discovery of the second warm carbon-chain-chemistry ... - IOPscience

The Astrophysical Journal, 697:769–786, 2009 May 20  C 2009.

doi:10.1088/0004-637X/697/1/769

The American Astronomical Society. All rights reserved. Printed in the U.S.A.

DISCOVERY OF THE SECOND WARM CARBON-CHAIN-CHEMISTRY SOURCE, IRAS15398−3359 IN LUPUS Nami Sakai1 , Takeshi Sakai2 , Tomoya Hirota3 , Michael Burton4 , and Satoshi Yamamoto1 1

Department of Physics and Research Center for the Early Universe, The University of Tokyo, Bunkyo-ku, Tokyo 113-0033, Japan 2 Nobeyama Radio Observatory, Minamimaki, Minamisaku, Nagano 384-1305, Japan 3 National Astronomical Observatory of Japan, Osawa, Mitaka, Tokyo 181-8588, Japan 4 School of Physics, University of New South Wales, Sydney, NSW 2052, Australia Received 2008 August 30; accepted 2009 March 4; published 2009 May 5

ABSTRACT We have conducted a search for carbon-chain molecules toward 16 protostars with the Mopra 22 m and Nobeyama 45 m telescopes, and have detected high excitation lines from several species, such as C4 H (N = 9–8), C4 H2 (J = 100,10 –90,9 ), CH3 CCH(J = 5–4, K = 2), and HC5 N(J = 32–31), toward the low-mass protostar, IRAS15398−3359 in Lupus. The C4 H line is as bright as 2.4 K measured with the Nobeyama 45 m telescope. The kinetic temperature is derived to be 12.6 ± 1.5 K from the K = 1 and K = 2 lines of CH3 CCH. These results indicate that the carbon-chain molecules exist in a region of warm and dense gas near the protostar. The observed features are similar to those found toward IRAS04368+2557 in L1527, which shows warm carbon-chain chemistry (WCCC). In WCCC, carbon-chain molecules are produced efficiently by the evaporation of CH4 from the grain mantles in a lukewarm region near the protostar. Our data clearly indicate that WCCC is no longer specific to L1527, but occurs in IRAS15398−3359. In addition, we draw attention to a remarkable contrast between WCCC and hot corino chemistry in low-mass star-forming regions. Carbon-chain molecules are deficient in hot corino sources like NGC1333 IRAS4B, whereas complex organic molecules seem to be less abundant in the WCCC sources. A possible origin for such source-to-source chemical variations is suggested to arise from the timescale of the starless-core phase in each source. If this is the case, the chemical composition provides an important clue to explore the variation of star formation processes between sources and/or molecular clouds. Key words: ISM: individual (IRAS15398−3359) – ISM: molecules

star-forming regions is not uniform, but shows significant variations between sources. Such chemical variations could originate from differences in the physical processes connected with the star formation. So far, L1527 is the only source which has shown WCCC. It could be that L1527 is an exceptional case. On the other hand, the chemical models by Aikawa et al. (2008) and Hassel et al. (2008) suggest that WCCC can occur more generally when CH4 ice is abundant. Therefore, it is important to investigate whether WCCC characteristics can be found in other starforming regions. Such a study would also be important for the understanding of the origin of chemical variations in starforming regions. With these ideas in mind, we have conducted extensive observations of C2 H and C4 H with the Mopra 22 m5 and C4 H with Nobeyama 45 m telescopes6 (hereafter Mopra and NRO 45 m, respectively) in order to search for new WCCC sources. From observations toward 16 low-mass protostars, we have found a second WCCC source, IRAS15398−3359, with similar characteristics to L1527. In order to investigate its nature further, follow-up observations have then been carried out including a map of C4 H with Mopra, a map of high frequency C2 H spectra with the ASTE7 10 m telescope (Ezawa et al. 2004), and a search for HCO+2 with NRO 45 m. This paper reports on the detailed chemical composition of IRAS15398−2259, along with results from our survey observations for carbonchain molecules.

1. INTRODUCTION We recently discovered an extraordinary richness of carbonchain molecules in the vicinity of the low-mass protostar in L1527 through extensive observations in the centimeter- to millimeter-wave regions. On the basis of this result, we proposed a new carbon-chain chemistry that occurs in a region of warm and dense gas around a protostar, which we called warm carbon-chain chemistry (WCCC; Sakai et al. 2008b). In WCCC, various carbon-chain molecules are produced by reactions of CH4 , which has been evaporated from grain mantles in a lukewarm region. This process is completely different from the classical production scheme for carbon-chain molecules which has been established for cold starless clouds. In such regions, carbon-chain molecules are efficiently produced in the early stages of prestellar evolution, when atomic carbon has not been completely converted to CO (e.g., Suzuki et al. 1992). Following our hypothesis, the concept of WCCC has been confirmed through chemical model simulations by Aikawa et al. (2008) and Hassel et al. (2008). WCCC is opening a new regime in astrochemistry. It provides an opportunity to search for new species. For instance, the molecular anions, C6 H− and C4 H− , were detected in L1527, demonstrating the importance of molecular anions in the chemistry of dense and warm gas in star-forming regions (Sakai et al. 2007, 2008d; Ag´undez et al. 2008; Millar et al. 2007; Herbst & Osamura 2008). It should be emphasized that L1527 is the first source in molecular clouds where C4 H− has been detected. Furthermore, we recently detected the HCO+2 lines in L1527 as well, which provides a valuable clue in understanding where and how gaseous CO2 is formed in low-mass star-forming regions (Sakai et al. 2008a). Most importantly, the discovery of WCCC in L1527 clearly demonstrates that the chemical composition of

5

The Mopra telescope is operated by the CSIRO Australia Telescope National Facility. 6 Nobeyama Radio Observatory is a branch of the National Astronomical Observatory of Japan, National Institutes of Natural Sciences, Japan. 7 The Atacama Submillimeter Telescope Experiment (ASTE) project is driven by NRO, a branch of NAOJ, in collaboration with University of Chile, and Japanese institutes including University of Tokyo, Nagoya University, Osaka Prefecture University, Ibaraki University, and Kobe University.

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Table 1 Source List Source

Cloud

Distance

Reference

(pc)

Position R.A. (J2000)

Telescopesa /Period

Other Names

Decl. (J2000)

“Northern” Sources L1448N L1448C NGC1333 IRAS2A NGC1333 IRAS4B B1 HH211-mm MC27 TMC-1A L1551IRS5

+30◦ 45 15. 5 +30◦ 44 05. 4 +31◦ 14 35. 0 +31◦ 13 09. 0 +31◦ 07 34. 4 +32◦ 00 50. 2 +26◦ 51 39. 0 +25◦ 41 45. 0 +18◦ 08 05. 1

Perseus Perseus Perseus Perseus Perseus Perseus Taurus Taurus Taurus

235 235 235 235 235 235 137 137 137

1 1 1 1 1 1 2 2 2

03h 25m 36.s 56 03h 25m 38.s 84 03h 28m 55.s 4 03h 29m 12.s 0 03h 33m 20.s 8 03h 43m 56.s 8 04h 28m 39.s 3 04h 39m 34.s 9 04h 31m 34.s 14

Cham. I I Lupus 1 ρ Oph. Aquila Rift Serpens R Cr A

200 150 155 120 200 310 170

3 4 5 6 7 7 8

12h 01m 36.s 7 13h 07m 38.s 0 15h 43m 02.s 3 16h 32m 22.s 8 18h 17m 29.s 8 18h 29m 57.s 1 19h 01m 56.s 4

−65◦ 08 48. 2 −77◦ 00 21. 0 −34◦ 09 07. 5 −24◦ 28 34. 0 −04◦ 39 38. 3 +01◦ 13 15. 1 −36◦ 57 27. 0

IRAS11590−6452 DC303.8−14.2 B228

RCrA SMM 1B

M/2007 Oct M/2007 Oct A/08Jun., N/08Jan.–Mar., M/2007 Oct M/2007 Oct N/2008 Jan, M/2007 Oct M/2007 Oct M/2007 Oct

Taurus

137

2

04h 39m 53.s 89

+26◦ 03 11. 0

IRAS04368+2557

M/2007 Oct, [9]

N/2008 Jan, 2007 Jun N/2008 Jan N/2008 Jan N/2008 Jan, 2005 Apr N/2008 Jan N/2008 Jan N/2008 Jan N/2008 Jan N/2008 Jan

L1448mm IRAS03258+3104

L1521F IRAS04365+2535 IRAS04287+1801

“Southern” Sources BHR71 IRS1 IRAS13036−7644b IRAS15398−3359 IRAS16293−2422 L483 Serpens SMM4 RCrA IRAS7B L1527

IRAS18140−0440

Notes. a N represents NRO 45 m, M Mopra, and A ASTE. b The observed position is slightly offset from the IRAS position by (Δα, Δδ) = (30 , −17 ). (1) Hirota et al. 2008; (2) Torres et al. 2007; (3) Bourke et al. 1997; (4) Lehtinen et al. 2005; (5) Lombardi et al. 2008; (6) Loinard et al. 2008; (7) Gregersen et al. 1997; (8) Knude & Hog 1998; (9) Sakai et al. 2008b.

2. OBSERVATIONS 2.1. Observed Sources and Molecules A search for WCCC sources was carried out toward the 16 low-mass star-forming regions listed in Table 1. They are arbitrarily selected from a list of Class 0 or Class I low-mass protostars. A brief description of each source is presented in the Appendix except for IRAS15398−3359, which will be described in Section 4.1. They are placed into three groups; “southern” sources, “northern” sources, and L1527, as shown in Table 1. The “southern” sources were observed with Mopra, although two of them were also observed with NRO 45 m. On the other hand, the “northern” sources were observed only with NRO 45 m. L1527 was observed with both telescopes. Molecular lines observed in this study are summarized in Table 2. The line strength and the upper state energy for each line are also given. The lines observed depend on the sources, as described below.

Table 2 List of Observed Molecules Molecule C2 H

c-C3 H2 l-C3 H2 C4 H2

C4 Hc CH3 CCH

2.2. Observations with Mopra We observed the N = 1−0 lines (87.3–87.4 GHz) of C2 H and the N = 9−8 lines (85.6 and 85.7 GHz for the F1 and F2 components,8 respectively) of C4 H toward the “southern” sources with Mopra. C2 H is the simplest carbonchain molecule, and hence, it is a good indicator for selecting prospective WCCC candidates. Furthermore, C2 H has hyperfine components, enabling the derivation of the column density and the excitation temperature from their intensities. In addition to C2 H, the C4 H lines were also observed. Since C4 H is a longer molecule than C2 H, it is sensitive to the mechanism of carbon-chain production, and therefore to WCCC. We also F1 and F2 represent the fine structure components with J = N + 1/2 and J = N − 1/2, respectively. 8

HC5 N HCO+2

Transition N N N N N N N N

= 1–0, J = 1–0, J = 1–0, J = 1–0, J = 1–0, J = 1–0, J = 4–3, J = 4–3, J

= 3/2–1/2, F = 3/2–1/2, F = 3/2–1/2, F = 1/2–1/2, F = 1/2–1/2, F = 1/2–1/2, F = 9/2–7/2, F = 9/2–7/2, F 43,2 –42,3 41,3 –31,2 40,4 –30,3 101,10 –91,9 100,10 –90,9 101,9 –91,8 N = 9–8, F1 N = 9–8, F2 J = 5–4, K = 0 J = 5–4, K = 1 J = 5–4, K = 2 J = 5–4, K = 3 J = 32–31 40,4 –30,3

= 1–1 = 2–1 = 1–0 = 1–1 = 0–1 = 1–0 = 5–4 = 4–3

Frequency (GHz)

Sa

87.284156 87.316925 87.328624 87.402004 87.407165 87.446512 349.337741 349.339067 85.656431 83.933699 83.165345 88.940237 89.314548 89.687047 85.634009 85.672580 85.457272 85.455622 85.450730 85.442528 85.201340 85.531512

0.17 1.67 0.83 0.83 0.33 0.17 4.89 3.89 1.75 3.75 4.00 9.90 10.00 9.90 9.47 8.47 5.00 4.80 4.20 3.20 32.00 4.00

Eu b (K) 4.2 4.2 4.2 4.2 4.2 4.2 41.9 41.9 29.1 23.4 10.0 37.0 23.6 37.2 20.5 20.6 12.3 19.5 41.1 77.1 67.5 10.3

Notes. a Intrinsic line strength. b Upper state energy. c F and F mean the transitions with J = N + 1/2 and J = N − 1/2, 1 2 respectively.

observed the l-C3 H2 (41,3 –31,2 and 40,3 –30,2 ; 83.9 and 83.2 GHz, respectively) and C4 H2 (100,10 –90,9 and 101,9 –91,8 ; 89.3 and 89.7 GHz, respectively) lines. Furthermore, the CH3 CCH

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DISCOVERY OF SECOND WCCC SOURCE IN LUPUS

(J = 5–4; 85.5 GHz), c-C3 H2 (43,2 –42,3 ; 85.7 GHz), HC5 N (J = 32–31; 85.2 GHz), and HCO+2 (40,4 –30,3 ; 85.5 GHz) lines were observed for IRAS15398−3359. In these observations, the 3 mm HEMT receiver was used as a frontend, whose system noise temperature ranged from 145 K to 175 K. The beam size and the main beam efficiency are 36 and 0.49, respectively, at 86 GHz. The telescope pointing was checked every hour by observing nearby continuum sources, and was maintained to be better than 10 . The back end was a digital autocorrelator, the 8 GHz wide UNSW-MOPS.9 We set the bandwidth and resolution of the individual zoom windows to be 137.5 MHz and 33.6 kHz, respectively. The frequency resolution corresponds to a velocity resolution of 0.12 km s−1 at 86 GHz. Taking advantage of the wide instantaneous bandwidth of the HEMT receiver, we observed the C2 H and C4 H lines simultaneously. The observations were made in the position switching mode, where the off position was taken at Δα = 30 , Δδ = 30 . The line parameters are shown in Tables 3–6 for C2 H, C4 H, l-C3 H2 , and C4 H2 , respectively.

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2.4. Observation with ASTE We also conducted C2 H (N = 4–3; 349.3 GHz) observations toward IRAS15398−3359 using ASTE. The front end for the ASTE observations was the SIS mixer receiver, CATS345. The typical system temperature in single-sideband (SSB) ranged from 250 K to 300 K. Observations were conducted remotely from Nobeyama, by using the N-COSMOS3 network observation system (Kamazaki et al. 2005). The absolute pointing accuracy and main beam efficiency were verified by observing the CO (J = 3–2) emission of II Lup at intervals of every 2 hr. They were better than 2 rms and 0.6 during the observation runs, respectively. The beam size is 22 at 345 GHz. Position switching was employed for the ASTE observation, where the off-position was taken at Δα = 7.◦ 5, Δδ = 30 . The back end was an autocorrelator, MAC, whose bandwidth is set to be 128 MHz. The velocity resolution is 0.11 km s−1 at 349.3 GHz. 3. RESULTS OF THE SURVEY 3.1. “Southern” Sources

2.3. Observations with NRO 45 m For the “northern” sources, we used NRO 45 m. We also observed two “southern” sources, L483 and IRAS15398−3359 with NRO 45 m. In our search for the WCCC sources toward the “southern” sources with Mopra, the C4 H (N = 9–8, F1 ) line (85.6 GHz) was found to be a better tracer than the C2 H (N = 1 − 0) line, and hence, it was employed to search for WCCC sources more efficiently in the “northern” sources. Only the F1 component was observed with NRO 45 m because of the limitation of the back end. The CH3 CCH (J = 5–4), c-C3 H2 (43,2 –42,3 ), HC5 N (J = 32–31), and HCO+2 (40,4 –30,3 ) lines were also observed for IRAS15398−3359. In these observations, we used two Superconductor–Insulator–Superconductor (SIS) mixer receivers (S80 and S100) simultaneously, and a dual-polarization sideband-separating SIS receiver newly installed on NRO 45 m (Nakajima et al. 2008). Typical system temperatures for the latter receiver were about 200 K and 300 K for the two orthogonal polarizations. Typical system temperatures for S80 and S100 were 300 K. Since the image rejection ratio of the new receiver was not well calibrated, the intensity scale was finally corrected by comparing the intensity of a strong line (e.g., the 21,2 –10,1 line of c-C3 H2 ; 85.3 GHz) toward L1527 with that observed with the standard receiver (S100). In addition, we were able to observe the C4 H2 (101,9 –91,8 ; 89.687047 GHz) line toward NGC1333 IRAS4B and L1448N fortuitously during our search for other molecules. In this observation, we employed S80 and S100 as frontends. The main beam efficiency and beam size for S100 are 0.44 and 20 , respectively, at 86 GHz. The telescope pointing was checked by observing the nearby SiO maser source every hour. The pointing accuracy was better than 7 except for L1551. Observing conditions for L1551 were unfortunately poor because of strong winds (∼12 ). Position switching was employed for all observations, where the off-position was taken at Δα = 30 , Δδ = 30 , except for L483, where it was Δα = 150 , Δδ = 0 . The back end was an acousto-optical radiospectrometer (AOS-H), whose bandwidth is 40 MHz. The velocity resolution is 0.13 km s−1 at 86 GHz. 9

The UNSW-MOPS is owned by the University of New South Wales, and funded through an Australian Research Council grant to the Universities of New South Wales, Macquarie and Sydney, together with the CSIRO-ATNF.

3.1.1. Observations of C2 H

As listed in Table 3, eight sources were observed with Mopra, including L1527 for comparison with our previous data. The C2 H lines were detected toward all the sources, with the spectra shown in Figure 1. Six fine and hyperfine components are seen, but their relative intensities vary from source to source. For IRAS15398−3359, they are quite different from the intrinsic relative line strengths listed in Table 2, indicating that the components are optically thick. On the other hand, for Serpens SMM4, RCrA IRAS7B, BHR 71, and IRAS13036−7644, the relative intensities of the components are similar to these intrinsic line strengths, implying the components are close to the optically thin case. The intensity pattern of the components toward IRAS15398−3359 is similar to those toward L1527 (Figure 1). Furthermore, the peak intensity of the weakest components (J = 3/2–1/2, F = 1–1, and J = 1/2–1/2, F = 1–0) toward IRAS15398−3359 are even higher than in L1527. Therefore, we have focused on IRAS15398−3359 as a good candidate for a WCCC source. The column density and excitation temperature for each source are determined from the integrated intensities of the components by the same method described in Sakai et al. (2008b). The line widths of the hyperfine components differ due to the different peak optical depth of each component. Use of the integrated intensities means that an average optical depth over the line shape is considered for each line. The results are summarized in Table 7. IRAS15398−3359 shows the highest column density, which is about half of the column density in L1527. The column densities of Serpens SMM4 and RCrA IRAS7B are comparable to that of IRAS15398−3359, but these are due to the wide line width rather than the peak brightness. In fact, their optical depths are much lower than those in IRAS15398−3359 and L1527, as shown in Table 7. C2 H is widely detected in various star-forming regions including Orion KL and high-mass protostellar objects (HMPOs; Beuther et al. 2008), where the optical depths are generally low (τ  1). Judging from the optical depth, the C2 H spectra observed toward Serpens SMM4 and RCrA IRAS7B are similar to the Orion KL and HMPO cases, rather than the L1527 case. Since simple carbon-chain molecules like C2 H are more easily produced than the longer members, it is difficult to discuss the contribution of WCCC only from the C2 H (N = 1–0) results. Thus, we

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Source BHR71 IRS1

IRAS13036−7644

IRAS15398−3359

IRAS16293−2422

L483

Serpens SMM4

RCrA IRAS7B

L1527

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Vol. 697

Table 3 Line Parameters of C2 H (N = 1−0) Observed with Mopra  TMB a dv a TMB dv (3σ ) −1 (K) (km s ) (K km s−1 )

3/2–1/2, 1–1 3/2–1/2, 2–1 3/2–1/2, 1–0 1/2–1/2, 1–1 1/2–1/2, 0–1 1/2–1/2, 1–0 3/2–1/2, 1–1 3/2–1/2, 2–1 3/2–1/2, 1–0 1/2–1/2, 1–1 1/2–1/2, 0–1 1/2–1/2, 1–0 3/2–1/2, 1–1 3/2–1/2, 2–1 3/2–1/2, 1–0 1/2–1/2, 1–1 1/2–1/2, 0–1 1/2–1/2, 1–0 3/2–1/2, 1–1 3/2–1/2, 2–1 3/2–1/2, 1–0 1/2–1/2, 1–1 1/2–1/2, 0–1 1/2–1/2, 1–0 3/2–1/2, 1–1 3/2–1/2, 2–1 3/2–1/2, 1–0 1/2–1/2, 1–1 1/2–1/2, 0–1 1/2–1/2, 1–0 3/2–1/2, 1–1 3/2–1/2, 2–1 3/2–1/2, 1–0 1/2–1/2, 1–1 1/2–1/2, 0–1 1/2–1/2, 1–0 3/2–1/2, 1–1 3/2–1/2, 2–1 3/2–1/2, 1–0 1/2–1/2, 1–1 1/2–1/2, 0–1 1/2–1/2, 1–0 3/2–1/2, 1–1 3/2–1/2, 2–1 3/2–1/2, 1–0 1/2–1/2, 1–1 1/2–1/2, 0–1 1/2–1/2, 1–0

0.312(15) 1.880(24) 1.092(17) 1.107(17) 0.561(15) 0.325(13) 0.555(22) 2.015d 1.119(24) 1.138(24) 0.783(21) 0.543(21) 1.740(14) 3.270(22) 2.499(16) 2.681(18) 2.214(18) 1.741(17) 0.312(13) 0.784(14) 0.650(11) 0.637(13) 0.434(13) 0.251(12) 0.952(21) 1.489(20) 1.253(20) 1.291(20) 1.138(21) 0.968(24) 0.290(13) 1.831(16) 1.013(14) 1.035(14) 0.509(13) 0.282(12) 0.294(13) 2.128(20) 1.192(15) 1.221(18) 0.521(14) 0.328(14) 1.699(30) 2.406(31) 1.701(28) 1.846(29) 1.622(28) 1.571(29)

1.06(6) 1.03(2) 1.03(2) 1.03(2) 1.03(3) 1.10(5) 0.54(2) ... 0.85(2) 0.91(2) 0.67(2) 0.52(2) 0.311(3) 0.426(3) 0.382(3) 0.377(3) 0.325(3) 0.306(3) 0.74(4) 1.55(3) 1.21(2) 1.25(3) 0.91(3) 0.94(5) 0.45(1) 0.78(1) 0.65(1) 0.62(1) 0.49(1) 0.41(1) 2.07(11) 2.28(2) 2.30(4) 2.28(4) 2.20(6) 2.18(11) 1.12(6) 1.36(1) 1.30(2) 1.27(2) 1.27(4) 1.05(5) 0.49(1) 0.57(1) 0.55(1) 0.54(1) 0.51(1) 0.47(1)

0.348(32) 2.119(29) 1.193(28) 1.200(27) 0.593(28) 0.369(29) 0.358(24) 1.830 1.105(30) 1.157(32) 0.647(25) 0.356(22) 0.585(9) 1.514(12) 1.043(11) 1.097(12) 0.778(11) 0.572(11) 0.251(19) 1.244(28) 0.865(23) 0.860(26) 0.454(22) 0.267(22) 0.444(19) 1.423(26) 0.731(23) 0.936(22) 0.663(19) 0.455(20) 0.607(56) 4.557(62) 2.465(59) 2.512(58) 1.195(55) 0.630(54) 0.363(29) 3.318(35) 1.790(34) 1.885(30) 0.753(33) 0.372(31) 0.892(26) 1.460(26) 0.968(28) 1.037(29) 0.858(28) 0.773(25)

rmsb (mK)

VLSR a,c (km s−1 )

10.1 9.4 9.0 8.8 9.0 8.8 15.0 15.3 11.8 11.8 12.5 14.0 9.8 9.3 9.4 10.6 11.0 11.5 8.6 6.0 6.2 6.8 8.1 7.7 14.3 11.1 11.8 12.0 13.3 16.3 9.0 9.1 8.6 8.4 8.4 8.3 8.6 8.6 8.8 8.0 8.7 10.0 17.8 15.1 16.8 18.0 18.5 18.2

−4.23 −4.41 −4.36 −4.40 −4.40 −4.25 3.82 ... 3.81 3.75 3.69 3.82 5.25 5.19 5.23 5.16 5.12 5.26 4.09 3.95 4.06 3.98 3.94 4.10 5.44 5.43 5.45 5.39 5.31 5.45 8.28 8.00 8.12 8.05 8.09 8.21 5.56 5.50 5.52 5.46 5.41 5.53 5.95 5.90 5.95 5.88 5.83 5.95

Notes. The numbers in parentheses represent the errors in units of the last significant digits. a Obtained by a Gaussian fit. b The rms noise averaged over the line width. c Rest frequency errors (typically an order of ∼0.01 km s−1 ) are much larger than Gaussian fitting errors. d This line cannot be fitted by a Gaussian. T MB represents the peak intensity.

have also examined the C4 H, l-C3 H2 , and C4 H2 spectra in the “southern” sources. These species are usually deficient in active star-forming regions like Orion KL. 3.1.2. Observations of C4 H

The N = 9–8, F1 and F2 lines of C4 H were detected with a signal-to-noise ratio (S/N) greater than 10σ in the integrated intensity toward all the sources except for RCrA IRAS7B (Figure 2). However, the line intensities of the C4 H lines vary

greatly from source to source, in contrast to the C2 H case. In particular, the C4 H lines observed toward IRAS15398−3359 are as strong as 1.37 ± 0.02 K (F1 ) and 1.26 ± 0.02 K (F2 ), even stronger than toward L1527 with Mopra (1.19 ± 0.03 K for F1 and 1.09 ± 0.04 K for F2 ). Note that the C4 H intensities observed toward L1527 with Mopra are weaker than those observed with NRO 45 m (1.70 ± 0.02 K for F1 and 1.61 ± 0.02 K for F2 ; Sakai et al. 2008b). Since the size of the C4 H distribution in L1527 is about 40 , beam dilution effects are likely more

No. 1, 2009

Source BHR71 IRS1 IRAS13036−7644 IRAS15398−3359 IRAS16293−2422 L483 Serpens SMM4 RCrA IRAS7B L1527

DISCOVERY OF SECOND WCCC SOURCE IN LUPUS

Transitiona F1 F2 F1 F2 F1 F2 F1 F2 F1 F2 F1 F2 F1 F2 F1 F2

Table 4 Line Parameters of C4 H (N = 9−8) Observed with Mopra  TMB dv (3σ ) TMB b dv b −1 (K) (km s ) (K km s−1 ) 0.326(15) 0.291(14) 0.248(23) 0.214(19) 1.367(17) 1.263(17) 0.129(14)e 0.112(13)e 0.650(24) 0.554(23) 0.110(10) 0.089(9) ... ... 1.188(33) 1.094(35)

0.93(2) 0.92(5) 0.49(5) 0.63(6) 0.279(4) 0.275(4) 0.58(7) 0.68(9) 0.35(2) 0.37(2) 1.40(15) 1.46(7) ... ... 0.40(1) 0.43(2)

0.315(29) 0.276(27) 0.128(24) 0.134(25) 0.404(10) 0.366(10) 0.079(16) 0.072(18) 0.259(18) 0.240(18) 0.162(31) 0.139(59)  0.039  0.043 0.519(25) 0.503(24)

773

rmsc (mK)

b,d VLSR (km s−1 )

10.3 9.7 15.9 13.2 12.1 11.8 9.2 9.0 17.0 16.1 7.3 11.4 9.3 10.2 20.5 18.5

−4.37 −4.35 3.81 3.83 5.17 5.14 4.03 3.96 5.35 5.31 8.08 8.08 ... ... 5.85 5.83

Notes. The numbers in parentheses represent the errors in units of the last significant digits. A line width of 1.4 km−1 is assumed for RCrA IRAS7B. a F and F denote the transitions with J = N + 1/2 and J = N − 1/2, respectively. 1 2 b Obtained by a Gaussian fit. c The rms noise averaged over the line width. d Rest frequency errors (typically an order of ∼0.01 km s−1 ) are much larger than Gaussian fitting errors. e It cannot be well fitted by a Gaussian. The peak intensities are 178 mK and 153 mK for F and F , respectively. 1 2

Source

Transition

BHR71 IRS1 IRAS13036−7644 IRAS15398−3359

41,3 –31,2 41,3 –31,2 41,3 –31,2 40,3 –30,2 41,3 –31,2 41,3 –31,2 40,3 –30,2 41,3 –31,2 41,3 –31,2

IRAS16293−2422 L483 Serpens SMM4 RCrA IRAS7B

Table 5 Line Parameters of l-C3 H2 Observed with Mopra  TMB a dv a TMB dv (3σ ) −1 (K) (km s ) (K km s−1 ) ... ... 0.204(18) 0.122(20) ... 0.188(30) 0.140(28) ... ...

... ... 0.31(3) 0.40(8) ... 0.27(5) 0.28(7) ... ...

0.021 0.029 0.068(11) 0.056(17) 0.027 0.060(17) 0.053(17) 0.069 0.045

rmsb (mK)

VLSR a,c (km s−1 )

14.3 19.4 12.3 14.3 9.0 21.0 20.0 11.4 10.7

... ... 5.09 5.03 ... 5.34 5.33 ... ...

Notes. The numbers in parentheses represent the errors in units of the last significant digits. The line widths of 0.5, 0.5, 1.0, 2.0, and 1.4 km s−1 are assumed for BHR71 IRS1, IRAS13036−7644, IRAS16293−2422, Serpens SMM4, and RCrA IRAS7B, respectively. a Obtained by a Gaussian fit. b The rms noise averaged over the line width. c Rest frequency errors (typically an order of ∼0.01 km s−1 ) are much larger than Gaussian fitting errors.

Source

Transition

BHR71 IRS1

100,10 –90,9 101,9 –91,8 101,9 –91,8 100,10 –90,9 101,9 –91,8 101,9 –91,8 101,9 –91,8 101,9 –91,8 101,9 –91,8

IRAS13036−7644 IRAS15398−3359 IRAS16293−2422 L483 Serpens SMM4 RCrA IRAS7B

Table 6 Line Parameters of C4 H2 Observed with Mopra  TMB dv (3 σ ) TMB a dv a (K) (km s−1 ) (K km s−1 )

rmsb (mK)

VLSR a,c (km s−1 )

0.019 0.040(24) 0.024 0.020(10) 0.036(11) 0.025 0.024 0.052 0.039

12.8 15.3 15.9 12.5 10.4 8.3 16.0 8.7 9.3

... −4.56 ... 4.98 5.13 ... ... ... ...

... 0.075(11) ... 0.071(18) 0.074(15) ... ... ... ...

... 0.51(8) ... 0.25(7) 0.37(8) ... ... ... ...

Notes. The numbers in parentheses represent the errors in units of the last significant digits. The line widths of 0.5, 0.5, 1.0, 2.0, 0.4, and 1.4 km s−1 are assumed for BHR71 IRS1 (for a para line), IRAS13036−7644, IRAS16293−2422, Serpens SMM4, L483, and RCrA IRAS7B, respectively. a Obtained by a Gaussian fit. b The rms noise averaged over the line width. c Rest frequency errors (typically an order of ∼0.01 km s−1 ) are much larger than Gaussian fitting errors.

SAKAI ET AL.

TMB [K]

774 3.5 3 BHR71 IRS1 2.5 2 1.5 1 0.5 0 -0.5 87.28 87.32 87.36 3.5 3 IRAS13036-7644 2.5 2 1.5 1 0.5 0 -0.5 87.28 87.32 87.36 3.5 3 IRAS15398-3359 2.5 2 1.5 1 0.5 0 -0.5 87.28 87.32 87.36 3.5 3 IRAS16293-2422 2.5 2 1.5 1 0.5 0 -0.5 87.28 87.32 87.36

87.4

87.4

87.4

87.4

Vol. 697

3.5 3 L483 2.5 2 1.5 1 0.5 0 87.44 -0.5 87.28 87.32 87.36 3.5 3 Serpens SMM4 2.5 2 1.5 1 0.5 0 -0.5 87.28 87.32 87.36 87.44 3.5 3 RCrA IRAS7B 2.5 2 1.5 1 0.5 0 87.44 -0.5 87.28 87.32 87.36 3.5 3 cf; L1527 2.5 2 1.5 1 0.5 0 -0.5 87.28 87.32 87.36 87.44

87.4

87.44

87.4

87.44

87.4

87.44

87.4

87.44

Frequency [GHz] Figure 1. Line profiles of C2 H (N = 1−0) observed with Mopra.

significant for the Mopra data than the NRO 45 m. In contrast, the C4 H lines are weak in Serpens SMM4, and are not detected in RCrA IRAS7B. Although these two sources show fairly high column densities of C2 H, we can now rule them out from the WCCC candidates on the basis of these C4 H results. 3.1.3. l-C3 H2 and C4 H2

The C4 H2 and l-C3 H2 lines were detected in IRAS15398−3359, whereas they were scarcely detected in the other sources (Tables 5 and 6). Exceptions are L483 for the l-C3 H2 line and BHR71 IRS1 for the C4 H2 lines (Figure 3). Since l-C3 H2 (4.1 D) and C4 H2 (4.5 D) have much larger dipole moments than C4 H (0.9 D), their lines have higher critical densities (greater than 106 cm−3 ) than the C4 H lines. Detection of these lines implies that the carbon-chain molecules exist in a dense region. Therefore, it is very likely that IRAS15398−3359 is the second WCCC source to be discovered. In view of detections of the C4 H2 and C3 H2 lines toward BHR71 IRS1 and L483, respectively, it is possible that they may be WCCC sources too.

Table 7 Column Density of C2 H Sources

Tex (K)

τa

N (1014 cm−2 )

BHR71 IRS1 IRAS13036−7644 IRAS15398−3359 IRAS16293−2422 L483 Serpens SMM4 RCrA IRAS7B L1527

6.4(6) 7.3(5) 7.2(4) 4.4(1) 6.1(7) 8.1(8) 12.1(17) 5.2(3)

0.08(2) 0.14(3) 0.42(11) 0.18(1) 0.25(12) 0.05(1) 0.04(1) 1.14(47)

1.6(2) 1.8(2) 3.1(7) 1.6(1) 1.9(7) 3.1(3) 2.5(2) 7.3(28)

Notes. The numbers in parentheses represent the errors in units of the last significant digits. a Optical depth of the weakest hyperfine components (J = 3/2–1/2, F = 1–1 and J = 1/2–1/2, F = 1–0).

even higher than in L1527 (1.70 ± 0.02 K; Sakai et al. 2008b).

3.2. “Northern” Sources

4. IRAS15398−3359 IN LUPUS

On the basis of our experience of the Mopra observations toward the southern sources, the C4 H (N = 9–8) line is better than the C2 H (N = 1–0) line for a search for WCCC candidates, and hence, we also searched for the C4 H (N = 9–8, F1 ) line toward the nine sources listed in Table 1 with NRO 45 m. In addition, IRAS15398−3359 and L483 were observed as well. The line was detected toward all the sources observed, except for NGC1333 IRAS4B. The line profiles are shown in Figure 4, with their line parameters listed in Table 8. While this line was detected with moderate intensity toward B1, L483, HH211mm, and MC27, the intensity toward IRAS15398−3359 is outstanding in comparison with the other sources. The C4 H line is as strong as 2.43 ± 0.04 K in IRAS15398−3359,

4.1. Source Description According to the above results, we regard IRAS15398−3359 as the best candidate for a WCCC source. IRAS15398−3359 is a protostar located in the Lupus 1 cloud (D = 155 pc; Lombardi et al. 2008), the largest component of the Lupus molecular cloud (Tachihara et al. 1996), and is believed to be a borderline object between Class 0 and Class I. A molecular outflow with a dynamical timescale of ∼2 × 103 yr is associated with the protostar (Tachihara et al. 1996). It is the only outflow source known in Lupus 1. The outflow axis is almost along the line of sight, and hence, the protostellar disk of IRAS15398−3359 should have a face-on configuration. Indeed, no signature of

No. 1, 2009

DISCOVERY OF SECOND WCCC SOURCE IN LUPUS

2.5 2

2.5 BHR71 IRS1

2

RCrA IRAS7B

2

1.5

1.5

1

1

1

0.5

0.5

0.5

0

0

0

0.4 0.3 0.2 0.1 0 -0.1

1.5

-15 -10 2.5 IRAS13036-7644 2

TMB [K]

2.5

IRAS16293-2422

-5

0

-5

5

0

5

10

-5

2.5 2

1.5

L483

1

1

0.5

0.5

0

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0

5

10

2

-5 2.5

IRAS15398-3359

2

1.5

0

1

1 0.5

0

0

-5

0

5

10

5

10

15

5

10

15

-5

0

5

10

15

Serpens SMM4

1.5

0.5

0

10

1.5

1

0

5

cf; L1527

2

0.5 -5

0

2.5

1.5

2.5

775

15

0.4 0.3 0.2 0.1 0 -0.1

0

5

10

0

15

5

10

15

VLSR [km/s] Figure 2. Line profiles of C4 H (N = 9 − 8, F1 ) observed with Mopra. The dashed lines represent VLSR of each source. The inserts for Serpens SMM4 and RCrA IRAS7B are the spectra with appropriate temperature and velocity scales showing their detailed features.

0.4

TMB [K]

0.3

0.4 l-C3H2

L483

41, 3-31, 2

0.3

0.4 l-C3H2

L483

40, 3-30, 2

0.3

0.2

0.2

0.2

0.1

0.1

0.1

0

0

0

-0.1

-0.1

-0.1

-0.2 -5

0

5

10

-0.2 15 -5

0

5

10

C4H2

BHR71 IRS1

-0.2 15 -15

-10

-5

101, 9-91, 8

0

5

VLSR [km/s] Figure 3. Line profiles of l-C3 H2 and C4 H2 observed with Mopra.

infall is seen in the H2 CO and CS observations reported by Mardones et al. (1997). Although it was first recognized more than 10 years ago, no systematic study of its chemical composition has yet been undertaken. 4.2. Abundant Carbon-chain Molecules in IRAS15398−3359 In order to confirm that IRAS15398−3359 is the second WCCC source, further observations were conducted with Mopra and NRO 45 m. First, lines from the related molecules detected in L1527 were searched for with Mopra. In addition to lines from C2 H, C4 H, C4 H2 , and l-C3 H2 mentioned before, high excitation lines of c-C3 H2 (43,2 –42,3 , Eu = 29.1 K) and CH3 CCH (J = 5– 4, K = 2, Eu = 41.1 K) were also detected (Table 9). Then, we undertook observations to confirm the C4 H, C4 H2 , c-C3 H2 , and CH3 CCH lines toward IRAS15398−3359 with NRO 45 m, for detailed comparisons with the equivalent data from L1527. Although IRAS15398−3359 is low in elevation as seen from Nobeyama, it can still be observed for 1.5 hr per day. These results are summarized in Table 9. The integrated intensities of the lines observed with NRO 45 m are generally

Table 8 Line Parameters of C4 H (N = 9–8, F1 ) Observed with NRO 45 m  a TMB dv (3σ ) rmsb Source TMB a dv a VLSR (K) (km s−1 ) (K km s−1 ) (mK) (km s−1 ) L1448N L1448C NGC1333IRAS2A NGC1333IRAS4B B1 HH211-mm MC27 TMC-1A L1551IRS5 L483 IRAS15398−3359

0.398(45) 0.316(48) 0.288(63) ... 0.357(36) 0.799(55) 1.161(76) 0.203(73) 0.301(51) 0.772(40) 2.425(35)

1.07(14) 0.98(17) 0.60(15) ... 1.31(15) 0.59(5) 0.49(4) 0.55(21) 0.72(6) 0.52(3) 0.282(5)

0.103(32) 0.299(100) 0.202 0.124 0.543(102) 0.601(63) 0.594(78) 0.147(86) 0.229(77) 0.467(44) 0.791(21)

31.9 34.1 45.1 41.2 25.9 35.8 53.7 52.0 35.9 28.2 24.6

4.83(6) 4.78(7) 7.83(7) ... 6.59(7) 9.08(2) 6.47(2) 6.10(10) 6.63(6) 5.34(1) 5.19(1)

Notes. The numbers in parentheses represent the errors in units of the last significant digits. The line width of 1.0 km s−1 is assumed for NGC1333 IRAS4B. a Obtained by a Gaussian fit. b The rms noise averaged over the line width.

776

SAKAI ET AL. 2.5

B1

2 1.5

1.5

1

1

1

0.5

0.5

0.5

0

0

2.5 2

0

5

10

0

15

-5

0

5

10

-5

15

IRAS15398-3359

2

2

1.5

1.5

1.5

1

1

1

0.5

0.5

0.5

0

0

-5 2.5 2

0

5

10

0

5

10

0

HH211-mm

2

2

1.5

1.5

1

1

1

0.5

0.5

0.5

0

0

0

5

10

2

15

0

5

10

MC27

2

2

1.5

1.5 1

1

1

0.5

0.5

0 0

5

10

15

10

15

0

5

10

15

5

10

15

cf; L1527

1.5

0.5 0

5

L483

-5 2.5

15

2.5 NGC1333IRAS4B

15

2.5

1.5

0

10

L1551IRS5

15

2.5

2.5

5

0 -5

15

NGC1333IRAS2

0

2.5

2.5 L1448C

TMC-1A

2

1.5

-5

TMB [ K]

2.5

2.5

L1448N

2

Vol. 697

0 -5

0

5

10

15

-5

0

VLSR [km/s] Figure 4. Line profiles of C4 H (N = 9−8, F1 ) observed with NRO 45 m.

a few times stronger than Mopra, most likely due to beam dilution in the larger Mopra beam. This suggests that the carbonchain molecules, Cn Hm , have a compact distribution around the protostar. The intensities of the C4 H2 lines are comparable to those in L1527, whereas the intensities of the CH3 CCH lines are approximately twice as bright as L1527 (Figure 5). The C4 H line is also stronger than in L1527, as mentioned before. The VLSR values for the lines range from 4.9 to 5.3 km s−1 , consistent with those derived from the N2 H+ and H2 CO lines (Mardones et al. 1997). Line widths are 0.2–0.4 km s−1 , except for those with poor S/N. In the case of C2 H, the lines are broader for the stronger components. This could arise if the brightest component is optically thick. The narrow line widths indicate that these molecules do not exist in the outflow. Note that a typical velocity shift of 5 km s−1 is seen for the outflow component toward IRAS15398−3359 (Tachihara et al. 1996). The gas kinetic temperature is estimated by using the integrated intensities of the K = 1 and K = 2 lines of CH3 CCH, as in the case of L1527 (Sakai et al. 2008b). It is found to be 12.6 ± 1.5 K from the NRO 45 m data, where the error is evaluated from three times the standard deviation of the integrated intensities. The total column density of CH3 CCH is then 13 determined to be (9.9+0.5 cm−2 , assuming local ther−0.2 ) × 10 modynamic equilibrium (LTE) conditions. Using this derived temperature, the integrated intensity of the K = 3 line is esti-

mated to be 0.004 K km s−1 , consistent with the upper limit (3σ ) given in Table 9. The gas kinetic temperature derived above is the effective temperature averaged over the telescope beam along the line of sight. If we use the Mopra data with a larger beam size (36 ), the gas kinetic temperature is derived to be 8.7 ± 1.7 (3σ ) K, implying a higher temperature for the inner region near the protostar. The temperature derived in IRAS15398−3359 is similar to that in L1527 (13.9 K; Sakai et al. 2008b). It is also consistent with the isothermal dust temperature at the 1000 AU scale reported by Shirley et al. (2002). In addition to the above molecules, we also detected a very high excitation (VHE) line of HC5 N (J = 32−31) toward IRAS15398−3359 (Figure 5). This line has an upper state energy of 67.5 K, and hence, it mainly traces warmer gas in the innermost region near the protostar. The intensity of the HC5 N line in IRAS15398−3359 is even higher than that observed in L1527 (N. Sakai et al. 2009, in preparation). If the rotational temperature of 26 K derived for the VHE lines (J = 32– 31, 33–32, and 40–39) of HC5 N seen in L1527 is assumed here, a column density of HC5 N toward IRAS15398−3359 of (2.0 ± 0.4) × 1012 cm−2 is determined, a factor of 2 higher than in L1527. Here the error is estimated from three times the standard deviation of the integrated intensity. On the other hand, if the rotational temperature is assumed to be as low as that for C2 H (7.2 K; Table 7), the column density is found to be

No. 1, 2009

DISCOVERY OF SECOND WCCC SOURCE IN LUPUS

Species C2 H

c-C3 H2 l-C3 H2

Table 9 Line Parameters Observed in IRAS15398−3359 in Lupus 1  TMB a dv a TMB dv (3σ ) −1 (K) (km s ) (K km s−1 )

Transition J J J J J J

= 3/2–1/2, F = 3/2–1/2, F = 3/2–1/2, F = 1/2–1/2, F = 1/2–1/2, F = 1/2–1/2, F 43,2 –42,3

= 1–1 = 2–1 = 1–0 = 1–1 = 0–1 = 1–0

C4 H2

41,3 –31,2 40,3 –30,2 101,10 –91,9 100,10 –90,9 101,9 –91,8

C4 H

N = 9–8, F1

CH3 CCH

N = 9–8, F2 J = 5–4, K = 0 J = 5–4, K = 1 J = 5–4, K = 2

HC5 N

J = 5–4, K = 3 J = 32–31

HCO+2

40,4 –30,3

1.740(14) 3.270(22) 2.499(16) 2.681(18) 2.214(18) 1.741(17) 0.069(13) 0.183(23) 0.204(18) 0.122(20) 0.203(46) 0.071(18) 0.175(35) 0.074(15) 1.367(17) 2.425(35) 1.263(17) 0.661(14) 1.822(25) 0.644(15) 1.769(27) 0.090(20) 0.289(25) ... 0.039d 0.231(22) 0.060d 0.181(20)

0.311(3) 0.426(3) 0.382(3) 0.377(3) 0.325(3) 0.306(3) 0.48(11) 0.61(9) 0.31(3) 0.40(8) 0.16(4) 0.25(7) 0.21(4) 0.37(8) 0.279(4) 0.282(5) 0.275(4) 0.365(9) 0.371(6) 0.357(9) 0.343(6) 0.21(6) 0.37(4) ... ··· d 0.45(5) ··· d 0.36(5)

0.585(9) 1.514(12) 1.043(11) 1.097(12) 0.778(11) 0.572(11) 0.032(14) 0.103(29) 0.068(11) 0.056(17) 0.036(16) 0.020(10) 0.040(17) 0.036(11) 0.404(10) 0.791(21) 0.366(10) 0.269(11) 0.767(19) 0.250(11) 0.715(18) 0.020(8) 0.127(20) 0.018 0.043(19)d 0.112(22) 0.019(13)d 0.067(16)

777

rmsb (mK) 9.8 9.3 9.4 10.6 11.0 11.5 9.4 15.8 12.3 14.3 32.3 12.5 26.7 10.4 12.1 24.6 11.8 10.0 17.2 10.2 17.9 13.0 17.7 ··· 13.2 15.7 11.8 14.6

VLSR a,c (km s−1 )

Telescope

5.25 5.19 5.23 5.16 5.12 5.26 5.16 5.12 5.09 5.03 5.22 4.98 5.19 5.13 5.17 5.19 5.14 4.98 5.00 4.93 4.94 5.03 4.98 ... 5.24d 5.22 5.12d 5.18

Mopra Mopra Mopra Mopra Mopra Mopra Mopra NRO 45 m Mopra Mopra NRO 45 m Mopra NRO 45 m Mopra Mopra NRO 45 m Mopra Mopra NRO 45 m Mopra NRO 45 m Mopra NRO 45 m NRO 45 m Mopra NRO 45 m Mopra NRO 45 m

Notes. The numbers in parentheses represent the errors in units of the last significant digits. a Obtained by a Gaussian fit. b The rms noise averaged over the line width. In the case of marginal detection with Mopra, the line width of the same line observed with NRO 45 m is assumed. c Rest frequency errors (typically an order of ∼0.01 km s−1 ) are much larger than Gaussian fitting errors. d Tentative detection. It cannot be fitted by a Gaussian. T MB represents the peak intensity and VLSR is its velocity.

(5.6±1.1)×1014 cm−2 . This is an order of magnitude higher than in TMC-1 (Takano et al. 1990), and so is likely to be unrealistic, although this should be confirmed with observations of lower J lines of HC5 N. According to the rotational diagram of HC5 N prepared for L1527 (N. Sakai et al. 2009, in preparation), there are at least two components providing the HC5 N emission, with the VHE lines coming mostly from the warmest component. Therefore, we adopt the first estimate for the column density for HC5 N, of (2.0 ± 0.4) × 1012 cm−2 . The beam-averaged column densities for l-C3 H2 , C4 H2 , and C4 H in IRAS15398−3359 are derived under the assumption of LTE at 12.6 K. They are listed in Table 10 along with the column densities of CH3 CCH and HC5 N derived above. The column densities of C4 H and CH3 CCH in IRAS15398−3359 are comparable to those in L1527, whereas the column densities of C4 H2 and l-C3 H2 are lower. However, the column density of l-C3 H2 in IRAS15398−3359 may be affected by beam dilution from the larger Mopra beam. Overall, the column densities of these molecules in IRAS15398−3359 are similar to those in L1527. In order to discuss the fractional abundances of the molecules relative to the H2 molecule, it is necessary to know the total column density of H2 through a comparable beam size (20 ). We estimate this from dust continuum measurements. According to Submillimeter Common-User Bolometric Array instrument (SCUBA) observations with James Clerk Maxwell Telescope

(JCMT) by Shirley et al. (2000), the dust continuum flux at 850 and 450 μm is 2.63 ± 0.15 and 19.6 ± 2.7 mJy, respectively, comparable to L1527 (3.19 ± 0.19 and 18.2 ± 3.2 mJy, respectively). Therefore, the column density of H2 toward IRAS15398−3359 is roughly approximated as the same as toward L1527 (3 × 1022 cm−3 ; Jørgensen et al. 2002). Thus, our conclusion is that the fractional abundances of C4 H and CH3 CCH are also comparable to those in L1527. Certainly, IRAS15398−3359 must be a carbon-chain rich source. 4.3. Distribution of C4 H and C2 H around the Protostar In L1527, carbon-chain molecules are found to exist in an infalling envelope on the basis of the intensity distribution of the C4 H line emission and the systematic broadening of the line toward the protostar’s position. Since IRAS15398−3359 has a face-on disk, it is difficult to assess a change in the line widths in a similar manner. However, it is essential to show that distributions of the carbon-chain molecules are concentrated around the protostar in order to establish IRAS15398−3359 as the second WCCC source. For this purpose, we obtained a five point map of C4 H (N = 9–8) using Mopra. Figure 6 shows the profile map obtained. The grid spacing is 18. 7, equal to that of the map obtained toward L1527 with NRO 45 m. This spacing is close to the half beam width of Mopra. In all the spectra, the intensities are similar. The C4 H

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NRO 45 m

(a) C4H2 101, 9 - 91, 8

0.4 0.3

1

0.1

0.5

0

0

-0.1 0.25

1

2

3

4

5

6

7

8

Mopra

(b) C4H2 100, 9 - 90, 8 0.15 0.2

2

0.5

0

0 1

2

3

4

5

6

7

8

Mopra

(c) l-C3H2 0.4 41, 3-31, 2 0.3

1

0.4 0.2

0

0

-0.1

-0.2

2

3

4

5

0.3

6

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8

NRO 45 m

(d) c-C3H2 43, 2-32, 3

0.4

0.5

0.2 0.1

0

0

-0.1

-0.1 3

1

2

1

6

7

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NRO 45 m

3

4

5

6

7

8

NRO 45 m

4

5

6

7

8

2

3

4

5

1

2

3

6

7

8

NRO 45 m

(h) HC5N J=32-31

0.3

0.1

2

5

(f) CH3CCH J=5-4, K=1

0.4

0.2

1

4

0.6

0.1

1

3

(g) CH3CCH 0.8 J=5-4, K=2

0.2

0.5

2

1

0.05

0.5

1

1.5

0.1

-0.05

NRO 45 m

(e) CH3CCH J=5-4, K=0

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VLSR [km/s] Figure 5. Spectral line profiles observed toward IRAS15398−3359.

line emission appears extended over the beam size of Mopra (∼36 ), and no apparent central concentration can be seen. In addition, no velocity shift from the systematic velocity of the IRAS15398−3359 core is seen, and no significant broadening is recognized. Considering that the C4 H (N = 9–8) line can be observed even in a starless core TMC-1 with a moderate intensity (1.2 ± 0.2 K; Sakai et al. 2008b), these spectra may reflect the ambient component as well as the emission from the WCCC region. Hence, we also conducted small mapping observations of the C2 H (N = 4–3, J = 9/2–7/2) line using the ASTE 10 m telescope. Since the critical density and the upper state energy of the C2 H line are ∼106 cm−3 and 41.9 K, respectively, this line traces warm and dense regions. Figure 7 shows a profile map of the C2 H line. The observing grid is 11 , the half beam size (22 ) of ASTE 10 m telescope at 340 GHz. Each spectrum shows a small splitting due to the hyperfine interaction of the H nucleus, and hence, the line looks broader than the other lines. Figure 7 indicates that the C2 H emission shows an apparent central concentration toward the protostar’s position, with a size of 20 –30 (3000–5000 AU). This result further supports IRAS15398−3359 being the second WCCC source. 4.4. Detection of HCO+2 toward IRAS15398−3359 We recently detected rotational emission lines of HCO+2 toward L1527, the first such detection in a star-forming region outside the Galactic center region (Sakai et al. 2008a). On the basis of this result, we suggested that CO2 production may occur

Table 10 Column Densities of Various Carbon-chain Molecules in IRAS15398−3359 Species CH3 CCH l-C3 H2 C4 H2 C4 H HC5 N (VHE)

IRAS15398

Telescope

L1527a

13b (9.9+0.5 −0.2 ) × 10 (3.3 ± 0.7) × 1011c (3.7 ± 1.6) × 1011c (1.4 ± 0.4) × 1014c (2.0 ± 0.4) × 1012e

NRO 45 m Mopra NRO 45 m NRO 45 m NRO 45 m

(6.0 ± 0.3) × 1013b (1.1 ± 0.4) × 1012d (1.6 ± 0.3) × 1012 (1.9 ± 0.4) × 1014d (1.3 ± 3.0) × 1012

Notes. Unit is cm−2 . a Taken from Sakai et al. (2008b). b Evaluated from the K = 1 and K = 2 lines. Errors are derived from three times the standard deviation of their integrated intensities. c The rotational temperature is assumed to be 12.6 K. Errors are derived from three times the rms noise of the integrated intensities considering the uncertainty in the excitation temperature (±1.5 K just as the 3σ errors for the CH3 CCH case). d The rotational temperature is assumed to be 12.3 K. Errors are derived from three times the rms noise of the integrated intensities considering the uncertainty in the excitation temperature (±2.3 K just as the 3σ errors for the C4 H2 case). e Derived from the VHE line, J = 32–31. The rotational temperature is assumed to be 26 K. See the text for details.

in the gas phase by oxidation of carbon-chain molecules in the WCCC region. Very interestingly, the HCO+2 (404 –303 ) line was also detected in IRAS15398−3359, with Mopra and NRO 45 m. The line profiles are shown in Figure 8. If we assume that the excitation temperature is 12.6 K, as before, the beam-averaged column density is evaluated to be (2.7 ± 0.6) × 1011 cm−2 from the NRO 45 m data. The error

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2

N TMB [K]

1.5

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1 0.5

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0 2

-0.5 1 2

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2

3

4

5

6

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Figure 6. Profile map of C4 H (N = 9−8, F1 ) observed toward Lupus 1 with Mopra. Grid spacing is 18. 7. The center position is toward IRAS15398−3359, and the beam size is 36 .

is evaluated from three times the standard deviation of the integrated intensity, considering the uncertainty in the excitation temperature (±1.5 K) just as for the CH3 CCH case. The column density is larger than the corresponding value derived in L1527 by a factor of 3.6. This might be because the column density in L1527 has been derived from IRAM 30 m observations, which have a larger beam size (29 ) than NRO 45 m (20 ). The fractional abundance of gaseous CO2 is estimated from this column density by the method described in Sakai et al. (2008a). This method employs a simplified chemical model to relate the fractional abundance of CO2 with the beamaveraged column density of HCO+2 under the assumption of a spherical distribution of CO2 . The fractional abundance of CO2 is proportional to the HCO+2 number density and to the square of the fractional abundance of CO. According to Pontoppidan et al. (2008), there is a large amount of solid CO2 in IRAS15398−3359. When we extrapolate the isothermal dust temperature distribution given by Shirley et al. (2002) to the inner region, it is higher than the evaporation temperature of solid CO2 (∼50 K) within a radius of 26 AU. In this case, the fractional abundance of gaseous CO2 is unrealistically high, 2.1 × 10−1 , exceeding the elemental abundance of carbon. Here we assume that fCO is the same as for L1527 (3.9 × 10−5 ; Jørgensen et al. 2002). Even if the radius is assumed to be 140 AU, that is the CO2 evaporation radius for L1527, the fractional abundance is still found to be 1.3 × 10−3 . Therefore, the gaseous CO2 cannot originate solely from the evaporation of grain mantles in a small inner warm region. On the other hand, if the radius is comparable to the beam size of NRO 45 m (∼1300 AU), the abundance is found to be 4.1 × 10−6 . Thus, the gaseous CO2 has to be abundant in low temperature regions, as far as the simplified model assumed by Sakai et al. (2008a) is correct. In this case, its presence could be due to WCCC or nonthermal desorption.

The HCO+2 lines have not been detected toward other representative star-forming regions such as Orion KL (e.g., Minh et al. 1988; Lee et al. 2001). In spite of the sensitive observations with Mopra, the HCO+2 line was not detected toward IRAS16293−2422, a well studied low-mass star protostar with high abundances of other O-bearing molecules like CH3 OH (e.g., Maret et al. 2005). The upper limit to the integrated intensity is 0.025 K km s−1 (3σ ), corresponding to a beamaveraged fractional abundance of HCO+2 of 7.7 × 10−14 . This upper limit is much lower than the HCO+2 fractional abundances in IRAS15398−3359 and L1527 (9 × 10−12 and 2 × 10−12 , respectively). Therefore, the detection of HCO+2 in the two WCCC sources, IRAS15398−3359 and L1527, implies that the gas-phase formation of CO2 is related to WCCC. The high abundance of solid CO2 reported by Pontoppidan et al. (2008) might originate from depletion of gaseous CO2 formed through WCCC. 5. WCCC As described above, various carbon-chain molecules are found to be abundant in L1527 and IRAS15398−3359. The observed features of these sources are summarized as follows. 1. High excitation conditions: a number of high excitation lines of carbon-chain molecules are detected. In particular, the C4 H N = 9–8 F1 line is as intense as 1.7 K and 2.4 K in TMB in L1527 and IRAS15398−3359, respectively. The beam-averaged gas kinetic temperatures are derived to be 13.9 ± 1.3 K and 12.6 ± 1.5 K, for L1527 and IRAS15398−3359, respectively, which are higher than typical values in cold starless cores, 10 K. Furthermore, the carbon-chain molecules must exist in warm dense gas near the protostar, because the critical density of the C4 H2 lines is ∼106 cm−3 .

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Figure 7. Profile map of C2 H (N = 4−3, J = 9/2−7/2, F = 5−4, 4−3) observed toward IRAS15398−3359 with ASTE. VLSR is calculated by using the F = 5−4 transition frequency (349.337741 GHz). Grid spacing is 11 . The mapping center is offset from IRAS15398−3359 by Δα = −3 .

TMB [K]

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NRO 45 m

0.2 0.05 0.1 0

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15

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VLSR [km/s] Figure 8. Line profiles of HCO+2 (404 − 303 ) observed toward IRAS15398−3359 in Lupus 1.

2. Central concentration: the distribution of the carbon-chain molecules is centrally peaked, when observed with the high excitation lines. In particular, carbon-chain molecules are

found to exist in an infalling envelope in L1527 on the basis of the line broadening toward the protostar position. In IRAS15398−3359, the C2 H (N = 4–3, Eu = 41.9 K)

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DISCOVERY OF SECOND WCCC SOURCE IN LUPUS

line emission shows a compact distribution around the protostar. 3. Enhancement of gaseous CO2 : spectral lines of HCO+2 are detected in L1527 and IRAS15398−3359. This implies that substantial amount of gaseous CO2 exists in a relatively low-temperature region where thermal evaporation of solid CO2 from grain mantles is inefficient. Gaseous CO2 may be produced by oxidation of carbon-chain molecules in the WCCC sources. 4. Existence of long carbon-chain molecules including anions: in addition to the above three points, an additional important characteristic is inferred from the L1527 result. Spectral lines of various long carbon-chain molecules such as C5 H, C6 H, C6 H2 , HC7 N, and HC9 N are seen in L1527 (Sakai et al. 2008b). Moreover, lines of the molecular anions C4 H− and C6 H− are seen as well (Sakai et al. 2007, 2008d; Ag´undez et al. 2008). These molecules have not been detected in other star-forming regions. All the above features are surprising in a conventional view of carbon-chain chemistry, because the carbon-chain molecules have been thought to be abundant only in the early stage of the starless core phase. A mechanism accounting for all the results could be the regeneration of carbon-chain molecules in a lukewarm region near the protostar. This regeneration mechanism is called WCCC (Sakai et al. 2008b), and is briefly summarized as follows. WCCC is triggered by evaporation of CH4 from grain mantles. The evaporation temperature of CH4 is ∼30 K. This is higher than that of CO (∼20 K), but is significantly lower than that of H2 O (∼100 K). Hence, CH4 can be abundant in a warm region which is somewhat extended around the protostar. Gaseous CH4 evaporated from grain mantles would then react with C+ to form the precursor ions for carbon-chain molecules as CH4 + C+ → C2 H+3 + H, (1) and

CH4 + C+ → C2 H+2 + H2 .

(2)

These ions then produce the shortest carbon-chain molecules, C2 H and C2 H2 , through dissociative electron recombination reactions: C2 H+3 + e → C2 H + H + H, (3)

and

C2 H+3 + e → C2 H2 + H,

(4)

C2 H+2 + e → C2 H + H.

(5)

+

Acetylene, C2 H2 , further reacts with C to form longer carbonchains. A condensation reaction between C2 H2 and C2 H+3 produces C4 H+3 , which leads to C4 H by a dissociative electron recombination reaction. A neutral–neutral reaction between C2 H2 and CN makes HC3 N. Through these reactions, various carbon-chain molecules are formed efficiently in a dense region when heated by an emerging protostar. This is a completely different mechanism to the classical carbon-chain chemistry occurring in cold starless clouds, which provides CH, CH2 , CH3 , and CH4 through hydrogenation reactions from C+ or C. Then the carbon chains start to grow through reactions of these hydrocarbons with C and C+ . Since C+ and C are abundant in the early stage of starless cores, before C is fixed into CO, the carbon-chain molecules can be efficiently produced. They later become deficient through

781

gas-phase destruction and depletion onto dust grains. In other words, a carbon-rich condition is temporally realized in the early stage of evolution of the core. In contrast, in WCCC it is evaporation of CH4 from the grain mantles that makes for the carbon-rich conditions in the warm, dense gas near the protostar, providing favorable conditions for the production of carbonchain molecules. This mechanism has been verified through chemical model simulations by Aikawa et al. (2008) and Hassel et al. (2008). Since WCCC makes substantial amount of carbon-chain molecules, the oxidation reactions of carbon-chain molecules in the gas phase such as C2 H + O2 → CH + CO2 ,

(6)

can then effectively occur to form gaseous CO2 . Solid O2 sublimates at temperatures higher than 19 K, and hence, O2 is expected to be abundant in the WCCC region. Thus, WCCC can play an important role in production of gaseous CO2 (Aikawa et al. 2008; Hassel et al. 2008). 6. WCCC AND HOT CORINO CHEMISTRY 6.1. Source-to-Source Variation of Chemical Components As we have described, we conducted observations with Mopra and NRO 45 m toward 16 low-mass star-forming regions, and discovered a new apparent WCCC source, IRAS15398−3359. In order to characterize the other sources, the column density and the fractional abundance of C4 H are derived for each source. The method of derivation is the same as for IRAS15398−3359, with LTE at 12.3 K assumed. The results are summarized in Table 11, together with H2 column densities taken from the literature. To show the sensitivity of these results to our assumption, the column density of C4 H toward B1, for instance, is calculated for excitation temperatures of 10, 15, and 20 K to be 10.6 × 1013 , 7.5 × 1013 , and 6.8 × 1013 cm−2 , respectively. This implies that our estimate of the C4 H column density is accurate to a factor of 2. On the other hand, the fractional abundances are less accurate because of the uncertainties of the available H2 column densities, perhaps as much as a factor of 10. Figure 9 shows a plot of the column densities of C4 H and H2 for the observed sources. Since some sources, including L1527 and IRAS15398−3359, were observed with both Mopra and NRO 45 m, the column densities obtained are plotted independently. As can be seen, the fractional abundance of C4 H is higher in L1527 and IRAS15398−3359 than those in other sources. This trend would remain, even if uncertainties of the H2 column density are taken into account. At the same time, several sources like B1, MC27, TMC-1A, L1448C, BHR71IRS1, Serpens SMM4, L1448N, and L483 also show relatively high abundances. Hence, WCCC may be occurring in these sources to some extent, and not just be confined to L1527 and IRAS15398−3359. However, it is not easy to distinguish the WCCC contribution from the cold extended component only by C4 H observations, as evident from our mapping toward IRAS15398−3359 (Figure 6). For example, the C4 H line detected toward IRAS16293−2422 is apparently narrow (0.58–0.68 km s−1 ) in comparison with the other lines coming from the vicinity of the protostar (e.g., Cazaux et al. 2003; typically a few km s−1 ), and hence, it may trace an ambient extended component. Further observations of lines with high critical densities are necessary to investigate the contribution of WCCC in other sources. Particularly, B1 and

782

SAKAI ET AL. Table 11 Column Densities and Fractional Abundances of C4 H N (1013 cm−2 ) f (10−10 )a Telescope References for N(H2 )

Source L1448N L1448C NGC1333IRAS2A NGC1333IRAS4B B1 HH211-mm MC27 TMC-1A L1551IRS5 BHR71 IRS1 IRAS13036−7644 IRAS15398−3359 IRAS16293−2422 L483 Serpens SMM4 RCrA IRAS7B L1527

7.3 ± 1.7 4.7 ± 1.6 3.1 ± 1.3 1.9 8.5 ± 1.6 9.7 ± 1.1 9.7 ± 1.4 2.3 ± 1.4 3.6 ± 1.2 4.9 ± 0.5 2.1 ± 0.4 6.6 ± 0.2 13.9 ± 0.4 1.2 ± 0.3 4.2 ± 0.3 7.5 ± 0.7 2.5 ± 0.6 0.7 8.8 ± 0.5 18.9 ± 0.5

20 3 0.6 0.6 30 ... 7 7 0.9 1 ... 20 50 0.08 0.5 0.8 8 0.2 30 60

NRO 45 m NRO 45 m NRO 45 m NRO 45 m NRO 45 m NRO 45 m NRO 45 m NRO 45 m NRO 45 m Mopra Mopra Mopra NRO 45 m Mopra Mopra NRO 45 m Mopra Mopra Mopra NRO 45 m

... 1 1 1 2 ... 3 1 1 4 ... ... ... 5 1 1 ... 6 1 1

Notes. The rotational temperature is assumed to be 12.3 K except for the IRAS15398−3359 case (12.6 K). N(H2 ) in IRAS15398−3359 is assumed to be equal to that in L1527, as described in the text. The errors are derived from the 3σ values of integrated intensities. a Fractional abundance of C H relative to H . 4 2 References. (1) Jørgensen et al. 2002; (2) Hirano et al. 1997; (3) Crapsi et al. 2005; (4) Chen et al. 2008; (5) Sch¨oier et al. 2002; (6) Sch¨oier et al. 2006.

Mopra NRO 45 m L1527 IRAS15398 B1 IRAS15398

N(C4H) [cm-2]

10

14

MC27

L1527 L1448N

L483 L1448C

Serp.SMM4

f

10

TMC-1A

BHR71 L483 L1551 NGC1333I2A NGC1333I4B

=10-9

IRAS16293

13

RCrAIRAS7B f =10-10

10

22

Vol. 697 −1

0.012 K km s ). The upper limit to the column density of C4 H2 is determined to be 1.2×1011 cm−2 , which corresponds to an upper limit to the fractional abundance of 3.9×10−13 . This is much lower than the values in L1527 and IRAS15398−3359, 5.3 × 10−11 and 1.2×10−11 , respectively. Similarly, the C4 H2 (101,9 − 91,8 ) line is not detected toward L1448N with an upper limit of 0.019 K km s−1 , corresponding to a column density limit of 1.9 × 1011 cm−2 . These results indicate that there is a significant variation in the C4 H and C4 H2 abundances among star-forming regions, although the observed sources are all Class 0 objects or transient objects from Class 0 to Class I. So far, such a chemical variation between star-forming regions of similar character has not been well recognized. This variation may be related to a variation in the physical processes of star formation, as we discuss in the following section. Interestingly, most of these sources without the WCCC features are so-called hot corino sources (e.g., Cazaux et al. 2003; Bottinelli et al. 2004a; Sakai et al. 2006). These are characterized by the existence of complex organic molecules like HCOOCH3 , (CH3 )2 O, and C2 H5 CN. These molecules are formed through gas phase reactions, evaporated from grain mantles in a hot region (T > 100 K) near the protostar, or they are directly formed from CH3 OH and H2 CO in grain mantles during a warm up phase and then are released into the gas phase (e.g., Garrod & Herbst 2006). Conversely, the complex organic molecule, HCOOCH3 , has not been detected in one of the WCCC sources, L1527, in spite of sensitive observations (Sakai et al. 2008c). In addition to the WCCC sources and the hot corino sources, there are intermediate sources in view of the variation of the C4 H abundance we have measured (Figure 9). Therefore, the WCCC activity can be varied from source to source, the highest being for the WCCC sources and the lowest for the hot corino sources. It is not clear whether the intermediate sources in view of the C4 H abundance are also intermediate sources in view of the hot corino chemistry, because the spectral lines of complex organic molecules have not been observed. Nevertheless, we can at least say that clear hot corino activity cannot be seen in the WCCC source, L1527. Hence, it seems likely that the WCCC activity and the hot corino activity can co-exist in general, with their relative importance varying from source to source. In this scenario, the WCCC sources and the hot corino sources represent the two extreme cases for the chemical composition of low-mass star-forming regions.

f =10-11

10

23

10

24

N(H2) [cm-2] Figure 9. Plot of the column densities of C4 H and H2 . Circles represent the data observed by Mopra, and triangles represent the data observed by NRO 45 m. The fractional abundances of 10−9 , 10−10 , and 10−11 are shown by the dashed lines.

L1448N are interesting sources for such studies, according to Figure 9. In contrast, the abundance of C4 H is very low in several sources like NGC1333 IRAS4B, IRAS16293−2422, and RCrA IRAS7B. In particular, the C4 H (N = 9–8) and C4 H2 (101,9 –91,8 ) lines are not detected toward NGC1333 IRAS4B in spite of sensitive observations with NRO 45 m (the 3σ noise level of the integrated intensity of the C4 H2 line is

6.2. Origin of the Chemical Variation What determines the relative importance between WCCC and hot corino chemistry? An important factor is apparently the chemical composition of the grain mantles, because WCCC and hot corino chemistry are both triggered by evaporation of molecules from the mantles. If CH4 is abundant there, in comparison with complex organic molecules, WCCC would naturally occur. If CH4 is deficient for some reason, and the complex organic molecules are abundant, hot corino chemistry would dominate. The chemical composition of the grain mantles is mostly determined in the starless-core phase, when atoms and molecules are depleted onto dust grains. In this scenario, the chemistry occurring after the onset of star formation would reflect the history of the starless-core phase, memorized in the grain mantles.

No. 1, 2009

DISCOVERY OF SECOND WCCC SOURCE IN LUPUS CO

C

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CO

H CO

(a) Long

783

CO

C

depleted as CO

CH 3OH

CO

CO

H

Hot Corino Chemistry

H

CH 4

H

H 2CO

CH 3OH

Complex organic molecules (HCOOCH3, (CH3)2O, etc.)

C CO H H (b) Short

C

C H

H C

depleted as C

CO

C

H C

CH 4 CH 4

WCCC CH 4

Carbon-chain molecules

CH 3OH

Figure 10. Our proposed scenario to account for the chemical variation which we observe in low-mass star-forming cores. Sketch (a) represents the long starless-core phase which favors hot corino chemistry. In this case, most of the C atoms are converted to the CO molecules in the gas phase, which then deplete onto dust grains. They then form CH3 OH, H2 CO, and more complex organic molecules through hydrogenation reactions and radical reactions on the grain surfaces, before being evaporated into the gas phase. Complex organic molecules such as HCOOCH3 are also produced from CH3 OH and H2 CO in the gas phase after their evaporation. Sketch (b) represents a short starless-core phase in which most of the C atoms deplete directly onto the dust grains, and then form CH4 through hydrogenation. This favors WCCC, with the long carbon-chain molecules built up in the gas phase after evaporation of CH4 from dust grains.

One possible explanation for the source-to-source variations in the chemical composition is the timescale of the starlesscore phase, which would be related to the timescale for the protostellar collapse. The major form of carbon changes from C to CO through various gas-phase chemical reactions in the starless cores, with a timescale of roughly 106 yr. This is comparable to, or longer than, the free-fall time (4 × 105 yr). Thus, the C atom can still be abundant in the late stage of a starless core, if the prestellar collapse time is close to that of the free-fall time. The timescale for the depletion of atoms and molecules onto dust grains is roughly 105 /(n/104 cm−3 ) yr (e.g., Burke & Hollenbach 1983). Hence, the C atom could be depleted onto dust grains before it is fixed into CO. In this case, CH4 will be efficiently produced in grain mantles by hydrogenation reactions of the C atom, and then will drive WCCC after the onset of star formation. In addition, some of the carbon-chain molecules produced in the early evolutionary stage of starless cores can still survive after the onset of star formation in the fast collapse case. This would explain a continuous distribution of carbon-chain molecules from the outer part of the infalling envelope to the inner portion. On the other hand, the situation can be different if the timescale of the starless-core phase is much longer than that of the free-fall time. In this slow collapse case, the C atom is mostly converted into CO in the gas phase, with the CO then being depleted onto dust grains. Therefore, few C atoms fall onto dust grains. Instead, CH3 OH and H2 CO are produced on dust grains by hydrogeneration of CO, whereas little CH4 is produced due to the deficiency of C atoms. The resulting CH3 OH and H2 CO molecules will then contribute to the production of more complex organic molecules like HCOOCH3 . After the onset of star formation, these molecules will be evaporated from the grain mantles to the gas phase, so driving hot corino chemistry. These two scenarios are schematically shown in Figure 10. In this picture, the relative weight of WCCC and hot corino chemistry is determined by the timescale for the starless-core phase. Therefore, the existence of sources with an intermediate chemical state between WCCC and hot corino chemistry can naturally explained. The above picture is consistent with the gas-grain model of gravitationally contracting clouds by Aikawa et al. (2005).

They compared the chemical compositions during the starless phase for the fast and slow collapse cases. According to their calculations, solid CH4 is more abundant for the fast collapse case than for the slow collapse case by an order of magnitude. On the other hand, solid CH3 OH is more abundant for the slow collapse case especially in the inner region. Although their model only treats the starless-core phase before WCCC occurs, this outcome supports the mechanism proposed above for the origin of the chemical variation. Solid CH4 in grain mantles has recently been observed toward several low-mass star-forming regions with the Spitzer Space ¨ Telescope. Oberg et al. (2008) surveyed the 7.7 μm absorption of CH4 toward 25 sources. L1527 was not included because the source is faint at 7.7 μm. On the other hand, IRAS15398−3359 was found to have strong CH4 absorption in comparison with the other sources. This observation supports the above scenario for WCCC. The timescale for the starless-core phase remains a problem for star-formation models, which needs to be addressed through observations. It is estimated to be ∼105 –106 yr on the basis of statistical arguments (e.g., Onishi et al. 1998). However, it is almost impossible to measure the timescale in an individual case. Hence, it is not clear whether there are source-to-source variations in the lifetime of starless cores. In principle, the timescale of the prestellar collapse can vary from source to source, or from region to region, depending on what controls the mechanism of cloud support, i.e., the strength of the magnetic fields, initial turbulence, etc. If the relative importance of WCCC and hot corino chemistry does reflect the collapse timescale, as suggested by the above scenarios, it will provide a new tool to understand and probe the processes occurring at the initiation of star formation. In this respect, it is worth noting that the observed line widths of the two WCCC objects are narrow, implying weak turbulent support against gravitational collapse. This is particularly true for the face-on source, IRAS15398−3359 (∼0.3 km s−1 ) in comparison with the line widths of C2 H and C4 H in other sources (Tables 3, 4, and 9). Another interesting fact is that the Taurus and Lupus regions, where the WCCC sources are found, are active star-forming regions. Table 12 shows the fraction of cores in which star formation occurs in each region.

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SAKAI ET AL. Table 12 Fraction of Star-forming Cores

Cloud Taurus Ophiuchus Southern Coal Sack RCrA Chameleon Lupus (Lupus 1

Number of Number of Fraction Dense Cores Star-forming Cores (%) Reference 40 40 6 8 23 36 12

19 9 0 2 6 6 4

48 23 0 25 26 17 33

1, 2 2 3 4 5 6 6)

References. (1) Onishi et al. 1998; (2) Tachihara et al. 2000; (3) Kato et al. 1999; (4) Yonekura et al. 1999; (5) Mizuno et al. 1999; (6) Hara et al. 1999.

Cores with protostellar IRAS sources and T-Tauri stars are regarded as star-forming cores (e.g., Tachihara et al. 2000). This fraction will reflect the timescale of the starless-core phase, a high fraction implying a short lifetime. Apparently, the Taurus region shows a higher fraction than the other regions. Although the Lupus region shows a relatively low fraction as a whole, the Lupus 1 cloud, where IRAS15398−3359 exists, has a higher fraction. These results are consistent with the short starless-core phase (or possibly the fast prestellar collapse) suggested for the WCCC sources. Note that the difference of the timescale of the starless-core phase has recently been pointed out by Jørgensen et al. (2008) between Ophiuchus and Pereseus regions, and by Enoch et al. (2008) among the Ophiuchus, Serpens, and Pereseus regions, on the basis of the dust continuum observations. Nevertheless, the above scenarios must be tested by further observational and theoretical investigations. A more systematic survey of the chemical compositions of star-forming regions using the lines of carbon-chain molecules and complex organic molecules is necessary to establish a possible anticorrelation between these two kinds of species, as we have hypothesized for WCCC and hot corino chemistry. Furthermore, chemical models simulating dynamically evolving cores under various conditions are also important to understand whether the evolution is indeed reflected in the chemical differences among star-forming regions. 7. SUMMARY 1. Observations of C2 H and C4 H line emission have been conducted toward 16 representative low-mass star-forming regions with NRO 45 m and Mopra in order to search for the WCCC source candidates. As a result, the Class 0 protostar, IRAS15398−3359 in the Lupus molecular cloud complex, is found to be a likely candidate. 2. In IRAS15398−3359, high excitation lines of various carbon-chain molecules, including C4 H, l-C3 H2 , C4 H2 , and HC5 N, are detected. The intensity of the C4 H line is particularly strong, 2.4 K. Considering the critical densities for the observed lines (∼106 cm−3 for C4 H2 ) and the gas kinetic temperature derived from the CH3 CCH lines (12.6 K), we conclude that these carbon-chain molecules exist in a region of warm and dense gas near to the protostar. In fact, the distribution of the high excitation line of C2 H clearly shows a concentration around the protostar with a size of about 20 (∼3000 AU). 3. All these observed features are similar to those found in the known WCCC source L1527. The peak temperature

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of the C4 H line is even higher in IRAS15398−3359 than in L1527. Moreover, the HCO+2 line, which was recently detected toward L1527 as the first example in star-forming regions, is also detected in IRAS15398−3359. Therefore, IRAS15398−3359 is unambiguously identified as the second WCCC source. This finding indicates that WCCC is no longer specific to L1527. 4. In our study, we have recognized that some sources like IRAS16293−2422 and NCG1333 IRAS4B show almost no indications of WCCC. These sources are so called hot corino sources. On the other hand, we have also found other sources showing intermediate characteristics between the WCCC and hot corino sources, in view of their C4 H abundances. The WCCC activity may be occurring, to some extent, in these latter sources. The chemical composition of low-mass star-forming cores is thus found to show a significant source-to-source variation. Characterization of the intermediate sources is left for future studies. 5. Since WCCC and the hot corino chemistry are triggered by the evaporation of CH4 and oxygen-bearing organic molecules from grain mantles, respectively, we suggest that the chemical variation may originate from the difference in chemical composition of grain mantles, which would reflect the chemical processes in the starless-core phase. Thus, an important parameter in determining this is the timescale for the starless-core phase. If it is close to the free-fall time, then C atom is depleted onto dust grains, resulting in CH4 by hydrogenation. If it is much longer than the free-fall time, CO is mainly depleted onto dust grains, giving rise to CH3 OH and H2 CO. The two cases lead to WCCC and hot corino chemistry, respectively. If this is the case, then the chemical composition in star-forming regions provides a diagnostic of their past history during their prestellar phases. We are grateful to Eric Herbst, Cecilia Ceccarelli, Emmanuel Caux, Jose Cernicharo, Nagayoshi Ohashi, and Yuri Aikawa for valuable discussions. We thank Osamu Saruwatari for his assistance in the ASTE observation. This study is supported by Grant-in-Aids from Ministry of Education, Culture, Sports, Science, and Technologies (14204013, 15071201, and 19-6825). APPENDIX A SOURCE DESCRIPTIONS A.1. L1448N, L1448C The L1448 cloud is a member of the Perseus molecular cloud complex, and includes four or more Class 0 objects (L1448C, L1448N(A), L1448N(B), and L1448IRS2). Among them, L1448C is the oldest, located in the south part of the cloud. It is believed that its powerful outflow could have triggered the formation of a binary protostar, L1448N(A) and L1448N(B) (e.g., Barsony et al. 1998). L1448N(B) is the brightest farinfrared source in the cloud (Bachiller & Cernicharo 1986). L1448C and L1448N are selected as good examples of outflowtriggered star-forming regions. A.2. NGC1333 IRAS2A, IRAS4B The NGC1333 cloud is also a member of the Perseus molecular cloud complex, located about 1◦ northeast of the L1448

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DISCOVERY OF SECOND WCCC SOURCE IN LUPUS

cloud. NGC1333 IRAS2 has two millimeter-wave continuum sources separated by ∼30 , each of which is associated with a molecular outflow (Jennings et al. 1987; Blake 1997). Source A (IRAS2A), the stronger of the two sources, shows strong emission from high excitation lines of CH3 OH (Maret et al. 2005). Moreover, the complex organic molecule, (CH3 )2 O, is also detected there (Jørgensen et al. 2005). NGC1333 IRAS4 contains three Class 0 protostars, IRAS4A1, IRAS4A2, and IRAS4B. IRAS4A and IRAS4B (e.g., Blake et al. 1995; Looney et al. 2000) are separated by 30 . According to the dynamical age of the molecular outflow, IRAS4B is considered to be in a younger evolutionary stage than IRAS4A (Choi 2001). Toward these sources, emission from HCOOCH3 is detected (Sakai et al. 2006; Bottinelli et al. 2004a, 2007). NGC1333 IRAS4B and NGC1333 IRAS2A are selected as typical hot corino sources which harbor complex organic molecules. A.3. B1, HH211-mm B1 is a dense core in the Perseus molecular cloud complex, which is associated with a Class 0 protostar IRAS03301+3057. Strong SiO emission is detected toward B1 (Bachiller et al. 1990; Yamamoto et al. 1992; Hirano et al. 1997), suggesting that the outflow in B1 is at an early evolutionary stage (Bachiller & G´omez-Gonz´olez 1992). Multiple deuterated molecules like ND3 and D2 CS are detected in the vicinity of B1, indicating heavy depletion of molecules onto dust grains (e.g., Lis et al. 2002; Vastel et al. 2003). HH211-mm is a well known star-forming region in the IC348 cloud in Perseus, having an archetypical outflow with a highly collimated jet powered by a young, low-mass Class 0 protostar (e.g., McCaughrean et al. 1994). The beautiful outflow features in the CO and SiO emissions are reported by Lee et al. (2007). By including these two sources, our source list covers most of the star-forming regions in Perseus.

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et al. 2006). However, molecular line observations are limited to fundamental species. IRAS13036−7644 is also a protostar formed in a Bok grobule (Hartley et al. 1986), located at the eastern part of the Chameleon II dark cloud complex. This source drives a bipolar molecular outflow detected in the 12 CO (J = 1–0) emission. Infalling signatures are seen in the CS (J = 2–1) and H2 CO (212 –111 ) lines (Lehtinen 1997). IRAS13036−7644 is also classified as a Class 0 to Class I protostar (Lehtinen et al. 2005). BHR71 IRS1 and IRAS13036−7644 are included in the source list as typical star-forming regions in Bok globules. A.6. IRAS16293−2422 IRAS16293−2422 is located in the ρ-Ophiuchus cloud complex. High angular resolution observations revealed that IRAS16293−2422 is a protobinary system with a projected separation of ∼5.2 (∼840 AU). This source harbors rich complex organic molecules like HCOOCH3 , (CH3 )2 O, and C2 H5 CN (Cazaux et al. 2003). According to interferometric observations (e.g., Kuan et al. 2004; Bottinelli et al. 2004b), these complex organic molecules are associated with both components of the binary system. It is interesting to examine whether WCCC occurs in this source. Hence, we include it in the target list. A.7. L483 L483-mm seems to be located in a flattened filament of Aquila Rift (e.g., Shirley et al. 2000). L483 shows gravitational infall motions in the H2 CO (212 –111 ) line (Myers et al. 1995), and is classified as a Class 0 to Class I source (Tafalla et al. 2000). Strong NH3 emission is observed toward this source (Fuller & Wootten 2000). Since the HC5 N (J = 17–16, 45 GHz) emission is relatively bright, this source could be a candidate for the WCCC source. Hence this source is included in the target list. A.8. Serpens SMM4

A.4. MC27, TMC-1A, L1551IRS5 In order to explore whether WCCC is specific to the Taurus region or not, a few sources are chosen from this region. TMC1A is a Class I protostar (Brown & Chandler 1999) in the HCL2 region. Since L1527 and TMC-1 are also in this region, TMC1A might have a similar chemical condition. MC27 (L1521F) has long been recognized as a starless core since its discovery by Onishi et al. (1999). Its HCO+ spectrum shows a self-absorption feature due to infall motion, and hence, MC27 was supposed to be in the late stage of prestellar collapse (e.g., Crapsi et al. 2004). Recently, a very faint infrared source was discovered with the SST, and is now recognized as a possible Class 0 source (Bourke et al. 2006). On the other hand, L1551IRS5 is a noted protostar, where the bipolar outflow phenomenon was first recognized (Snell et al. 1980). It is known to be a binary protostar consisting of Class 0 and Class I objects. A.5. BHR71 IRS1, IRAS13036−7644 The dark cloud BHR71 is an isolated Bok globule, located near the southern Coalsack region. BHR71 IRS1 has a highly collimated outflow, lying almost in the plane of the sky (e.g., Bourke et al. 1997; Parise et al. 2006). The driving source, IRS1 (IRAS11590−6452), is classified as a Class 0 to Class I protostar (Froebrich 2005). The outflow strongly interacts with an ambient cloud and produces a shocked region, where various molecules like SiO and CH3 OH are detected (Parise

The Serpens core is known to involve about a half-dozen Class 0 protostellar candidates within a size of ∼0.25 pc (e.g., Casali et al. 1993; Hurt & Barsony 1996). Among them, SMM4 is the brightest submillimeter continuum object in the southeastern part of the Serpens core (Casali et al. 1993), and shows the strongest H2 CO (303 –202 ) emission among all of the observed sources in Serpens (Hurt & Barsony 1996). SMM4 accompanies a CO outflow (White et al. 1995). A.9. RCrA IRAS7B The R Coronae Australis dark cloud is one of the nearest active star-forming regions. Among many submillimeter sources in the cloud, IRS7B is the only known Class 0 protostar (Nutter et al. 2005). Toward this source, high excitation lines of H2 CO, CH3 OH, and CS are strongly detected, as in the case of IRAS16293−2422 (Sch¨oier et al. 2006). Furthermore, the C2 H (N = 4–3) line emissions are also as strong as 4 K in TMB , and hence, it is included in the target list. REFERENCES Ag´undez, M., Cernicharo, J., Gu´elin, M., Gerin, M., McCarthy, M. C., & Thaddeus, P. 2008, A&A, 478, L19 Aikawa, Y., Herbst, E., Roberts, H., & Caselli, P. 2005, ApJ, 620, 330 Aikawa, Y., Wakelam, V., Garrod, R. T., & Herbst, E. 2008, ApJ, 674, 984 Bachiller, R., & Cernicharo, J. 1986, A&A, 168, 262 Bachiller, R., & G´omez-Gonz´olez, J. 1992, A&AR, 3, 257

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