Apr 21, 2010 - a very early stage of high-mass star formation. High-resolution observations have been carried out toward several IRDCs, where evidences of ...
The Astrophysical Journal, 714:1658–1671, 2010 May 10 C 2010.
doi:10.1088/0004-637X/714/2/1658
The American Astronomical Society. All rights reserved. Printed in the U.S.A.
A SURVEY OF MOLECULAR LINES TOWARD MASSIVE CLUMPS IN EARLY EVOLUTIONARY STAGES OF HIGH-MASS STAR FORMATION Takeshi Sakai1 , Nami Sakai2 , Tomoya Hirota3 , and Satoshi Yamamoto2 2
1 Institute of Astronomy, The University of Tokyo, Osawa, Mitaka, Tokyo 181-0015, Japan Department of Physics and Research Center for the Early Universe, The University of Tokyo, Tokyo 113-0033, Japan 3 National Astronomical Observatory of Japan, Osawa, Mitaka, Tokyo 181-8588, Japan Received 2009 September 20; accepted 2010 March 29; published 2010 April 21
ABSTRACT We have observed the CH3 OH J = 2–1, SiO J = 2–1, C34 S J = 2–1, H13 CO+ J = 1–0, HN13 C J = 1–0, CCH N = 1–0, OCS J = 8–7, and SO JN = 22 –11 lines toward 20 massive clumps, including Midcourse Space Experiment (MSX) 8 μm dark sources (infrared dark clouds) and MSX 8 μm sources, by using the Nobeyama Radio Observatory 45 m telescope. We have found that the velocity widths of the CH3 OH and C34 S lines are broader than those of the H13 CO+ line in the MSX dark sources. On the other hand, they are comparable to the velocity width of the H13 CO+ line in the MSX sources. In addition, the [SiO]/[H13 CO+ ] abundance ratio is found to be enhanced in the MSX dark sources in comparison with the MSX sources. These results suggest that shocks caused by interaction between an outflow and an ambient dense gas would have substantial impact on the chemical composition of the MSX dark sources. The velocity widths of the CH3 OH and C34 S lines relative to that of the H13 CO+ line as well as the [SiO]/[H13 CO+ ] abundance ratio could be used as good tools for investigating evolutionary stages of massive clumps. On the basis of the results, we discuss the chemical and physical evolution of massive clumps. Key words: ISM: clouds – ISM: molecules – stars: formation
several molecules, such as CH3 OH and NH3 , evaporate from dust grains after the onset of star formation. Furthermore, the chemical composition will tell us about physical phenomena occurring in the deep inside of clumps because the emitting region could be different from molecule to molecule. Hence, we can learn star formation activities by observations of molecular lines, even if the internal structure is not resolved with singledish observations. For example, SiO traces shocked regions such as an interacting region between an outflow and an ambient gas, whereas complex organic molecules trace hot molecular cores. With the above motivation, we carried out a systematic survey of several molecular lines, such as CCS JN = 43 –32 , N2 H+ J = 1–0, and CH3 OH JK = 7K –6K , toward 55 massive clumps associated with IRDCs (Sakai et al. 2008, hereafter Paper I). We found that most of the massive clumps are chemically more evolved than low-mass starless cores. In addition, we found that the velocity widths of the CH3 OH JK = 7K –6K line toward several MSX dark sources are broader than those toward the MSX sources. Such broadening would originate from a strong interaction between an outflow and a dense gas in the early stage of protostellar evolution. If so, the difference in the CH3 OH velocity widths reflects a difference in evolutionary stages, and the CH3 OH velocity width can be a useful indicator tracing evolutionary stages of massive clumps. However, we detected the CH3 OH lines only for three MSX sources in Paper I because we employed very high excitation lines in the submillimeter region. Thus, we need more observations to confirm that the large CH3 OH velocity width really reflects the early evolutionary stage of massive clumps. In this paper, we have carried out a survey of the SiO J = 2–1 and CH3 OH JK = 2K –1K lines in the 3 mm wavelength region toward 20 massive clumps, including seven MSX 8 μm dark sources and nine MSX 8 μm sources, in order to investigate the origin of the large velocity width reported in Paper I. If the origin is related to shocks as suggested in Paper I, a similar feature is expected for the SiO line, which is known to be a good shock tracer. In addition, we have also surveyed the C34 S, H13 CO+ ,
1. INTRODUCTION High-mass stars (>8 M ) play important roles in the evolution of a galaxy. However, their formation process is a long standing problem. Toward its resolution, it is very important to understand the initial conditions for high-mass star formation. Since high-mass stars are known to form in massive and dense clumps (e.g., Plume et al. 1997), such clumps without active star formation are thought to be good targets for this purpose. Recently they have been found toward infrared dark clouds (IRDCs). IRDCs are extinction features against the background mid-IR emission (e.g., P´erault et al. 1996; Egan et al. 1998; Simon et al. 2006). Hundreds of massive clumps associated with IRDCs are found by mapping observations with bolometer arrays (Carey et al. 2000; Beuther et al. 2002a; Fa´undez et al. 2004; Hill et al. 2005; Rathborne et al. 2006; Schuller et al. 2009; Vasyunina et al. 2009). Although their typical mass is similar to that of massive clumps with active high-mass star formation, they are cold (10–20 K; Carey et al. 1998; Teyssier et al. 2002; Pillai et al. 2006) and dense (∼106 cm−3 ; Carey et al. 1998). Therefore, massive clumps associated with IRDCs are likely in a very early stage of high-mass star formation. High-resolution observations have been carried out toward several IRDCs, where evidences of high-mass star formation are found (Rathborne et al. 2007; Beuther et al. 2005a; Beuther & Steinacker 2007). Molecular line observations have also been carried out toward several IRDCs (Teyssier et al. 2002; Pillai et al. 2006, 2007; Purcell et al. 2006; Ragan et al. 2006; Beuther & Sridharan 2007; Leurini et al. 2007b). In spite of the above observations, the evolutionary stages of many massive clumps are still controversial. For detailed understanding of high-mass star formation, it is essential to identify their evolutionary stages. As shown in studies of the low-mass starless cores (e.g., Suzuki et al. 1992), chemical composition will be useful for this purpose because it is sensitive to evolutionary stages. In addition, it reflects star formation activities; 1658
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Source
R.A. (J2000.0)
Decl. (J2000.0)
Table 2 Observed Lines VLSR D Reference (km s−1 ) (kpc)
MSX Dark Sources G019.27+00.07 MM1 G022.35+00.41 MM1 G023.60+00.00 MM2 G034.43+00.24 MM3 I18151-1208 MM2 I18223-1243 MM3 I18337-0743 MM3
18 25 58.5 18 30 24.4 18 34 21.1 18 53 20.4 18 17 50.4 18 25 08.3 18 36 18.2
G034.43+00.24 MM2 I18089-1732 MM1 I18151-1208 MM1 I18182-1433 MM1 I18223-1243 MM1 I18264-1152 MM1 I18272-1217 MM1 I18337-0743 MM1 I18385-0512 MM1
18 53 18.6 18 11 51.5 18 17 58.0 18 21 09.2 18 25 10.5 18 29 14.6 18 30 02.9 18 36 41.0 18 41 13.3
−12 03 59 −09 10 34 −08 18 07 01 28 23 −12 07 55 −12 45 28 −07 41 01
26.8 52.7 53.3 59.7 29.7 45.7 56.4
2.3 3.7 3.7 3.5 2.6 3.7 3.8
1 1 1 1 2, 3, 4 2, 3, 4 2, 3, 4
57.8 33.1 33.2 60.0 45.3 43.7 34.3 58.4 26.0
3.5 2.0a 2.8 4.6 3.6 3.4a 2.8a 3.9 1.7a
1 2, 3, 4 2, 3, 4 2, 3, 4 2, 3, 4 2, 3, 4 2, 3, 4 2, 3, 4 2, 3, 4
53.1 57.8 57.7 22.4
3.6 3.5 3.5 2.7
1 1 1 2, 3, 4
MSX Sources 01 24 40 −17 31 29 −12 07 27 −14 31 57 −12 42 26 −11 50 22 −12 15 17 −07 39 20 −05 09 01 Others G024.60+00.08 MM1 G034.43+00.24 MM1 G034.43+00.24 MM4 I18102-1800 MM1
1659
18 35 40.2 −07 18 37 18 53 18.0 01 25 24 18 53 19.0 01 24 08 18 13 11.0 −17 59 59
Note. a Derived from our observation data in this paper by using the rotation curve of Clemens (1985). References. (1) Rathborne et al. 2006; (2) Sridharan et al. 2002; (3) Beuther et al. 2002a; (4) Sridharan et al. 2005.
HN13 C, CCH, OCS, and SO lines in order to explore chemical abundance variation among the massive clumps. On the basis of the results, we discuss how chemical and physical conditions of a massive clump vary with evolution.
Species
Transition
ν(rest) (GHz)
Eu /k (K)
CH3 OH CH3 OH CH3 OH CH3 OH SiO C34 S H13 CO+ HN13 C CCH OCS SO
JK = 2−1 –1−1 E JK = 20 –10 A+ JK = 20 –10 E JK = 21 –11 A− J = 2–1 J = 2–1 J = 1–0 J = 1–0 N = 1–0,J = 3/2–1/2,F = 2–1 J = 8–7 JN = 22 –11
96.739362 96.741375 96.744550 97.582804 86.846960 96.412940 86.754288 87.090850 87.316925 97.301209 86.093950
4.6a 7.0 8.5a 21.6 6.3 6.9 4.2 4.2 4.2 21.0 19.3
Note. a Upper state energy from the lowest level of the E state (1−1 ).
for the observations. The SiO, H13 CO+ , HN13 C, CCH, and SO lines were simultaneously observed with one frequency setting, whereas the CH3 OH, C34 S, and OCS lines were simultaneously observed with another frequency setting. The half-power beam width is 18 and 17 at 87 and 96 GHz, respectively. The main beam efficiency is 0.55 and 0.50 at 87 and 96 GHz, respectively Acousto-optical radiospectrometers (AOS) were employed as a backend. We used two types of AOSs, AOS-H and AOS-W. AOS-Hs have a bandwidth and a frequency resolution of 40 MHz and 37 kHz, respectively, while AOS-Ws have a bandwidth and a frequency resolution of 250 MHz and 250 kHz, respectively. The telescope pointing was checked by observing the nearby SiO maser source, IRC+00363, every 1–2 hr, and was maintained to be better than 5 . The line intensities were calibrated by the chopper wheel method. The system noise temperature was typically 200 K. All the observations were carried out with the position switching mode. The emission-free regions in the Galactic Ring Survey 13 CO J = 1–0 data (Jackson et al. 2006) were employed as the OFF positions, as reported in Paper I.
2. OBSERVATIONS
3. RESULTS AND ANALYSIS
We observed the 14 objects toward which the CH3 OH J = 7–6 line were detected in Paper I. These objects include seven MSX 8 μm dark sources (hereafter MSX dark sources) and three MSX 8 μm sources (hereafter MSX sources). The remaining four sources are not associated with point-like sources and dark parts of the MSX 8 μm and show no contrast in the MSX 8 μm map toward these sources. Thus, these four sources have been identified only by the millimeter-continuum data. We classify these four sources as “Others.” In addition to the above 14 objects, we observed six MSX sources, which are selected from the list by Beuther et al. (2002a). In total, we observed 20 massive clumps listed in Table 1. Distances to the sources are all less than 4.6 kpc. The observed MSX sources, except for G34.43+00.24 MM2, are recognized as high-mass protostellar objects (HMPOs). Note that the Spitzer 24 μm sources are associated with all the observed sources, indicating that star formation has already started. The CH3 OH J = 2–1, SiO J = 2–1, C34 S J = 2–1, H13 CO+ J = 1–0, HN13 C J = 1–0, CCH N = 1–0, OCS J = 8–7, and SO JN = 22 –11 lines were observed by using the Nobeyama Radio Observatory (NRO) 45 m telescope in 2008 March. The line parameters are summarized in Table 2. We used the two side-band separating SIS receiver, T100 (Nakajima et al. 2008),
3.1. Spectra 3.1.1. MSX Dark Sources
Figure 1 shows the spectra of the observed MSX dark sources. In Figure 1, the top trace of each panel shows the CH3 OH spectrum, where three K components of the J = 2–1 transition can be seen. The three components correspond to the JK = 20 –10 E, 20 –10 A+ , and 2−1 –1−1 E from left to right. Hereafter we call these three lines as J = 2–1 triplet. In this figure, we also present the CH3 OH JK = 21 –11 A− line, whose upper state energy (22 K) is slightly higher than those of the JK = 20 –10 E, 20 –10 A+ , and 2−1 –1−1 E lines (5–9 K). The SiO line, which is a good shock tracer, was detected toward all the observed MSX dark sources. The detection rate is higher than that reported by Beuther & Sridharan (2007, ∼40 %). This may originate from the difference in the source selection criteria. We select the MSX dark sources toward which the CH3 OH J = 7–6 line is detected, while Beuther & Sridharan observed almost all the MSX dark sources found by the bolometer array observation (Beuther et al. 2002a). Our observation therefore seems to be biased on active MSX dark sources. In fact, the Spitzer 24 μm sources are associated with all our observed sources.
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4 G19.27+00.07 MM1 G22.35+00.41 MM1 G23.60+00.00 MM2 G34.43+00.24 MM3 2
CH3OH
0
Ta* [K]
CH3OH SiO(x2)
-2 C34S
-4
H13CO+ HN13C
-6
CCH OCS
-8
SO
10
30
50
30
50
70
30
50
70
40
60
80
4 I18151-1208 MM2
I18223-1243 MM3
I18337-0743 MM3
20
30
2
CH3OH
0
Ta* [K]
CH3OH SiO(x2)
-2 C34S
-4
H13CO+ HN13C
-6
CCH OCS
-8
SO
10
30
50
40
60
50
70
VLSR [km s-1] Figure 1. Spectra of the CH3 OH (JK = 20 –10 E, JK = 20 –10 A+ , JK = 2−1 –1−1 E), CH3 OH JK = 21 –11 A− , SiO J = 2–1, C34 S J = 2–1, H13 CO+ J = 1–0, HN13 C J = 1–0, CCH N = 1–0, J = 3/2–1/2, F = 2–1, OCS J = 8–7, and SO JN = 22 –11 lines from top to bottom observed toward the MSX dark sources. The CH3 OH JK = 20 –10 E, JK = 20 –10 A+ , JK = 2−1 –1−1 E lines shown on the top of this figure overlap one another. For clarity, the intensity of the SiO line is multiplied by 2.
The SiO lines look broader than the other molecular lines. The broad linewidth would reflect shocks due to molecular outflows. According to Chambers et al. (2009), the “green fuzzy” feature, which is a 4.5 μm excess due to the shock-excited lines, such as H2 0–0 S(9) and CO ν = 1–0, is seen toward most of our observed MSX dark sources. This supports that shocks are occurring in the observed sources. The C34 S peak intensities are generally weaker than the 13 H CO+ peak intensities, except toward G23.60+00.00 MM2. Toward G34.43+00.24 MM3, the C34 S line is clearly broader than the H13 CO+ line. This indicates that the C34 S line comes from relatively active regions.
The H13 CO+ , HN13 C, and CCH lines are relatively narrow, showing a Gaussian shape. The H13 CO+ and HN13 C line shapes are very similar to each other. This suggests that they come from relatively quiescent regions. The OCS and SO emissions are not detected toward all the MSX dark sources; very weak OCS emission can be seen toward a few sources. 3.1.2. MSX Sources
Figure 2 shows the spectra observed toward the MSX sources. We have detected strong CH3 OH emission toward five sources; G34.43+00.24 MM2, I18089-1732 MM1, I18182-1433 MM1, I18264-1152 MM1, and I18337-0743 MM1. It is reported that
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4 G34.43+00.24 MM2 I18089-1732 MM1
I18151-1208 MM1
I18182-1433 MM1
I18223-1243 MM1
2
CH3OH
0
Ta* [K]
CH3OH SiO(x2)
-2 C34S
-4
H13CO+ HN13C
-6
CCH OCS
-8
SO
30
50
70
10
30
50
10
30
50
40
60
80
20
40
60
4 I18264-1152 MM1
I18272-1217 MM1
I18337-0743 MM1
I18385-0512 MM1
10
30
0
2
CH3OH
0
Ta* [K]
CH3OH SiO(x2)
-2 C34S
-4
H13CO+ HN13C
-6
CCH OCS
-8
SO
20
40
60
30
50
50
70
20
40
VLSR [km s-1] Figure 2. Same as Figure 1, but toward the MSX sources.
hot molecular cores are associated with I18089-1732 MM1 (Beuther et al. 2004, 2005b) and I18182-1433 MM1 (Beuther et al. 2006). In the five sources, the other observed molecular line emissions are also strong, although there is a variation in the intensities among them. This variation suggests that the chemical compositions are slightly different from source to source. The CH3 OH emission is relatively weak toward four sources; I18151-1208 MM1, I18223-1243 MM1, I18272-1217 MM1, and I18385-0512 MM1. However, strong H13 CO+ and CCH emissions are detected toward I18151-1208 MM1 and I182231243 MM1, indicating that a large amount of dense gas exists there. On the other hand, all the emissions are weak toward I18272-1217 MM1 and I18385-0512 MM1, although their CCH intensities are comparable to those of the MSX dark sources. For these two sources, most of the dense gas has been dissipated by
various star formation activities. Hence, they are more evolved than the others, as discussed by Beuther et al. (2002a) on the basis of their density profiles. In fact, the dust temperatures of I18272-1217 MM1 and I18385-0512 MM1 are 52 K and 49 K, respectively, being higher than those of the others. Furthermore, the centimeter-wave continuum emission is relatively strong toward these two objects (Sridharan et al. 2002). 3.1.3. Others
Figure 3 shows the spectra toward the other clumps. Strong CH3 OH lines were detected toward all the four sources. All the molecular lines are strong toward G34.43+00.24 MM1, where even the OCS and SO lines are clearly detected. It is known that there is a hot core containing a high-mass protostar (Rathborne et al. 2008), and hence, the strong emissions toward G34.43+00.24 MM1 seem to be related to the hot core activity.
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I18102-1800 MM1
2 CH3OH
0
Ta* [K]
CH3OH SiO(x2)
-2 C34S H13CO+
-4 HN13C
-6
CCH OCS SO
-8 30
50
70
30
50
70
30
VLSR [km
50
70
0
20
40
s-1]
Figure 3. Same as Figure 1, but toward the “Others” sources.
The SiO line was detected toward all the sources in this category. All the SiO lines show a broad feature, except for G34.43+00.24 MM4. Even for G34.43+00.24 MM4, a weak broad emission could be seen. The C34 S line shapes are different from one another among the four sources. G34.43+00.24 MM1 and G24.60+00.08 MM1 show a broad feature of the C34 S line. On the other hand, the C34 S lines observed toward G34.43+00.24 MM4 and I18102-1800 MM1 have a narrow line shape that is similar to the H13 CO+ line shape. These differences may reflect the difference in evolutionary stages, as discussed in a later section. 3.2. Velocity Widths We fit the observed spectra with a single Gaussian function, except for the CH3 OH J = 2–1 triplet lines. Since the CH3 OH J = 2–1 triplet lines overlap one another, we fit the lines with a triple Gaussian function (see Appendix A). The results are listed in Table 3. Since the observed spectra sometimes show nonGaussian shape, the velocity widths obtained here are “effective” ones. We tried to fit the spectra to the double Gaussian function with broad and narrow components without success because of insufficient signal-to-noise ratios. To investigate the difference from the Gaussian shape, we performed the reduced chi-squared test: if the large velocity width would be due to deviation from the Gaussian shape, the reduced chisquare should be large for the lines with large velocity widths. However, we did not find any dependence of the velocity widths on the reduced chi-square, and hence, we cannot say that the large velocity widths are due to non-Gaussian features, such as a wing-like feature. To investigate how the wing-like feature dominates in the spectra, more sensitive observations are necessary. Figure 4 shows plots of the velocity widths of the observed molecular lines against that of the H13 CO+ line. In these plots, it is apparent that the velocity widths of the CH3 OH, SiO, C34 S,
OCS, and SO lines tend to be larger than that of H13 CO+ for several sources. In contrast, the velocity widths of CCH and HN13 C are similar to that of H13 CO+ . The enhancement of the CH3 OH, SiO, C34 S, OCS, and SO velocity widths cannot originate from the high optical depth because the line shape does not show a flat top profile nor a self-absorption profile. These features imply that these molecular lines come from different parts within a dense clump. In contrast, the situation is different between the MSX dark sources and the MSX sources. We find that the CH3 OH and C34 S velocity widths of the MSX sources are similar to each other and are well correlated with the H13 CO+ velocity width, where the correlation coefficient is 0.82, 0.98, and 0.92 for the CH3 OH J = 2–1 triplet, CH3 OH JK = 21 –11 A− , and C34 S J = 2–1 lines, respectively. On the other hand, the MSX dark sources show the broad widths of the CH3 OH and C34 S lines and no correlation with the H13 CO+ linewidth, where the correlation coefficient is −0.24, 0.11, and 0.22 for the CH3 OH J = 2–1 triplet, CH3 OH JK = 21 –11 A− , and C34 S lines, respectively. Thus, this difference in the velocity width seems to reflect a difference of evolutionary stages. This confirms the results in Paper I, as mentioned in Section 1. Recently, Purcell et al. (2009) also reported a similar trend toward several MSX dark sources and the MSX sources. Figure 5(a) shows a plot of the velocity width of CH3 OH against that of SiO. In the MSX dark sources, the SiO velocity width is broader than the CH3 OH velocity width. Since the broad velocity width of SiO is likely to originate from shocks due to outflows, it would be natural that the large velocity width of CH3 OH is also due to the shocks. The difference of the velocity widths reflects the difference of distributions of SiO and CH3 OH in shocked regions. This point is later discussed in Sections 4.1.1 and 4.1.2. Figure 5(b) shows a plot of the velocity width of CH3 OH against that of C34 S. We can see a rough correlation between them (correlation coefficient ∼0.75), although the CH3 OH
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Table 3 Velocity Widthsa,b Source
CH3 OHc
CH3 OHd
SiO
C34 S
G019.27+0.07 MM1 G022.35+00.41 MM1 G023.60+00.00 MM2 G034.43+00.24 MM3 I18151-1208 MM2 I18223-1243 MM3 I18337-0743 MM3
6.4(0.2) 4.2(0.1) 6.1(0.2) 9.1(0.3) 5.4(0.1) 4.5(0.1) 5.6(0.2)
5.8(0.5) 3.9(0.5) 2.5(0.5) 7.7(0.4) 9.0(1.0) ... ...
13.0(1.0) 7.3(0.5) 28.4(2.8) 12.6(0.5) 8.5(1.0) 6.1(0.3) 14.2(1.3)
Average
5.9(1.6)
5.8(2.7)
12.9(7.5)
H13 CO+
HN13 C
CCH
OCS
SO
7.0(1.0) 3.5(0.7) 3.7(0.3) 6.1(0.6) 3.2(0.6) 4.0(1.0) ...
3.0(0.2) 2.8(0.2) 2.3(0.3) 2.3(0.1) 2.7(0.1) 2.5(0.1) 4.2(0.5)
2.8(0.2) 2.8(0.3) 2.4(0.3) 1.9(0.2) 3.3(0.3) 2.5(0.1) 1.9(0.3)
3.4(0.2) 2.6(0.2) 3.2(0.4) 3.6(0.2) 4.2(0.1) 3.2(0.1) 4.5(0.4)
9.0(3.0) ... ... 4.7(0.9) 4.0(1.0) ... ...
... ... ... ... ... ... ...
4.6(1.6)
2.8(0.7)
2.5(0.5)
3.5(0.6)
5.9(2.7)
...
5.4(0.2) 3.8(0.1) 2.6(0.2) 3.5(0.2) 2.6(0.3) 3.0(0.2) 2.2(0.3) 2.9(0.3) 4.1(0.8)
4.8(0.1) 3.7(0.1) 2.1(0.1) 3.3(0.2) 2.5(0.1) 2.8(0.1) 1.7(0.4) 3.5(0.1) 3.5(0.4)
4.1(0.2) 3.7(0.2) 2.4(0.2) 3.4(0.3) 2.2(0.1) 2.4(0.1) ... 3.4(0.2) ...
5.5(0.2) 3.8(0.2) 2.7(0.1) 3.5(0.1) 3.0(0.1) 3.1(0.1) 3.7(0.2) 3.7(0.1) 4.2(0.2)
5.0(0.8) 5.1(0.4) ... 2.7(0.6) ... 5.0(1.0) ... ... ...
4.7(0.3) 4.0(0.3) ... 3.9(0.4) ... 7.5(0.5) ... ... 7.1(0.8)
3.3(1.0)
3.1(0.9)
3.1(0.7)
3.7(0.8)
4.5(1.2)
5.4(1.7)
MSX Dark Sources
MSX Sources G034.43+00.24 MM2 I18089-1732 MM1 I18151-1208 MM1 I18182-1433 MM1 I18223-1243 MM1 I18264-1152 MM1 I18272-1217 MM1 I18337-0743 MM1 I18385-0512 MM1
5.9(0.1) 3.6(0.1) 3.3(0.2) 3.8(0.1) 4.1(0.4) 4.0(0.1) ... 3.9(0.1) ...
5.8(0.4) 4.3(0.2) 2.4(0.6) 4.4(0.3) ... 3.4(0.1) ... ... ...
10.1(0.3) 15.4(1.7) 2.5(0.4) 5.2(0.4) ... 9.5(0.6) ... 5.8(0.6) ...
Average
4.1(0.8)
4.1(1.3)
8.1(4.6)
Others G24.60+00.08 MM1 G034.43+00.24 MM1 G034.43+00.24 MM4 I18102-1800 MM1
5.4(0.1) 4.7(0.1) 4.6(0.1) 5.0(0.1)
4.8(0.7) 5.2(0.1) 6.5(0.6) 5.1(0.6)
16.1(2.0) 10.9(0.3) 6.6(0.5) 13.0(1.0)
5.1(0.2) 5.4(0.2) 3.5(0.3) 3.9(0.4)
2.5(0.2) 3.4(0.1) 3.2(0.1) 3.4(0.2)
3.2(0.4) 3.0(0.1) 2.6(0.2) 3.2(0.3)
3.7(0.2) 4.0(0.2) 4.0(0.2) 3.3(0.1)
... 5.9(0.8) ... ...
... 6.0(0.3) ... ...
Average
4.9(0.4)
5.4(0.8)
11.7(4.0)
4.5(0.9)
3.1(0.4)
3.0(0.3)
3.8(0.3)
5.9(–)
6.0(–)
Notes. a The numbers in parentheses represent the errors of one standard deviation. b km s−1 . c Obtained by fitting a triple Gaussian function to the CH OH J = 2 –1 E, 2 –1 A+ , and 2 –1 3 K 0 0 0 0 −1 −1 E lines. d Obtained from the CH OH J = 2 –1 A− data. 3 K 1 1
velocity widths tend to be broader than the C34 S velocity width. This may indicate that the emitting regions of the CH3 OH line are similar to that of the C34 S lines within a clump. 3.3. Integrated Intensities The integrated intensities of the observed sources are listed in Table 4. Since millimeter-wave and submillimeter-wave continuum emissions have been observed for most IRDC clumps, we examine the relationship between molecular line emissions and dust emissions. In Figure 6, we plot the integrated intensities of each molecular line against the 1.2 mm peak emissions observed by Beuther et al. (2002a) and Rathborne et al. (2006). The CH3 OH and SiO integrated intensities show no systematic trends between the MSX dark sources and the MSX sources. For these two molecules, the integrated intensities of the MSX dark sources are comparable to those of the MSX sources, although the scatter of the integrated intensity is larger for the MSX sources. This is due to relatively large velocity widths of the MSX dark sources, as shown in Figure 4, and is not due to the optical depth effect of the lines of the MSX sources. A range of the SiO J = 2–1 integrated intensities for the MSX dark sources obtained here is comparable to that for the massive infrared-quiet cores in the Cygnus X region obtained by Motte et al. (2007), which are thought to be in a very early stage of high-mass star formation.
On the other hand, the integrated intensities tend to be larger toward the MSX sources than toward the MSX dark sources for H13 CO+ , C34 S, CCH, and SO and possibly for HN13 C. The integrated intensities of OCS are weaker than those of the other molecules and are not very different between the MSX dark sources and the MSX sources except for I18089-1732 MM1, although many of them are upper limits. The SO intensities of several MSX sources are clearly higher than those of the MSX dark sources. The SO emission is strong toward all the objects whose 1.2 mm peak flux is larger than 1 Jy. These results may suggest that SO is abundant in a hot region. 3.4. Column Densities For simplicity, we derive the column densities, N, of the observed molecules from the integrated intensities by assuming the local thermodynamic equilibrium conditions. To derive the column densities, we have to assume the excitation temperatures of the molecular lines. Since we have observed the CH3 OH J = 7–6 lines toward almost all the objects, except for five sources, in Paper I, we can derive the rotational temperature of CH3 OH from the CH3 OH J = 2–1 and J = 7–6 E-type lines by using the rotation diagram method (e.g., Turner 1991). In Table 5, we list the derived rotation temperatures. Figure 7 shows the plot of the CH3 OH rotation temperature against the NH3 rotation temperature reported in Paper I. The CH3 OH rotation temperature, Trot (CH3 OH), is found to be comparable to the
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SAKAI ET AL. CH3OH (3 lines)
CH3OH JK=21-11 A-
C34S J=2-1
HN13C J=1-0
OCS J=8-7
SO JN=22-11
Vol. 714 SiO J=2-1 ( 3)
CCH N=1-0, J=3/2-1/2, F=2-1
ΔV(X) [km s-1]
12
MSX Dark
8
MSX Source Others
4
0
0
2
4
6
ΔV(H13CO+) [km s-1] Figure 4. Plots of the velocity width of various observed lines against that of the H13 CO+ J = 1–0 line. 12
(a)
30
10
24
8
ΔV(C34S)
ΔV(SiO)
36
18
6
12
4
6
2
0
0
2
4
6
ΔV(CH3OH)
8
10
12
(b)
0
0
2
4
6
ΔV(CH3OH)
8
10
12
Figure 5. (a) Plot of the velocity width of CH3 OH vs. that of SiO J = 2–1. (b) Plot of the velocity width of CH3 OH vs. that of C34 S J = 2–1. The CH3 OH velocity width plotted here is obtained by fitting the three overlapping lines (CH3 OH JK = 20 –10 E, JK = 20 –10 A+ , JK = 2−1 –1−1 E).
NH3 rotation temperature within ±5 K. Hence, we assume that the excitation temperature ranges from Trot (CH3 OH)−5 K to Trot (CH3 OH)+5 K for all the molecules except for CH3 OH. As for CH3 OH, we adopt the 1σ error in the rotation diagram analysis (Table 5) as the error range of Trot (CH3 OH). Since we have observed two hyperfine components of the CCH line, we can derive the optical depth, τ , from the two lines (see Appendix B). The derived optical depths of the CCH line are summarized in Table 7, where the maximum optical depth is 2.17 for G34.43+00.24 MM4. By using the optical depth, we correct the column density by multiplying a factor of
τ /(1−e−τ ); the optical depth of 2.17 corresponds to the correction factor of 2.45. The effect of the optical depth on the derived column densities will be discussed below. The uncertainties of the derived column densities are evaluated by incorporating the following contributions. (1) The statistical error due to the spectrum noise and the intensity calibration error—the former is derived from the observed rms noise of each spectrum, (2) an uncertainty due to unknown excitation temperature—this is estimated by the method described above, and (3) an uncertainty due to the unknown optical depth. This is difficult to estimate. Nevertheless, we know that the optical depth
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Table 4 Integrated Intensitiesa,b Source
CH3 OHc
CH3 OHd
G019.27+0.07 MM1 G022.35+00.41 MM1 G023.60+00.00 MM2 G034.43+00.24 MM3 I18151-1208 MM2 I18223-1243 MM3 I18337-0743 MM3
20.1(0.2) 15.2(0.2) 16.5(0.2) 31.5(0.2) 16.4(0.2) 8.2(0.3) 15.3(0.4)
1.2(0.1) 0.8(0.1) 0.7(0.1) 2.3(0.1) 1.3(0.1) < 0.5 1.0(0.2)
SiO
C34 S
H13 CO+
HN13 C
CCH
OCS
SO
MSX Dark Sources 4.0(0.3) 3.7(0.3) 2.7(0.3) 6.0(0.2) 3.3(0.3) 2.1(0.2) 2.9(0.3)
1.1(0.2) 0.7(0.2) 2.0(0.2) 1.6(0.2) 0.8(0.2) 1.1(0.2) < 0.7
1.3(0.2) 1.5(0.1) 0.5(0.1) 2.0(0.2) 2.4(0.2) 2.3(0.1) 1.1(0.2)
0.9(0.1) 1.2(0.1) 0.6(0.1) 0.7(0.1) 1.4(0.2) 2.3(0.1) 0.6(0.2)
1.4(0.1) 1.1(0.1) 1.4(0.1) 2.7(0.1) 4.6(0.1) 2.8(0.1) 0.9(0.1)
0.8(0.2) < 0.7 < 0.5 1.0(0.2) 0.7(0.2) < 0.8 < 0.7
0.4(0.1) < 0.3 0.3(0.1) 0.5(0.1) 0.4(0.1) < 0.4 < 0.4
5.0(0.2) 2.7(0.2) 2.4(0.2) 2.9(0.2) 3.0(0.1) 4.7(0.2) 0.6(0.2) 2.4(0.1) 0.8(0.2)
2.3(0.2) 2.4(0.1) 1.5(0.1) 1.1(0.2) 1.4(0.1) 1.7(0.2) < 0.5 1.5(0.1) < 0.4
4.4(0.1) 3.3(0.1) 4.8(0.1) 3.3(0.1) 4.5(0.1) 5.2(0.1) 2.0(0.1) 3.3(0.1) 1.8(0.1)
1.2(0.1) 2.3(0.2) < 0.5 0.7(0.2) < 0.7 1.0(0.2) < 0.7 < 0.9 < 0.6
1.2(0.1) 1.7(0.1) 0.4(0.1) 0.9(0.1) < 0.4 1.7(0.1) < 0.4 0.4(0.1) 1.0(0.1)
1.3(0.1) 3.7(0.2) 3.2(0.2) 2.0(0.2)
1.0(0.2) 2.2(0.2) 1.3(0.2) 1.4(0.2)
1.4(0.1) 3.8(0.1) 2.3(0.1) 2.6(0.1)
< 0.5 3.1(0.1) < 0.7 < 0.7
< 0.3 2.0(0.1) 0.6(0.1) 0.5(0.1)
MSX Sources G034.43+00.24 MM2 I18089-1732 MM1 I18151-1208 MM1 I18182-1433 MM1 I18223-1243 MM1 I18264-1152 MM1 I18272-1217 MM1 I18337-0743 MM1 I18385-0512 MM1
32.0(0.2) 10.8(0.2) 4.5(0.2) 11.5(0.3) 3.4(0.3) 19.1(0.3) < 0.8 18.9(0.4) < 0.9
2.1(0.1) 2.0(0.1) 0.6(0.1) 1.9(0.1) < 0.6 1.2(0.1) < 0.4 < 0.5 < 0.4
6.8(0.2) 2.8(0.3) 1.0(0.3) 2.8(0.3) < 0.8 4.8(0.3) < 0.7 2.4(0.2) < 0.8
4.7(0.2) 6.7(0.2) 1.8(0.2) 3.3(0.2) 1.5(0.2) 2.7(0.2) 0.8(0.2) 1.1(0.2) 0.9(0.2)
G24.60+00.08 MM1 G034.43+00.24 MM1 G034.43+00.24 MM4 I18102-1800 MM1
12.7(0.2) 38.7(0.2) 21.2(0.3) 16.7(0.3)
0.8(0.1) 5.6(0.1) 1.0(0.1) 1.0(0.1)
1.8(0.2) 7.8(0.2) 3.9(0.3) 3.1(0.3)
1.3(0.2) 4.6(0.2) 1.9(0.2) 2.1(0.2)
Others
Notes. a The numbers in parentheses represent the errors of one standard deviation. b [K km s−1 ] in T ∗ . a c Total integrated intensity of the CH OH J = 2 –1 E, 2 –1 A+ , and 2 –1 3 K 0 0 0 0 −1 −1 E lines. d Integrated intensity of the CH OH J = 2 –1 A− line. 3 K 1 1
of CCH is less than 2.17 on the basis of the hyperfine analysis. Since the abundances of H13 CO+ , HN13 C, SiO, SO, and OCS are generally lower than that of CCH (e.g., Bergin et al. 1997) and the intensities of these molecular lines are rather weak, it is likely that the observed molecular lines of these molecules are not optically very thick. As for CH3 OH, we use the integrated intensity of the CH3 OH JK = 21 –11 A− line to derive the column density. Since the intensities of this line are rather weak for the observed sources, this line is likely to be optically thin. In Figure 6, the C34 S, H13 CO+ , and CCH integrated intensities are weakly correlated with the 1.2 mm peak emission, where the correlation coefficient is 0.49, 0.66, and 0.54 for C34 S, H13 CO+ , and CCH, respectively. Since the 1.2 mm peak emission traces the total column density along the line of sight, the weak correlation means no heavy saturation of these line emissions, further supporting the optically thin conditions for the lines. We think that the column densities derived here are accurate within a factor of 2. In the following, we use the fractional abundances relative to H13 CO+ (Table 5). Although there are five sources for which we cannot derive the rotation temperature, we derive the column density ratios, assuming that the excitation temperature ranges from 10 K to 50 K. In the calculation of the systematic error, we also assume that the excitation temperature difference between H13 CO+ and the other molecules is less than 10 K. Judging from the velocity width, the H13 CO+ emission would trace relatively quiescent regions. In addition, the H13 CO+ abundance does not vary so much with time according to the chemical model calculations (Bergin et al. 1997; Nomura & Millar 2004). Thus, we think that the H13 CO+ column density
reflects the total amount of dense gas in a clump. Therefore, the N(SiO)/N(H13 CO+ ) ratio, for instance, represents the fraction of the shocked gas in a dense clump. In Figure 8, we plot the abundance ratios of the observed molecules against N(CH3 OH)/ N(H13 CO+ ). The errors indicated in Table 5 and Figure 8 include all the uncertainties mentioned above. 4. DISCUSSION 4.1. Origin of the Molecules 4.1.1. SiO
The SiO molecule is known to be absent in cold quiescent clouds like TMC-1, because Si is almost depleted onto dust grains (e.g., Ziurys et al. 1989). When a shock occurs in a cloud, SiO is released into the gas phase by disruption of silicate grains through sputtering and grain–grain collisions (e.g., Seab & Shull 1983; Schilke et al. 1997). Alternatively, the Si atom is ejected into the gas phase, which reacts with O2 or OH to produce SiO (e.g., Caselli et al. 1997; Schilke et al. 1997). Thus, SiO is generally abundant in a shocked region caused by an interaction between an outflow and a surrounding medium (Mikami et al. 1992; Bachiller & P´erez Guti´errez 1997). In Figure 8, it is remarkable that the N(SiO)/N(H13 CO+ ) ratios of the MSX dark sources are higher than those of the MSX sources. This indicates that the fraction of shocked gas is lower in the MSX sources than in the MSX dark sources, although the observed MSX sources have powerful outflows (Beuther et al. 2002b). Since the H13 CO+ integrated intensity tends to be higher in the MSX sources (Figure 6), the lower N(SiO)/
SAKAI ET AL.
Integrated intensity [K km s-1]
1666
CH3OH (triplet)
CH3OH JK=21-11 A-
C34S J=2-1
H13CO+ J=1-0
Vol. 714
SiO J=2-1
HN13C J=1-0
100 MSX Dark MSX Source
10
Others
1 CCH N=1-0, J=3/2-1/2, F=2-1 OCS J=8-7 0.1 10 102 103 104
SO JN=22-11
F1.2 [mJy] Figure 6. Plots of the observed integrated intensities of various molecules against the 1.2 mm peak flux (Beuther et al. 2002a; Rathborne et al. 2006).
30
Trot(NH3) [K]
25
MSX Dark MSX Source Others
20
15
10 10
15
20
25
30
Trot(CH3OH) [K] Figure 7. Plot of Trot (CH3 OH) vs. Trot (NH3 ).
N(H13 CO+ ) ratios in the MSX sources could partly be due to a large amount of dense gas. However, it seems more likely that the SiO abundance decreases in the MSX sources, because
the SiO integrated intensities of some MSX sources are actually much lower than those of the MSX dark sources (Figure 6). Motte et al. (2007) also indicate that the SiO integrated intensities of the infrared-quiet cores are higher than those of the luminous infrared sources (see Figure 8 in Motte et al. 2007). As mentioned above, a strong shock is necessary for the SiO production in the gas phase. In MSX dark sources, which are in early stages of star formation, a dense gas would still surround the protostars. Hence a strong shock can occur by an impact of an outflow on the ambient dense gas. On the other hand, the outflow cavity may have been relatively large in evolved sources, like MSX sources, so that a strong shock does not occur frequently in the vicinity of the protostars. Thus, SiO cannot be provided efficiently into the gas phase in the evolved sources. Furthermore, it has been suggested that the SiO abundance in the gas phase would decrease with time by depletion onto dust grains or the gas-phase destruction with OH, where the timescale of destruction is estimated to be about 104 yr (Pineau des Forets et al. 1997). Since the lifetime of SiO is shorter than the dynamical timescale of the outflows of the MSX sources (several times 104 yr; Beuther et al. 2002b), SiO tends to be less abundant in the MSX sources. On the basis of the SiO observations of 15 high-mass star forming cores, Miettinen et al. (2006) suggest that the SiO abundance decreases with rising temperature. This can also be interpreted as the SiO abundance decreasing with time. Codella et al. (1999) pointed out that the SiO line profile reflects evolutionary stages; the SiO line has a large velocity width due to shock in the early stage, while it becomes narrower
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Table 5 Abundancesa Source
Trot (CH3 OH) N(H13 CO+ ) N(CH3 OH)b N(SiO) N(C34 S) N(HN13 C) N(CCH) N(OCS) N(SO)) 12 −2 13 + 13 + 13 + (K) 10 (cm ) /N(H CO ) /N(H CO ) /N(H CO ) /N(H13 CO+ ) /N(H13 CO+ ) /N(H13 CO+ ) /N(H13 CO+ ) MSX Dark Sources
G019.27+0.07 MM1
15.0 (1.3)
3.0+1.2 −0.9
165+79 −55
5.8+4.0 −2.4
3.6+2.7 −1.5
1.5+1.2 −0.7
25+21 −12
51+64 −22
G022.35+00.41 MM1
14.9 (1.8)
3.6+1.3 −1.1
91+50 −33
4.5+3.0 −1.8
2.0+1.7 −1.0
1.6+1.2 −0.7
15+14 −8
