the W 3 GMC core. The non-detections toward IRS4 and IRS5 may have several reasons. The most important one is that the lines available in the 230 and 345 GHz windows originate from levels that are rather high up in energy (), making it difficult to populate these levels in gas with temperatures below 100 K. Moreover, the molecule has a large dipole moment, so that the critical densities are large (). The upper limits on the lines still provide useful limits on the beam-averaged column density, if the temperature and density from Paper I are adopted. For IRS4 the best upper limit comes from the 24-23 line and is . For IRS5 the same line gives . The column densities are an order of magnitude less than those of HCN toward the two sources.
is detected toward W 3(), and is one of the few more complex unsaturated molecules identified here. Although one of the detected lines is blended, it still gives an upper limit on density () and on temperature (). These limits are not very stringent, because an error of 30% can already bring these values down to the physical parameters of Paper I. Using the Paper I values gives a beam-averaged column density of .
HNCO is not detected toward IRS4, but is marginally present toward IRS5 and clearly toward W 3(). Care has to be taken in analyzing its excitation, since it is well known from observations of the Galactic Center that infrared pumping can affect the populations (Churchwell et al. 1986). Since only a few lines have been detected in this survey, a study like that of Churchwell et al. is not possible, even though intense far-infrared radiation is present toward these two sources and pumping is likely. An alternative is the rotation diagram method, which gives and toward W 3().
Toward IRS5, the rotation temperature was assumed equal to the kinetic temperature from Paper I, because only two weak lines have been observed. The fit of the strength of the 351.633 GHz line results in a (highly uncertain) beam-averaged column density of . If is assumed for IRS4, the non-detections give an upper limit .
CCH is a linear molecule with strong doublet lines toward IRS4. The measurements allow determination of the density from the 4-3/3-2 line ratio, and the results are consistent with those found in Paper I. Using these values, the beam-averaged column density is . Toward IRS5, the molecule is also detected and the line ratio indicates a similar density as found from formaldehyde. The corresponding column density is .
Toward W 3(), the 349 GHz lines of CCH are blended with the lines, but the contribution can be estimated rather well so that this does not pose a problem but only increases the error bars. The density derived from the line ratio is again consistent with that found in Paper I and implies .
is a non-saturated molecule, which can be made quite easily through pure gas-phase chemistry. Three lines have been detected toward IRS4, but the corresponding rotation temperature is uncertain since one of the lines is on the edge of the spectrum. The results are and with an uncertainty of a factor of two (see Fig. 2 (click here)). is not detected toward IRS5, but an upper limit of is found if an excitation temperature of 50 K is assumed.
Many lines are detected toward W 3(), allowing a more accurate rotation diagram than for IRS4. The inferred parameters are and . The rotational temperature seems to indicate that the resides in the same gas as the CO (), and not in the compact ``hot core'' with a much higher temperature, since no lines were detected from very high energy levels (e.g. like methanol, see Paper I).
Formaldehyde and methanol are two particularly useful organic molecules, which can provide a large amount of information about the physical and chemical characteristics of the gas they reside in. Other (saturated) organics like , and are typical for ``hot-core'' type regions and are only detected toward W 3(). Ethanol () also belongs in this category but has not been detected in this survey.
The results of Paper I for IRS4 and IRS5 have not changed, but more information on the isotopomer has been obtained toward W 3(). The line has been detected in two independent sets of observations and indicates that the ratio of integrated intensities lies in between 10 and 40. The column density for ortho- is , which implies for an ortho/para ratio of 3. This value is consistent with derived in Paper I.
The source size determined from the (slightly) optically thick line is larger than the 220 GHz beam, suggesting that the emission comes largely from the warm core surrounding W 3() and W 3(OH). Some is likely to be present as well in the compact ``hot core'', but observations of higher excitation lines and/or isotopomer are needed to probe this region.
Methanol was discussed extensively in Paper I, and the results remain valid except for small corrections based on more extensive data. See Table 3 (click here) for beam-averaged column densities and rotational temperatures. It should be noted that the values for IRS5 are still uncertain, since the detections have large error bars. The methanol lines have been grouped in Tables 8 (click here)-10 (click here). For W 3() the lines are grouped according to the quantum number of the torsional mode (). The different sets were fitted separately, but this did not improve the fit, nor was any systematic trend visible. Therefore, just as in Paper I, a single temperature was used and the scatter in the diagram is explained best with the methanol being subthermally excited.
The lines toward W 3() were fitted with and with large uncertainties. Especially the two lines coming from levels with energies higher than 300 K seem to deviate from the fit. If these two lines are left out, the rotational temperature becomes and the column density . Thus, the column density is fairly robust, whereas the rotational temperature is not. The first estimate is favoured since there is no good reason to delete the two detected lines. The scatter can come from subthermal excitation such as found for methanol, but perhaps infrared pumping and optical depth effects can play a rôle as well.
Statistical equilibrium calculations have been performed and show that the kinetic temperature must be higher than 120 K and the density of order . The ortho and para column densities are 1.0 and respectively. The ortho-para ratio is close to unity, as expected for a warm region.
The above analysis assumed that the emission fills the 230 and 345 GHz beams. However, the interferometer data of Wink et al. (1994) show that the emission comes from an unresolved region, with most of the single dish flux recovered in the interferometer. Assuming a source size of 1, they derive a column density and an excitation temperature of 105 K. If the size is indeed as small as , our inferred column density would increase to if the lines remain optically thin. For such high column densities, however, optical depth effects start to become important and the column density may well be an order of magnitude larger. Recent interferometer observations by Wyrowsky & Walmsley (1996, private communication) suggest an column density a factor larger than that listed in Table 2 (click here) in a region. The corresponding abundance is compared with the beam-averaged value of found here.
Dimethyl ether, , was detected toward W 3() through a number of lines, which is however only a small subset of the transitions available to this molecule. Since no collisional rates are available, a rotation diagram was constructed (Fig. 2 (click here)). This resulted in a rotational temperature of and a beam-averaged column density of . Just as for , two outlying points at high energies affect the fit. Although these appear to be clear detections, a fit has also been made without them, giving and .
Figure 2: Rotation diagram for (upper panel) toward W 3 IRS4 (solid symbols, dashed line) and W 3() (open symbols, solid line). The panels for , and are for W 3() only. Squares represent detected lines, while triangles denote upper limits. For an explanation of the stars and the two fits in the and rotation diagrams, see the text
The same approach was used for methyl formate, and gives and . Note that these column densities are only about one order of magnitude less than found for methanol. They could be increased considerably if most of the emission originates from a source size as small as that of .
The chemistry of sulfur-bearing molecules, especially , is still poorly understood. In this work, a large number of species containing sulfur have been detected, especially toward IRS5.
Carbon monosulfide is well known as a tracer of the densest parts of molecular clouds and indeed CS and 3 of its isotopomers are found toward the three sources in several lines. Tieftrunk et al. (1995) show in their of the W 3 core that the emission is strongest toward IRS4 and southward, but that there is relatively little J=3-2 emission in the direction of IRS5 and the rest of the core. This trend is also found in the higher excitation lines.
Although blending is seldom a problem toward IRS4, it happens that the line is blended with a (weak) methanol line. The integrated line strength can, however, be used as an upper limit. Together with the measured 7-6 line, the ratio has been employed to constrain the density to , similar to that indicated by . A confirmation of this value comes from the line and the upper limit on its 7-6 line. From the absolute strengths of the lines, the column densities are determined and listed in Table 3 (click here). Using the cosmic abundance ratios, . The different isotopomers give values which agree within a factor two, indicating that only CS is optically thick. Note that this column density is in excellent agreement with the results of Tieftrunk et al. (1995). The source size was estimated from the optically thick CS 7-6 line and was found to be beam-filling, in agreement with the map of Tieftrunk et al. (1995).
The density found from the CS isotopes toward IRS5 is also in excellent agreement with the results. The column densities were derived using a kinetic temperature of 100 K, and are listed in Table 3 (click here). From the rare isotopomers, , somewhat lower than toward IRS4. The emission was again found to fill the beam. The column density is accurate to better than a factor of two, but it is six times lower than the value given by Tieftrunk et al. (1995). The reason is unclear since our column densities can explain their measurements of the and 3-2 lines well for the inferred physical parameters.
For W 3() the density found from the 7-6/5-4 lines is also in excellent agreement with the formaldehyde results of Paper I. The column densities for the different isotopomers are listed in Table 3 (click here), and give . The source size was found to be approximately 10.
Using the column densities listed in Table 2 (click here), our CS abundances are larger by up to an order of magnitude than the values given by Tieftrunk et al. (1995) for IRS4 and IRS5, and Wilson et al. (1991) for W 3().
Several isotopomers of SO have been detected toward our sources, including , and . Like CS, the SO itself is optically thick, as is also implied by the large line width of the main lines, although outflowing gas could play a rôle. Although there are many lines available to construct a rotation diagram, it proved difficult to do so: the SO lines suffer from optical depth effects, whereas the lines often turn out to be blended.
The excitation of SO can be studied through statistical equilibrium calculations using the recent collisional cross sections of Green (1994). For IRS4, the observed SO line ratios are not very sensitive to density and temperature, but the values agree with the parameters derived from , although somewhat higher temperature and density are favoured. The beam-averaged column density is , when a temperature of 80 K and a density of are used.
Toward IRS5 not only and lines have been detected but also some , indicating that the main SO lines are very optically thick, . For example, the ratio of the lines of SO and is only 50 instead of the cosmic ratio . Even the lines are still somewhat optically thick. Statistical equilibrium calculations were performed for the lines, resulting in and . Assuming that the SO, just like (see Paper I and below), originates in hotter () and denser () gas than , the column density was calculated to be , implying . From the excitation calculations, the source size is estimated to be . Again, this number is somewhat uncertain, but it shows that the SO likely comes from a compact, warm and dense region.
The SO emission from W 3() is very intense, just as that from IRS5, but no lines are detected and only one line. The statistical equilibrium calculations were therefore performed for . The density and temperature can be constrained to and . Using the temperature and density derived from , is found, resulting in . This beam-averaged column density is fairly robust since it does not change by more than 10% if the density is increased to . The source size was calculated to be large, .
is one of the most interesting molecules found in the W 3 region. In Paper I, its abundance was found to differ greatly between the three sources, with IRS5 showing the strongest emission. Indeed, the abundance has been found to be highly variable in other sources as well (Groesbeck 1994; McMullin et al. 1994). In some cases such as Orion/KL and Sgr-B2(N), the lines are so overwhelmingly strong that it can be one of the major coolants in the 345 GHz range. The excitation temperature was also found to be surprisingly high toward IRS5. With the extended data set presented here, a high temperature has been found toward IRS4 as well. Another big improvement compared with Paper I stems from the recent completion of the catalogue of lines.
Figure 3: Rotation diagram for (open symbols) and (filled symbols) in the three sources. Squares indicate lines that have been detected; triangles denote upper limits; stars are detected lines with , see text. The solid lines represent the least-squares fits using all detected transitions. The dashed line indicates the fit through the lines only. The dotted lines represent the least-squares fits using all detected transitions. The rotation temperatures are those derived from all detected transitions (IRS4 and W 3()) or from the transitions only (IRS5)
Toward IRS4, several additional lines have been detected, mainly because of the better DAS backend. The inclusion of these lines in the rotation diagram (Fig. 3 (click here)) results in and . No lines from have been detected. The column density is, within the errors, consistent with Paper I, but the rotational temperature is increased. The higher temperature is no longer consistent with the kinetic temperature derived from the formaldehyde lines. This could, as for IRS5, be due to optical depth effects, but the upper limits on the column density are not stringent enough to provide clues to the optical depth since is >3.4. More likely, the emission comes from a somewhat warmer region in the direction of IRS4. This possibility is discussed further in Sect. 6.
For IRS5, the complete data set gives and , which is within the errors consistent with Paper I. For the optically thin the following parameters are found: and . From these values it is clear that the ratio of the beam-averaged column densities is not equal to the solar abundance ratio of 22. This was already found in Paper I, where it was attributed to optical depth effects. Since the lines have the lowest intrinsic line strengths and thus the lowest optical depth, a separate rotation diagram has been made for these lines. The results are and . This rotational temperature matches very well with the results, and the beam-averaged column densities are consistent if the cosmic  ratio of 22 is used. The fact that the rotational temperatures are somewhat lower than found in Paper I is thus explained by the high optical depths which flatten the slope of the fit.
In Paper I it was also argued that the lower rotational temperature of may be caused by lack of data at higher energies. Since then, the catalog has been completed, and this lower temperature still persists. At the same time, our list of detected lines has increased and this more complete data set now nicely confirms the results, giving more weight to the optical depth argument. The rotational temperature of is no longer consistent with the rotational temperature of of the outflowing CO (Mitchell et al. 1991). However, a recent re-analysis of the CO excitation by HMMT shows that there may be multiple components with . Furthermore a kinetic temperature of more than 200 K can still result in a rotational temperature of if the is subthermally excited. Since the densities toward W 3 appear to be lower than those toward the Orion Plateau region and since has a large number of (optically thin) lines to radiate through, subthermal excitation cannot be excluded. The optically thick lines lead to a crude determination of the source size. For , the most optically thick line gives a source size of if the optical depth were very high. However, the ratio suggests that the optical depth is still less than unity, so that a larger source size of , similar to that found for SO, is more likely.
The results for W 3() from the rotation diagram are and . The lines give and respectively. Only few lines with have been identified at low signal to noise, but the same trend was found as in IRS5, implying that the main lines are somewhat optically thick. The inferred source size is .
, like , is a near-prolate rotor with transitions that can be used to determine both density and kinetic temperature (see Fig. 9 of Blake et al. 1994). Toward IRS4, two ortho and two para lines have been detected, and good upper limits on several other lines have been obtained. The corresponding temperature range is and density interval . Using 80 K and as the best parameters, and , so that the ratio differs somewhat from the theoretical high-temperature ortho/para ratio of 3. The total column density is the same as found by the rotation diagram method. The rotational temperature is , indicating subthermal excitation.
Toward IRS5 only the line has been detected, prohibiting an analysis as given for IRS4. Using the physical parameters from , is found. Assuming an ortho/para ratio of 3, this results in .
Many more lines were detected toward W 3() allowing a rotation diagram to be constructed. The fitting parameters are and . Statistical equilibrium calculations were performed as well, resulting in and . The column densities are and . Again, the ortho/para ratio differs from three but the sum of the two column densities is close to the rotation diagram value. Most significant however is the low temperature found from the statistical equilibrium calculations. It suggests that the emission does not come from the compact ``hot core'', but from the warmer core surrounding the W 3(OH)/W 3() clumps. This interpretation is strengthened by the fact that the emission from IRS4 comes from a similar region, which is not compact and hot, but extended and warm.
has only few lines available in atmospheric windows, and is therefore not much studied. Therefore the analysis of this molecule is largely based on the single line at 216 GHz. Since collisional rates for this molecule are not known and since no rotational temperature can be derived from a single line, the excitation temperature has to be assumed. Fortunately, the upper energy level of this line lies at 84 K, nicely in the temperature range of the three sources. Moreover, for excitation temperatures between 50 and 250 K, the column density changes by only a factor of three. Taking , we obtain ; ; and for IRS4, IRS5 and W 3() respectively. The upper limits on the line at 214.7 GHz indicate that the 216 GHz line is not very optically thick toward IRS5 and W 3().
For IRS5 the analysis can be improved by using the line at 168.763 GHz toward this source observed by Minh et al. (1991a) in a 35'' beam. They also measured the corresponding line and found the main line to be optically thick. If the two lines are combined, and is derived. While the rotational temperature is of the same order as their adopted excitation temperature, the inferred column density is more than an order of magnitude smaller than that of Minh et al. Using the lines of the rarer isotopomer (assuming the upper limit at 214.7 GHz to be a two sigma detection) gives a somewhat higher rotational temperature and a column density of . Although this value is uncertain, it suggests that the column density of Minh et al. is too large. The reason for this discrepancy is unclear. Therefore the value of derived above is kept.
Minh et al. (1991a) also observed W 3(OH) in , and their 35'' beam included W 3() as well. The rotation diagram for this line and the two sigma upper limit at 214.7 GHz (used as a detection) gives and , corresponding to . This number is highly uncertain, but does not differ much from the value adopted above.
If toward IRS4 is lowered to , similar to the values found for IRS5 and W 3(), its column density would be increased to .
OCS has not been detected toward IRS4. The best limit on the column density comes from the 19-18 line, which gives for the physical parameters of Paper I.
The situation for IRS5 is somewhat better, since two lines were detected which are reasonably far apart in energy. As can be seen from Fig. 11 of Jansen (1995), such transitions are extremely good density tracers. The statistical equilibrium calculations give a density as high as at temperatures of more than 200 K, but the signal-to-noise ratio on the data (especially for the 19-18 line) is poor. If this line were in error by a factor of two or if the source size were small, the density would decrease to and the temperature to . These conditions are similar to those found for SO and , providing additional evidence for a separate sulfur chemistry. If the lower temperature and density are adopted, the beam-averaged column density is found to be .
Toward W 3() many OCS lines are detected and their strengths are large enough to place useful limits on the ratio. The non-detection of at the level indicates that the OCS lines are not very optically thick. The OCS lines often turn out to be blended, but in most cases OCS is the dominant component of these blends. Therefore line ratios can again be used to determine the physical conditions. Both the 28-27/19-18 and 29-28/19-18 ratios point at and , consistent with the results. The inferred is .
Silicon monoxide is a special molecule, because of its intimate connection with shocks in the interstellar medium (Bachiller et al. 1991; Martın-Pintado et al. 1992). SiO has been detected toward all three sources, often in more than one line and in more than one isotopomer. Statistical equilibrium calculations of the 8-7/6-5 ratio were used to constrain the densities.
Toward IRS4, the 8-7 line is very weak. The fit of the 6-5 line and the conditions of Paper I give a column density of .
Toward IRS5, the SiO emission is strong, but the lines are not detected, indicating that the main isotope is not very optically thick. A line has been found coincident with , but this is probably spurious. Therefore, only the SiO lines themselves can be used. The 8-7/6-5 ratio is large, indicating high densities. A smaller source size of , comparable to that for SO and , decreases both parameters somewhat, but they remain high, for . Only for much higher temperatures would lower densities be obtained. This is not unlike the situation close to Orion IRc2 where brightness temperatures of about 1000 K are found for the thermal SiO emission (Wright et al. 1995), indicating a very hot and dense environment. For and the inferred SiO column density is in the 15'' beam. If the physical parameters of Paper I are used together with 8-7 line, decreases to .
W 3() shows not only emission from the main isotopomer, but also from the rarer ones. The lines indicate a density higher than and temperatures higher than 100 K, not unlike the results. is calculated to be , corresponding to . The source size can be estimated since both optically thick and thin lines are available, resulting in >7, so that the SiO emission is relatively widespread.
Although the high densities derived above are most likely evidence for the shock scenario, there remains some uncertainty. MacKay (1995, 1996) provides an alternative explanation in which silane () is evaporated from icy grain mantles. The silane rapidly forms SiO in the subsequent chemistry. On basis of the survey data it is difficult to judge which scenario applies. Higher spatial resolution data are clearly needed.
The observed ions toward IRS5 have been discussed in detail by de Boisanger et al. (1996) and are only briefly summarized here.
The 3-2 and 4-3 lines of and its isotopomers lie at wavelengths with good atmospheric transmission, and are readily detected toward the three sources. The main lines are optically thick and the profile is self-reversed toward W 3(). The ratio of the two lines of the more optically thin variety is a good density indicator.
For IRS4 it is found that the inferred density is rather high, at 55 K. This high value can be due to optical depth effects (even for ), to a small source size, or to calibration errors. The influence on the column density is not very large, however, and is obtained. This value is consistent with that derived from the line. An estimate for the source size of 17 is found from the optically thick line.
De Boisanger et al. (1996) found that the density indicated by the 4-3/3-2 line ratio toward IRS5 is somewhat higher than that found from but still consistent with the latter result within the error bars. The column density has been determined using the parameters of Paper I. No has been detected. The source size is estimated to be 18.
Toward W 3() the 4-3/3-2 ratio indicates a density of at 220 K, somewhat lower than found from . From the line is inferred. This is consistent with , so that . Note that this result does not depend sensitively on temperature. Decreasing the temperature to 100 K increases the column density by only 15%. The source size was calculated from the optically thick 4-3 line. The peak temperature is found to be more than 15 K, from fits obtained by putting a mask around the self-absorption. This results in a source size of . The interferometer map of Wink et al. (1994) shows that some fraction of the emission indeed comes from a compact clump of size , but that most of the emission originates in the more extended core surrounding W 3() and W 3(OH).
Much less information is available on . There has been an intensive search for this ion toward IRS5 (see de Boisanger et al. 1996), and the 8-7 line has been detected in the survey toward W 3().
Assuming that the stems from the same region as toward IRS4, the upper limit on the 8-7 line implies . If the high density found for is taken, the column density decreases to .
Toward IRS5, the lines are very weak, but certainly present. It is interesting that their widths are very small, just as found for the CS isotopomers. The 8-7/6-5 ratio implies an (uncertain) density of , comparable to that found from the . The column density is .
The line was detected toward W 3(), but is unfortunately blended with a methyl formate line. If the methyl formate excitation is in LTE at , the line is not expected to be very strong, but subthermal excitation may enhance its contribution. If we assume that is the dominant component and use the physical parameters from Paper I, is obtained.
is interesting because its abundance provides an indirect determination of that of , which cannot be observed at (sub-)millimeter wavelengths. has been detected toward the Galactic Center in its line (Minh et al. 1991b). In this work, the line was searched toward IRS5, which lies at an energy . The non-detection implies an upper limit on the column density of , using a line strength and assuming an excitation temperature of 50 K.
In the search for ions (de Boisanger et al. 1996), a deep integration on the line at 222 GHz was made toward IRS5, but the line was not detected, giving an upper limit on the column density in a 20'' beam of .
has been detected toward IRS4 and IRS5 by Turner (1994) in 4 lines, who derived a column density of and for IRS4 and IRS5 respectively, assuming a source size of 30''. In this survey the 347740.0 MHz line is detected toward IRS5, and absent toward W 3(). The 348115.2 MHz line is unfortunately blended with , and is thus of limited use. Therefore we adopt the column densities of Turner (1994).
has been searched for toward all three sources by Phillips et al. (1992) through observations of three lines The detection toward IRS5 is confirmed in this survey by the 364 GHz line. The total column density is calculated assuming a high temperature ortho/para ratio of unity. Using the standard physical parameters, toward IRS5, neglecting radiative excitation. However, the 396/364 line ratios observed by Phillips et al. (1992) indicate that higher densities () are more appropriate. This results in , close to the value calculated by de Boisanger et al. (1996). The fact that the 364 GHz line is a factor of 2 stronger in the 15 JCMT beam than in the 20 CSO beam implies a small source size of a few .
The non-detection towards IRS4 implies ) less than .
In the case of W 3() a marginal feature at 364 GHz is present, with a width much smaller than expected from other lines in this survey and the data for the 396 GHz line of Phillips et al. (1992). Using the main beam temperature and the data of Phillips et al. (1992) the uncertain column density is .
De Boisanger et al. (1996) have determined the electron fraction toward IRS5 to be , using many lines and specialized models. For the other two W 3 sources such detailed information is not available, prohibiting a thorough study of the electron abundance toward IRS4 and W 3(). Nevertheless, a lower limit is easily estimated by adding the abundances of the observed molecular ions, which is primarily . For IRS4 this results in an abundance and for W 3() .
Only few deuterated molecules have been detected toward the three sources. One of the most prominent is DCN, which has been seen toward W 3(). Unfortunately the 5-4 line is blended with , but it is estimated that contributes less than 30% of the integrated line flux. Assuming standard physical parameters, is obtained, giving with a factor of four uncertainty. Toward the other two sources, DCN has not been detected. The amount of deuteration, [DCN]/[HCN] and , appears less than for W 3(), but can still be the same within the uncertainties.
From the upper limits on the DNC 3-2 line the following upper limits on N(DNC) are derived: and toward IRS4 and IRS5 respectively. Toward W 3() the DNC 3-2 line has been marginally detected and using the standard physical conditions a column density N(DNC) of (in a 20 \ beam) is inferred. The [DNC]/HNC] is , , and respectively, similar to the [DCN]/[HCN].
is commonly found in star-forming regions (Wootten et al. 1982). Toward IRS4 the line is blended with and thus no column density could be derived. Toward IRS5, the 5-4 line was observed sepearately by de Boisanger et al. (1996). Standard physical parameters give . In the same way, the beam-averaged column density toward W 3() was found to be . This implies ratios of and respectively, again close to the deuteration values for HCN and HNC.
The high beam-averaged column densities found toward the three sources make an inspection of the HDS upper limits of some interest. If the same approach is used as for , i.e., if is assumed, the following HDS upper limits are found from the line for IRS4, IRS5 and W 3() respectively: ; ; and and ; < 0.14 and respectively. Unfortunately, these values provide no useful limits on the deuteration of . If the lower excitation temperature of K is adopted (see Sect. 5.5.5), the limits for IRS5 and W 3() both become but the ratio still does not give useful information.
Another deuterated ion is for which only upper limits on the 3-2 line were obtained, although a hint of the line is present toward IRS5 (de Boisanger et al. 1996). Taking the standard physical parameters, the limits on the column densities are ; ; and for IRS4, IRS5 and W 3() respectively.
Several lines of HDO have been detected toward both W 3() and W 3(OH). The excitation and abundance of this species are described in a separate paper (Helmich et al. 1996). The inferred ratio of is lower than that of other species.
Using the upper limits on the HDCO column density for IRS4, IRS5 and W 3() are and respectively. The ] ratios are ; and respectively. The same method yields and for the column densities of IRS4, IRS5 and W 3() respectively. The ratios are ; < 0.25 and .
W 3 IRS4 W 3 IRS5
X n(H) N(X) n(H) N(X) n(H) N(X) (K) ( cm) (cm) (K) ( cm)
(cm) (K) ( cm)
(cm) - - 7 . 0(22) - - 1 . 3(23) - - 9 . 6(22)
CO 55 1.0 7 . 3(15) 100 1.0 1 . 3(16) 100 2.0 1 . 0(16)
CO 55 1.0 2 . 3(16) 100 1.0 3 . 7(16) 100 2.0 5 . 3(16)
CO - - - - - - 100 2.0 1 . 1(17)
CO - - 1 . 9(19) - - 3 . 3(19) - - 2 . 6(19)
CN 55 1.0 3 . 3(14) 100 1.0 1 . 8(14) 220 2.0 1 . 1(14)
HCN 55 1.0 3 . 1(12) 100 1.0 2 . 1(12) 220 2.0 6 . 6(12)
HCN 55 1.0 1 . 1(13) 100 1.0 5 . 9(12) 220 2.0 1 . 1(13)
HCN - - 8 . 4(14) - - 5 . 3(14) - -
1 . 2(15)
HNC 55 1.0 3 . 3(12) 100 1.0 3 . 2(12) 220 2.0 1 . 5(12)
HNC - - 2 . 0(14) - - 1 . 9(14) - -
9 . 0(13)
HCN 55 1.0 <2 . 5(13) 100 1.0 <9 . 8(12) 220 2.0 1 . 5(13)
HNCO 55 - <2 . 8(13) 100 - 4 . 4(13) 53 - 4 . 8(14)
CCH 55 1.0 1 . 9(14) 100 1.0 1 . 2(14) 220 2.0 1 . 3(14)
CHCH 25 - 7 . 7(14) 50 - <1 . 0(14) 63 - 6 . 1(14)
HCO - - - - - - 220 2.0 5 . 1(12) 55 1.0 8 . 0(13) 100 1.0 8 . 0(13) 220 2.0 4 . 0(14) 28 - 4 . 2(14) 47 - 1 . 6(14) 265 - 8 . 8(15)
CHCN 55 1.0 <1 . 2(13) 100 1.0 <9 . 4(12) 120 4.0 2 . 7(13) - - - - - - 193 - 2 . 0(15) - - - - - - 141 - 6 . 7(14)
CS 55 1.0 6 . 4(12) 100 1.0 1 . 4(12) 220 2.0 1 . 5(13)
CS 55 1.0 6 . 6(12) 100 1.0 2 . 5(12) 220 2.0 2 . 0(13)
CS 55 1.0 1 . 9(13) 100 1.0 1 . 0(13) 220 2.0 4 . 3(13)
CS - - 4 . 0(14) - - 2 . 0(14) - -
1 . 0(15)
SO - - - 200 10 1 . 0(13) - - -
SO - - - 200 10 5 . 3(13) 220 2.0 4 . 3(13)
SO - - - 200 10 8 . 0(13) 220 2.0 5 . 6(13)
SO 80 2.0 1 . 4(14) - - 5 . 0(15) - -
1 . 3(15) 102 - <4 . 4(13) 147 - 3 . 0(14) 179 - 2 . 4(14) 102 - 1 . 5(14) 154 - 5 . 3(15) 184 - 1 . 0(15) 80 2.0 4 . 4(13) 100 1.0 1 . 3(13) 100 2.0 1 . 5(14) 55 - 1 . 1(14) 100 - 2 . 0(14) 220 -
1 . 0(15)
OCS 55 1.0 <5 . 0(13) 150 2.0 8 . 9(13) 220 2.0 2 . 3(14)
SiO - - - - - - 220 2.0 3 . 0(12)
SiO 55 1.0 3 . 7(12) 250 10 5 . 3(12) - -
5 . 9(13)
HCO 55 10 6 . 7(11) - - - 220 2.0 4 . 3(11)
HCO 55 10 6 . 3(12) 100 1.0 3 . 3(12) 220 2.0 3 . 9(12)
HCO - - 3 . 8(14) - - 1 . 9(14) - -
2 . 3(14)
HCS 55 1.0 <2 . 0(13) 100 1.0 6 . 3(11) 220 2.0
1 . 7(13)
HOCO - - - 50 - <2 . 0(12) - - -
HCNH - - - 100 1.0 <1 . 5(13) - - -
HO 55 1.0 <4 . 6(13) 100 5.0 1 . 2(14) 220 2.0 1 . 0(14)
DCN 55 1.0 <3 . 6(12) 100 1.0 <2 . 6(12) 220 2.0 9 . 0(12)
DNC 55 1.0 <1 . 4(12) 100 1.0 <7 . 7(11) 220 2.0 4 . 3(11)
DCO - - - 100 1.0 4 . 8(11) 220 2.0 9 . 5(11)
HDS 55 - <1 . 0(13) 100 - <2 . 7(13) 220 -
<8 . 1(13)
ND 55 1.0 <3 . 1(11) 100 1.0 <3 . 3(11) 220 2.0 <2 . 7(11)
HDCO 55 - <3(12) 100 - <4(12) 220 - <1(13)
CHOD 55 - <3(13) 100 - <4(13) 220 - <1(14) . . .
Apart from the torsional modes of , several other lines originating in vibrationally-excited levels were found in the W 3 survey. Two J= 4-3 lines within the first vibrational level of the bending mode of HCN were detected toward W 3(). This mode lies at much higher energies, approximately 1000 K, than the torsional modes of , so that the levels are unaccessible in most molecular cloud environments. Vibrationally excited HCN was first detected by Ziurys & Turner (1986) in Orion. In principle there are two ways to populate such high states. The first interpretation, favoured by Ziurys & Turner, is through absorption of photons into the bending mode. In the case of W 3(), these photons can be provided by the embedded late O/early B star which heats the dust in the inner part of the envelope and the compact clump surrounding the star. Depending on the actual physical conditions and geometry, the rotational level populations in the ground vibrational state can also be controlled by this infrared pumping (see e.g., Hauschildt et al. 1993). Note that the continuum from W 3(OH) is not able to pump the HCN around W 3(), as Turner & Welch (1984) already concluded.
The second mechanism for populating the vibrationally excited levels is through collisions at sufficiently high densities and temperatures. Wink et al. (1994) estimate a density of on a 1 scale, whereas Wilner et al. (1995) obtain within 0.5''. Although the optically thick dust emission prevents any submillimeter lines from <1'' to be observed, the density in a 1-2'' region is obviously very high and close to the critical density for excitation. More detailed modeling is needed to assess which process dominates in the case of W 3().
The J = 7-6 line within the level of CS was detected in three independent spectra toward W 3(). This vibrational level lies at approximately 1900 K () and is thus even more difficult to populate than that of HCN. This suggests that close to W 3(), either the dust must be very hot or the conditions quite extreme (see also Helmich et al. 1996).
CCH has its lowest vibrational level at or (Hsu et al. 1993; Kanamori & Hirota 1988). At these wavelengths the infrared radiation is intense enough to pump this vibrational level. Black (private communication) predicts the lines within the 345 GHz window to lie at 346248, 346929, 348974 and 349650 MHz. However, no CCH lines within this vibrational level are found in the survey. This suggests that the CCH abundance in the inner region is lower than those of HCN and CS. CCH is also unique in that it has its first electronic state at low energies ().
As in other surveys, some lines remain unidentified, but the fraction is small. The U-lines are generally quite weak, but are thought to be true detections. Because the data are taken in double side-band mode, there are two possible frequencies listed, unless the U-line turns up in more than one spectrum with different local oscillator settings. In those cases, a definite determination of the frequency can be made and only one value is given. Most of the U-lines (7) are found toward IRS5.
|(MHz)||(K)||(km s)||(K km s)|
W 3 IRS4
|W 3 IRS5|
The fact that the three sources have different chemical characteristics provides hints on the identity of the carriers of the lines. Small, non-saturated molecules are mostly found toward IRS4, so there is a large probability of identifying the two U-lines toward this source with a simple molecule. Since the same U-lines are also found toward IRS5, the molecule may involve sulfur. Indeed, the other U-lines toward IRS5, which are not found toward the other sources, most likely contain sulfur or silicon. The U-lines toward W 3() can be due to any species, but most likely originate from some complex organic molecule.
The U-lines were compared with those found in surveys for other sources but there are no U-lines in common with either the Orion surveys of Sutton et al. (1995) and Schilke et al. (1996), the Sgr-B2 survey of Sutton et al. (1991) or the IRAS 16293-2422 survey of Blake et al. (1994).