A particular carbon-chain-producing region: L1489 starless core
Yuefang Wu, Lianghao Lin, Xunchuan Liu, Xi Chen, Tie Liu, Chao Zhang, Binggang Ju, Jinghua Yuan, Junzhi Wang, Zhiqiang Shen, Kee-Tae Kim, Sheng-Li Qin, Juan Li, Hongli Liu, Tianwei Zhang, Ye Xu, Qinghui Liu
aa r X i v : . [ a s t r o - ph . S R ] M a y May 23, 2019
A particular carbon-chain-producing region:L1489 starless core
Yuefang Wu , Lianghao Lin , , , Xunchuan Liu , Xi Chen , , Tie Liu , , Chao Zhang , , Binggang Ju , JinghuaYuan , Junzhi Wang , Zhiqiang Shen , Kee-Tae Kim , Sheng-Li Qin , Juan Li , Hongli Liu , Tianwei Zhang , YeXu , and Qinghui Liu Department of Astronomy, Peking University, 100871, Beijing, Chinae-mail: [email protected] School of Astronomy and Space Sciences, University of Science and Technology of China, 96 Jinzhai Road, Hefei, 230026, China Purple Mountain Observatory and Key Laboratory of Radio Astronomy, Chinese Academy of Sciences, 8 Yuanhua Road, Nanjing,210034, China Shanghai Astronomical Observatory, Chinese Academy of Sciences, Shanghai 200030, China Center for Astrophysics, GuangZhou University, Guangzhou, 510006, China Korea Astronomy and Space Science Institute, 776 Daedeokdae-ro, Yuseong-gu, Daejeon 34055, Korea East Asian Observatory, 660 North A’ohoku Place, Hilo, HI 96720, USA; Department of Astronomy, Yunnan University, Kunming, 650091, China National Astronomical Observatories, Chinese Academy of Sciences, 20A Datun Road, Chaoyang District, Beijing 100101, China Department of Physics, The Chinese University of Hong Kong, Shatin, NT, Hong Kong SARMay 23, 2019
ABSTRACT
We detected carbon-chain molecules (CCMs) HC n + N (n = S in K u band as well as high-energy excitation lines includingC H N = = / /
2, 19 / /
2, and CH CCH J = = H and CH CCH are close to those of L1527, and the CH CCH columndensities of the EMC and L1527 are slightly higher than those of TMC-1. The EMC and L1527 have similar C S column densities,but they are much lower than those of all the starless cores, with only 6.5% and 10% of the TMC-1 value, respectively. The emissionsof the N-bearing species of the EMC and L1527 are at the medium level of the starless cores. These comparisons show that the CCMemissions in the EMC are similar to those of L1527, though L1527 contains a protostar. Although dark and quiescent, the EMC iswarmer and at a later evolutionary stage than classical carbon-chain–producing regions in the cold, dark, quiescent early phase. ThePACS, SPIRE, and SCUBA maps evidently show that the L1489 IRS seems to be the heating source of the EMC. Although it islocated at the margins of the EMC, its bolometric luminosity and bolometric temperature are relatively high. Above all, the EMC is arather particular carbon-chain-producing region and is quite significant for CCM science.
Key words.
ISM: molecules — ISM: abundances — stars: formation — ISM: individual object (L1489)
1. Introduction
Carbon-chain molecules (CCMs), including radicals, have thelargest number of atoms among interstellar molecules found upuntil now . They have a large mass range and have di ff erent ex-citation energies and active states, and as a result they play im-portant roles in interstellar chemical and physical processes. Al-though it is di ffi cult for them to exist in terrestrial conditions,they were detected in star-forming regions and circumstellar en-velopes of evolved stars (Turner 1971; Avery et al. 1976; Suzukiet al. 1992; Taniguchi et al. 2018; Winnewisser & Walmsley1978; Ziurys 2006; Zhang et al. 2017).In molecular clouds, CCMs were found to be abundant indark and quiescent cores. TMC-1 is the typical carbon-chain-producing region. Molecules of HC N, HC N, and HC N werefirst detected in this core (Kroto et al. 1978; Broten et al. 1978;Bell et al. 1997). Almost all CCMs found so far, includingC n O and C n S, were discovered in TMC-1 (Herbst et al. 1984; / cdms / molecules Matthews et al. 1984; Irvine et al. 1988). High-resolution mapsof a number of S-bearing and N-bearing species were made andat least six cores were detected in the TMC-1 ridge (Hiraharaet al. 1992). Changes in molecular abundances along the ridgewere analyzed with the dynamical-chemical model (Markwicket al. 2000). In addition to TMC-1, L1521B, L1498, L1544,and L1521E were also detected as rich starless carbon-chain-producing regions (Suzuki et al. 1992; Kuiper et al. 1996; Ohashiet al. 1999; Hirota et al. 2002). All these cores in the molec-ular complex Taurus are in an early stage of CCM chemistry.Meanwhile a number of starless cores outside of Taurus, suchas L492, Lupus-1A, L1512, and Serpens South 1a (Serp S1a),have been found to be abundant carbon-chain-producing regions(Hirota & Yamamoto 2006; Sakai et al. 2010b; Cordiner et al.2011; Friesen et al. 2013; Li et al. 2016). Starless and dark coresCB 130-3 in the Aquila rift region and L673-SMM4 in CloudB of L673 were also identified as carbon-chain-producing re-gions (Hirota et al. 2011). Recently, 17 high-mass starless cores(HMSCs) and 35 high-mass protostellar objects (HMPOs) were
Article number, page 1 of 16 & Aproofs: manuscript no. ms surveyed with HC N and HC N by Taniguchi et al. (2018). Themolecule HC N was detected in 15 HMSCs and 28 HMPOs,and HC N was found in 5 HMSCs and 14 HMPOs (Taniguchiet al. 2018).All these cores are in an early chemical phase. However, theirproperties and evolutionary states could be somewhat di ff erent.In L1498, radial chemical di ff erentiation has been detected withC S and NH distributed in an onion shell-like structure withNH at the inner part and C S at the outer part, which may haveresulted from a slowly contracting dense core with a growingouter envelope (Kuiper et al. 1996). A similar structure of C Semission was discovered in L1544, which was interpreted as aresult of infall and rotation by Ohashi et al. (1999). Their di ff er-ent evolutional states can be sensitively traced with abundanceratios of CCMs and NH (Olano et al. 1988; Suzuki et al. 1992;Hirota & Yamamoto 2006).In star-forming cores, CCMs are usually less abundant thanin early cold and dark cores (Sakai et al. 2008). In particular, theabundance of S-bearing species is significantly lower in proto-stellar cores (Suzuki et al. 1992).However, in 2008 the protostellar core L1527 containingan infrared source IRAS 04368 + E up >
20 K; Sakai et al. 2008) of carbon-chain molecules such as C H (10 , − , ), C H N = − l -C H (4 , − , ), and CH CCH J = − , K = T MB ) of the line C H N = − J = / − / is about 30 K, it can be abundant in somewarmer regions around protostellar objects (Sakai et al. 2009).Subsequently, CH reacts with C + to form hydrocarbon ions.Such processes were proposed as warm carbon-chain chemistry(WCCC) by Sakai et al. (2008), which is di ff erent from thechemistry in early cores. More recently, formation of CCMs inWCCC (lukewarm corinos) was modeled with the macroscopicMonte Carlo method and it was found that the amount of CH can di ff use inside the ice mantle, and therefore sublimation uponwarm-up plays a crucial role in the synthesis of carbon-chainspecies in the gas phase (Wang et al. 2019).A second WCCC source, IRAS 15398-3359, was found bySakai et al. (2009) very shortly after the discovery of L1527. Amassive star-forming region NGC 3576 was also revealed as aWCCC source according to the detected C H J = / − / N, HC N, HC N, and evenHC N have been seen in L1527. Furthermore, HC N J = N J = −
31, the latter being a very-high-energy ex-citation transition, were detected in the second WCCC source(Sakai et al. 2008, 2009; Saul et al. 2015; Benedettini et al.2012). HC N J = −
10 was detected in NGC 3576 (Saulet al. 2015).L1489 is a famous low-mass star-forming source located inTaurus with a distance of 140 PC (Myers et al. 1988). In this pa-per we report CCM emissions detected towards the L1489 star-less core, that is, the eastern molecular core (EMC) of the L1489IRS (Benson & Myers 1989; Caselli et al. 2002). We made ob-servations of multiple spectral lines of CCMs for this core, in-cluding emissions of N- and S-bearing CCMs in the K u bandand high-energy excitation transitions in the 3 mm band. The observations are described in the following section. InSect. 3 we present the results. The discussions are presented inSection 4 and a summary is given in Sect. 5.
2. Observation
First we observed the spectral lines toward R.A.(J2000) = = N J = −
1, HC N J = − N J = − J = − J = −
15 as well asC S J = − u band, and C H N = − J = / − / J = / − /
2, CH CCH J = −
4, K = J = − =
1, and J = −
4, K =
2, c-C H (2 , − , ) as well as HC N J = − that is a compilation of the Jet Propulsion Labora-tory (Pickett et al. 1998), Cologne Database for Molecular Spec-troscopy (CDMS; (Müller et al. 2005)), and Lovas / NIST (Lovas2004) catalogs. This source was searched for HC N J = − J = −
4, HC N J = − J = −
16 as well as C S J N = − , J N = − (Suzuki et al. 1992; Fuller & Myers 1993; Brinchet al. 2007). All the transitions in this work were observed forthe first time. The spectral lines at K u band were observed with the Tian MaRadio Telescope (TMRT) of Shanghai Observatory on Jan 25,2016, and Dec 2-4, 2017. The TMRT is a newly built 65 m di-ameter fully steerable radio telescope located in the western sub-urb of Shanghai (Li et al. 2016). The front end of the K u bandis a cryogenically cooled receiver covering the frequency rangeof 11.5 − ′′ .An FPGA-based spectrometer based upon the design of Versa-tile GBT Astronomical Spectrometer (VEGAS) was employedas the Digital backend system (DIBAS) (Bussa & VEGAS De-velopment Team 2012). For molecular line observations, DIBASsupports a variety of observing modes, including 19 single sub-band modes and 10 modes with eight sub-bands each. The centerfrequency of the sub-band is tunable to an accuracy of 10 kHz.For our observation, the DIBAS mode 22 was adopted. Each ofthe eight side-bands has a bandwidth of 23.4 MHz and 16384channels. The main beam e ffi ciency is 60% at the K u band (Wanget al. 2015; Li et al. 2016). The beam sizes and the equivalentvelocity resolutions are given in the last two columns of Table1, respectively. After the spectra at O point were observed, wemade nine-point mapping observations on Dec 2-3, 2017. Theobservations were performed in point-by-point mode around apoint 1 ′ south of the O point to cover the L1489 IRS. We startthe map in a square pattern and with a grid separation of 1 ′ ; itsdiagonal lines are along E-W and N-S directions. To cover thenorthern part of the emission region, we added four samplingpoints: (1,0), (0,0.5), (0,-0.5), and (-0.5,0) on Dec 4, 2017. The spectral lines at 3 mm were observed with the 13.7 m tele-scope of the Qinghai Station of the Purple Mountain Observatory(PMO) on March 31, 2017. The pointing and tracking accuracieswere both better than 5 ′′ . The main beam e ffi ciency is 59% .A Superconducting Spectroscopic Array Receiver with sidebandseparation was employed at the front end (Shan et al. 2012) whileat the back end a Fast Fourier Transform Spectrometer with a to-tal bandwidth of 1 GHz allocated to 16,384 channels was used.The observed lines were covered in the two sidebands with fre-quencies from 85 to 86 GHz and 90 to 91 GHz, respectively.The spectral resolution is about 61 kHz. The system tempera-tures are in the range of 139 −
149 K, with a mean value of 144 K.The position-switch mode was adopted. The on-source time wasabout 25 minutes for all observed lines except CH CCH J = ′ × ′ for the lines of the CCMs in the 3 mm band with a 20 ′′ × ′′ grid.The IRAM software package GILDAS including CLASSand GREG was used for all the line data reduction (Guilloteau& Lucas 2000).
3. Results
All the transitions in K u band and 3 mm band are detectedtowards the EMC. Hyperfine lines of HC N J = − N J = − F = − F = − N J = − ∼ / s.Kaifu et al. (2004) detected HC N J = − ∼ / s). Li et al. (2016) detected this line to-wards Serpens South 1a with TMRT 65 m. However, only theHC N J = − F = − / s. Three hyperfine compo-nents F = − F = − , and F = − N J = − N J = − N J = − N map. In total, fivehyperfine lines of HC N J = − J = − F = − MB ashigh as 3.23 K. When zooming in on the hyperfine componentHC N J = − F = − − spanning from 6.75 to 7.08km s − .Panel (c) shows the contours of the integrated intensity of the redwing with a maximum value of 0.14 K km s − and a σ of 0.02K km s − overlaid on the map of the integration of line center.The red wing may belong to high-velocity gas since the profileof the wing is rather smooth and the ratio of the wing range tothe FWHM of the line is similar to that of the red wing of themolecular outflow S140 (Lada 1985). However, foreground andbackground cold gas as well as additional components cannot beexcluded (Wu et al. 2005). High-resolution observations may be useful to identify its origin. Further discussions are excluded inthe following analysis. The spectral lines of HC N J = − N J = − J = − J = −
15 as well as the C S J = − N J = − F = − H (2 , -1 , ), CH CCH J = − , and HC N J = − =
0, 1, and 2 of CH CCH J = − / N) of the weakest K = LSR , themain beam temperature T MB , the full width at half maximum(FWHM), and integrated intensities are listed in columns 3-5 ofTable 2. One can see that all the lines have a V LS R of about 6.6km s − except for C H and c-C H whose V LS R are about (6.8-6.9) km s − . The FWHMs of these two lines (C H and c-C H )are also broader than those of other detected CCM lines.Integrated intensity maps of all the detected lines were made.The maps of the lines in K u and the 3 mm band are presented inthe right panels of Figs. 2 and 3, respectively.One can see from the right-hand panels of Figs. 2 and 3that the emission peaks (black hollow triangle) of transitionsof N-bearing species in K u band are relatively consistent witheach other. This peak point (P point hereafter) is located atR.A.(J2000) = = H (green filled hollow triangle in Fig. 3) is located at < ff erent molecularspecies and can often be seen in molecular cores such as the NH cores in Orion, Cepheus, and H O maser sources as well as var-ious gas cores in the carbon-chain-producing region NGC 3576(Harju et al. 1993; Wu et al. 2006; Saul et al. 2015, see also Fig.7 and Sect. 4.1). We take the P point as the peak of the emissionsof all the CCMs except C S J = − Single point observations towards O point were made before themap. The highest S / Ns of the data at this point were compared tothose of other points. The comparison enabled us to analyze theprofile of molecular lines in detail. The lines at O point are thusalso shown in Fig. 4. The left panel of Fig. 4 presents the spectrallines of HC N, HC N, HC N, and C S in the K u band. The rightpanel presents the spectral lines of C H, CH CCH, c-C H , andHC N in 3 mm band.The spectral peak of the HC N J = −
13 at O point shownin Fig. 4 seems to have a dip at the center with an S / N ≈
3. Thespectrum of HC N J = −
14 may also be split but interferenceof noise cannot be excluded. This was not seen in the spectrumsof other CCMs in the K u or 3 mm bands at O point. For theemission of C S J = −
2, the T MB at O point is higher than thatat P point, and therefore O point is adopted as the peak of theC S core.
Article number, page 3 of 16 & Aproofs: manuscript no. ms
The column densities of the observed CCMs except C S werecalculated from the integrated intensities of lines of P point. ForC S, parameters at O point are adopted. Assuming the gas is inlocal thermodynamic equilibrium (LTE) and the lines are opti-cally thin, the column densities are calculated with the solutionof the radiation transfer equation (Garden et al. 1991; Mangum& Shirley 2015) N = k π ν Q P S i j µ exp E up kT ex ! J ( T ex ) Z τ dV , (1) T MB /η = f h ν k [( e h ν kTex − − − ( e h ν kTbg − − ][1 − e − τ ] , (2)where S i j , µ , and Q are the line strength, the permanent dipolemoment, and the partial function, respectively, which are quotedfrom “Splatalogue”, and T MB is the main beam temperature, η isthe e ffi ciency of the main beam, and f is the beam filling factor,assumed to be unity.The excitation temperature was calculated using di ff erentmethods in this work for comparison. Using the hyperfine struc-ture (HFS) fitting program in GILDAS / CLASS, we performedhyperfine structure fitting toward spectra of HC N J = − τ = . ± .
12 and an excitation temperature of 11 . ± . bg = =
0, 1, and 2 of CH CCH J = − rot and a value of 12.6 ± .
0K was obtained (see the middle panel of Fig. 5). Since the K = CCH J = (1,1) (2,2), the kinetic temperature T k is as-sumed equal to T rot , 12.6 ± . . The dust temperature (T d ) andH column density were derived via modeling the PACS andSPIRE data of Herschel at 160, 250, 350, and 500 µ m (see Table3 and the right panel of Fig. 5) to a modified black body: S ν = B ν ( T )(1 − e − τ ν ) Ω , (3)where τ ν = µ H m H κ ν N H / R gd ; here µ H = ff mann et al. (2008), m H is themass of a hydrogen atom, N H is the column density, and R gd =
100 is the ratio of gas to dust. The dust opacity κ ν can be ex-pressed as a power law of frequency, κ ν = . ν/ GHz ) β cm g − , (4)with κ ν (850 GHz) = g − adopted from Ossenkopf &Henning (1994). The free parameters are the dust temperature,dust emissivity index β, and column density. The fitting resultsgive T d = ± N H = (1.02 ± × cm − with β = . d vari-ation from the inner to outer parts of the core, which was modelfitted from the SCUBA 850 µ m image (Ford & Shirley 2011),is also listed in Table 4. These values are comparable with thekinetic temperature of L1527 (12.3 ± For the column density of the detected species, only that ofCH CCH can be derived by analysis of multiple transitions atdi ff erent energy levels. For other species including the hyperfinelines of HC N and multiple rotation lines of HC N , the up-per level energies of the multiple lines are close to one anotherand this method is impractical. The average values of columndensities calculated from multiple lines weighted by T MB / σ wastherefore adopted as the species column density.The T ex uncertainty will bring in about 10% error for calcu-lation of column densities of all the species.The column densities are listed in Table 2. The uncertain-ties of column densities are derived from errors of T ex and lineintegrated intensities. The column densities of our N-bearingmolecules range from 1 . × to 4 . × cm − . Among allthe detected molecules in the K u band, C S has the lowest col-umn density, 0.8 × cm − , which is much lower than those ofstarless cores such as TMC-1, L1544, and L1498 (Suzuki et al.1992); this is also lower than that of the stellar core L1251A(Cordiner et al. 2011).
4. Discussion
The species C H N = = / /
2, 17 / / CCH J = =
0, 1, 2 were detected in the WCCC sourceL1527 (Sakai et al. 2008). CH CCH J = = S is very weak.Figure 6 presents the CCM column densities of the EMC andL1527 as well as the five early carbon-chain-producing regionsnormalized by the values of TMC-1. One can see that the columndensities of the species with high-energy excitation lines in theEMC and L1527 are close to each other. The CH CCH columndensities of the two cores are slightly higher than that of TMC-1while their values of C H are lower than that of TMC-1. Amongfive starless cores, L1521B is the only one where the C H isdetected. It is worth noting that the column density of C H inL1521B is even higher than that of TMC-1. In this source thedetected C H line is N = − J = / − / F = − , − N J = up N column densityin the EMC is higher than that of L1527. The N-bearing speciesseem to be abundant in WCCC sources in general. In the WCCCsource IRAS 15398-3359, a very-high-energy excitation line ofHC N J = up N column densityin IRAS 15398-3359 is 1.5x10 cm − (Wu et al. in preparation).In the massive core WCCC source NGC 7536 the HC N columndensity is 1x10 cm − which is derived from the detected line ofHC N J = up of the HC N J = S column densities of both sources, about 6.5% (EMC) and10% (L1527) of the TMC-1 value, respectively, are much lower
Article number, page 4 of 16uefang Wu et al.: A particular carbon-chain-producing region: L1489 starless core than all of the compared early starless cores. It was recognizedpreviously that S-bearing species are usually deficient in star-forming cores. A number of low-mass star-forming cores wereexamined for emissions of S-bearing CCMs by Suzuki et al.(1992). Results showed no or marginal detection for most proto-stellar cores and L1489 was taken as an example. In the EMC,the emission of C S J = − S J = − S J = − (Myers et al. 1988), similar to the cases ofL1498 and L1544. But even at the O point, the column densityof the C S is still much lower than those of L1544 and L1498.All these comparisons present the following characteristicsof the EMC.1. Emissions of the species with high-energy excitation linesof the EMC are close to those of L1527. For these two sourcesthe column densities of the species with high-energy excitationlines are comparable with those of TMC-1.2. For the N-bearing species, the column densities of theEMC and L1527 are at the medium level among all the samples.3. The column density of C S of the EMC is close to that ofL1527 but is much lower than those of all starless cores.In short, the CCM emissions of the EMC are similar to thoseof L1527 but deviate from the early cold starless cores.
L1489 is one of the 90 small visually opaque regions chosenfrom the Palomar Sky Atlas prints (Myers et al. 1983). NH (1,1)mapping presents a gas core which is about half an arcminutesouth of the high visual opacity and about one arcminute eastof the IRAS 04016 + H + J = (1,1) and N H + J = S distributions of L1498 and L1544 have a centralhole which can be explained with infall and rotation (Hirota &Yamamoto 2006; Ohashi et al. 1999). In the EMC no sign of col-lapse or rotation has been detected so far and it is quiescent. At Ppoint there is strong NH emission together with weak C S emis-sion, which may be related to the evolutional state of the core(Suzuki et al. 1992; Hirota & Yamamoto 2006).The abundance ratios of NH / C S of TMC-1, L1521B, L492,L1498 and L1544 range from 2.9 to 25 (Hirota & Yamamoto2006). For the EMC, using the NH column density of My-ers et al. (1988) and the C S column density estimated fromC S emission with an average abundance ratio C S / C S of4.3 ± / C S is 289. This is one order of magnitude larger than thelargest of the seven cores listed in Hirota & Yamamoto (2006).It is also larger than the NH / C S ratio of 37 for Serp S1a, whichwas derived from NH and C S column densities with the ratioof C S / C S 4.3 (Friesen et al. 2013; Li et al. 2016; Suzuki et al.1992). Serp S1a has a more active and complex environment thanstarless cores in Taurus and infall was detected (Friesen et al.2013). These indicate that the EMC has the latest evolutionarystate among all the starless cores shown in Fig. 6.A prominent di ff erence between the EMC and other earlystarless cores is the temperature and the thermalisation of thegas. In TMC-1 the rotational temperature ranges from 4 to 8 K while the kinetic temperature ranges from 9 to 10 K (Kawaguchiet al. 1991; Kalenskii et al. 2004; Snyder et al. 2006; Sakai et al.2008). The dust temperature is 10.5-12 K (Fehér et al. 2016)and 10.6 ± k (1,1) (2,2) and 7 K was adopted as T ex (Friesen et al. 2013; Li et al. 2016). The T ex of L492 was derivedas 6.4 K from HC N and 8.3 K from CO, respectively (Hirota& Yamamoto 2006). L1521B has a T ex of 5.5-6.5 K and its T K is about 10 K (Hirota et al. 2004; Suzuki et al. 1992). The T ex of C S of L1498 ranges from 5.5 to 10 K (Suzuki et al. 1992;Kuiper et al. 1996), and the T k is 10 K derived from NH emis-sion in this core. The T d of L1498 is 10 K (Tafalla et al. 2004).The T ex of L1544 is from 5 to 6.5 K (Suzuki et al. 1992) and T K derived from NH is ∼
10 K , and the T d of L1544 is also 10K (Tafalla et al. 2002). In the EMC however, the T ex is 11.5 Kderived from HC N and T rot is 12.6 K from CH CCH, whichare similar to the T rot of C H (12.3 ± CCH (13.9K) for L1527 (Sakai et al. 2008). These comparisons show thatthe gas of the EMC is warmer and closer to thermalization sta-tus than those in early starless cores and comparable to thatof L1527. These may explain why the carbon-chain molecularemissions in the EMC are similar to those of WCCC sources.However, there is a fundamental di ff erence in that L1527contains a protostar which is the heating source of the core mate-rial (Goldreich & Kwan 1974; Sakai et al. 2008). For the EMC,there is no protostar inside, only an association with a protostel-lar object L1489 IRS at one arcminute to the west.Figure 7 displays the images of 850 µ m, 500 µ m, and 250 µ m continuum emissions as well as molecular line emissions inthe L1489 region. One can see that the P point (black triangle),the centers of larger 850 µ m continuum core, and the NH core(blue triangle) are close to each other but not overlaid exactly.Such deviation among peaks of di ff erent molecular line emissionregions are very common in prestellar and stellar cores even inthe TMC-1 region (Sánchez-Monge et al. 2014; Keown et al.2016; Pratap et al. 1997). Cores of other molecular lines suchas N H + J = = µ m map andthe SPIRE 500 µ m image covered the EMC entirely.The L1489 IRS is a Class I source while the one in L1527 isa Class 0 object. The L bol is 3.6-3.8 L ⊙ for the L1489 IRS and1.3-1.9 L ⊙ for L1527, respectively (Kristensen et al. 2012; Chenet al. 1995). The T bol of L1527 is 44 K (Kristensen et al. 2012;Chen et al. 1995), while the T bol of the L1489 IRS is 200-283K and more than five times of that of L1527 on average (Kris-tensen et al. 2012; Chen et al. 1995). The C H emission extendsfrom the center part of the envelope to 5600 AU in L1527. High-energy excitation CCM lines in the EMC are possible to be ex-cited, though L1489 IRS is located at the margins of the EMCsince L1489 IRS is has relatively high bolometric luminosity andtemperature. Further more, H O 1 , -1 , observed with HIFI onHerschel presents an inverse P-cygni profile in L1527 but a P-Cygni profile in L1489 IRS, indicating that L1489 IRS lacks anouter, cooler gas layer, which is in favor of the heating of itsneighboring material.From Fig. 7 one can see that at each wavelength, thestrongest emission is around the IRS. The higher the frequency,the more the contours concentrate at the IRS. The tails of thecontours of the 250 µ m and 500 µ m maps present a trajectoryformat starting from the IRS, which has not previously been seenin low-mass protostellar cores. These may be the evidence thatthe EMC is heated by the L1489 IRS. The size of the 250 µ memission region is 169.1 arcsec and T d from the spectral energy Article number, page 5 of 16 & Aproofs: manuscript no. ms distribution (SED) fitting based on data of PACS-SPIRE bandsis 13.8 K in the whole region.The heating of L1489 IRS can be more significant consider-ing the possible evolutionary history of the relationship betweenL1489 IRS and the EMC; they may be neighbors since the birthof L1489 IRS. However, another possibility may be that the IRShas moved away from the EMC center to the cloud margin, or ithas dispersed the surrounding gas and has separated itself fromthe EMC (Harju et al. 1993; Wu et al. 2006). The influences ofthe L1489 IRS might what gives the EMC its WCCC character-istics.However, WCCC began from reactions of C + and sublimatedCH from dust. A precondition is that the temperature needs tobe ∼
30 K (Sakai et al. 2009). For L1527, the kinetic tempera-ture is 20-30 K derived from the c-C H , -4 , line detectedwith the Plateau de Bure Interferometer (Sakai et al. 2010a).From the c-C H and SO lines measured with ALMA withbeam sizes 0.8 ′′ × ′′ and 0.7 ′′ × ′′ and analyzed with non-LTE large-velocity-gradient code, the kinetic temperatures of thec-C H emitting region was found as 30 K at 100 AU (Sakaiet al. 2014). The kinetic temperatures of L1527 measured fromthe interferometers are much higher than 13.9 K derived fromthe CH CCH J = = ∼ ′′ . These results present e ff ects of tele-scope beams. The kinetic temperature of the EMC derived fromCH CCH J = = ′′ is 12.6 ± ± bol and T bol of the L1489 IRS arehigher than those of L1527. These comparisons indicate that ina smaller region within the EMC the kinetic temperature may behigher than the current results. Higher-resolution observationsare needed to measure the temperature.In the WCCC source L1527, CCMs of the prestellar phasecould survive if the prestellar collapse is faster than that of otherstar-forming cores (Sakai et al. 2008). The EMC should have anearly and a cold phase, like those starless cores shown in Fig. 6.No collapse or infall have been found in the EMC so far. TheCCMs detected in the EMC may have survived from its earlyphase. In particular, the chemical activities of the CCMs seem tobe related to their existence in the WCCC phase. Because of thedi ff erent polarities, the N- and S-bearing species are more ac-tive than C H and CH CCH and more likely to react with theirpartners. For example, the major loss route of C H is throughreaction with C + (Millar & Freeman 1984): C H + C + → C + + H,which has a rate coe ffi cient of 2.0 × − cm at 10 K. Whilethe reaction of HC N with C + : HC N + C + → C HN + + C has arate coe ffi cient of 8.7 × − cm at the same temperature (Leunget al. 1984). However, both species are formed from hydrocarbonions and the reaction rates are not significantly di ff erent fromone another (Leung et al. 1984). This suggests that the molecu-lar species with high-energy excitation lines have a higher rateof survival than low-excitation lines, which might be a reasonfor the emissions of species with high-energy excitation lines inL1527 (Sakai et al. 2008) and the late core EMC. This may alsobe the reason for the detection of the HC N J = up H with T MB N J = T MB S show consistent emission peaks at about 1 ar-cmin east of the L1489 IRS. These results show that the EMCis an abundant CCM laboratory. Its comparable and contrast-ing conditions with the WCCC source L1527 and with earlycold carbon-chain-producing regions may promote further CCMsearches and help to constrain model analyses of CCMs.
5. Summary
Abundant carbon-chain molecules were detected toward theEMC of L1489 IRS, which is identified as a particular carbon-chain-producing region.1. With the TMRT, HC N J = N J = N J = S J = MB of the hyperfine line J = = N is3.23 K. Five hyperfine lines of HC N J = − = = = N J = − MB of thethree rotational lines of HC N ranges from 0.25 to 0.32 K. Theemission of C S J = H N = = / /
2, 19 / /
2, CH CCH J = =
2, and HC N J = u and 3 mm band aredetected for the first time in the EMC of L1489 IRS.2. Maps of the observed transition lines were also obtained.Emission peaks of our detected lines except C S are all locatedat about 1 ′ east of the L1489 IRS, which is consistent with pre-viously detected gas cores and the 850 µ m continuum east core.3. The CCM column densities of the EMC were comparedwith those of TMC-1 together with five carbon-chain-producingregions in early phase (Serp S1a, L492, L1521B, L1498 andL1544) and WCCC source L1527. Results show that the columndensities of the species with high-energy excitation lines includ-ing CH CCH and C H in the EMC are close to those of L1527.The C S column density of the EMC is slightly lower than thatof L1527 but much lower than the five starless cores. The columndensities of the N-bearing species are close to those of L1527,and the values of the both sources are at the intermediate level ofthe starless cores.4. Similarly to early carbon-chain-producing regions, theEMC is dark, starless, and quiescent. However, the EMC israther at a late evolutionary stage ( N (NH ) / N (C S) = µ m and the SPIRE 500 µ m imagescovered the EMC completely. The tails of the contours of the 250 µ m and 500 µ m maps present a trajectory format starting fromthe IRS. The dust continuum emission area and the morphologyof the contours show that the EMC is externally heated by theL1489 IRS. Acknowledgements.
We are grateful to the sta ff of PMO Qinghai Station andSHAO. We also thank Shanghuo Li, Kai Yang and Bingru Wang for theirassistance during the observation period. This project was supported by thegrants of National Key R&D Program of China No. 2017YFA0402600, NSFCNos. 11433008, 11373009, 11373026, 11503035, 11573036, U1331116 andU1631237, and the Top Talents Program of Yunnan Province. This researchused the facilities of the Canadian Astronomy Data Centre operated by the theNational Research Council of Canada with the support of the Canadian SpaceAgency. Article number, page 6 of 16uefang Wu et al.: A particular carbon-chain-producing region: L1489 starless core
References
Askne, J., Hoglund, B., Hjalmarson, A., & Irvine, W. M. 1984, A&A, 130, 311Avery, L. W., Broten, N. W., MacLeod, J. M., Oka, T., & Kroto, H. W. 1976,ApJ, 205, L173Bell, M. B., Feldman, P. A., Travers, M. J., et al. 1997, ApJ, 483, L61Benedettini, M., Pezzuto, S., Burton, M. G., et al. 2012, MNRAS, 419, 238Benson, P. J., & Myers, P. C. 1989, ApJS, 71, 89Brinch, C., Crapsi, A., Hogerheijde, M. R., & Jørgensen, J. K. 2007, A&A, 461,1037Broten, N. W., Oka, T., Avery, L. W., MacLeod, J. M., & Kroto, H. W. 1978,ApJ, 223, L105Bussa, S., & VEGAS Development Team. 2012, in American Astronomical So-ciety Meeting Abstracts, Vol. 219, American Astronomical Society MeetingAbstracts ff mann, J., Bertoldi, F., Bourke, T. L., Evans, II, N. J., & Lee, C. W. 2008,A&A, 487, 993Kawaguchi, K., Kaifu, N., Ohishi, M., et al. 1991, PASJ, 43, 607Keown, J., Schnee, S., Bourke, T. L., et al. 2016, ApJ, 833, 97Kristensen, L. E., van Dishoeck, E. F., Bergin, E. A., et al. 2012, A&A, 542, A8Kroto, H. W., Kirby, C., Walton, D. R. M., et al. 1978, ApJ, 219, L133Kuiper, T. B. H., Langer, W. D., & Velusamy, T. 1996, ApJ, 468, 761Lada, C. J. 1985, ARA&A, 23, 267Leung, C. M., Herbst, E., & Huebner, W. F. 1984, ApJS, 56, 231Li, J., Shen, Z.-Q., Wang, J., et al. 2016, ApJ, 824, 136Lovas, F. J. 2004, Journal of Physical and Chemical Reference Data, 33, 177Mangum, J. G., & Shirley, Y. L. 2015, PASP, 127, 266Markwick, A. J., Millar, T. J., & Charnley, S. B. 2000, ApJ, 535, 256Matthews, H. E., Irvine, W. M., Friberg, P., Brown, R. D., & Godfrey, P. D. 1984,Nature, 310, 125Millar, T. J., & Freeman, A. 1984, MNRAS, 207, 405Müller, H. S. P., Schlöder, F., Stutzki, J., & Winnewisser, G. 2005, Journal ofMolecular Structure, 742, 215Myers, P. C., Fuller, G. A., Mathieu, R. D., et al. 1987, ApJ, 319, 340Myers, P. C., Heyer, M., Snell, R. L., & Goldsmith, P. F. 1988, ApJ, 324, 907Myers, P. C., Linke, R. A., & Benson, P. J. 1983, ApJ, 264, 517Ohashi, N., Lee, S. W., Wilner, D. J., & Hayashi, M. 1999, ApJ, 518, L41Olano, C. A., Walmsley, C. M., & Wilson, T. L. 1988, A&A, 196, 194Ossenkopf, V., & Henning, T. 1994, A&A, 291, 943Pickett, H. M., Poynter, R. L., Cohen, E. A., et al. 1998,J. Quant. Spectr. Rad. Transf., 60, 883Pratap, P., Dickens, J. E., Snell, R. L., et al. 1997, ApJ, 486, 862Sakai, N., Sakai, T., Hirota, T., Burton, M., & Yamamoto, S. 2009, ApJ, 697, 769Sakai, N., Sakai, T., Hirota, T., & Yamamoto, S. 2008, ApJ, 672, 371—. 2010a, ApJ, 722, 1633Sakai, N., Shiino, T., Hirota, T., Sakai, T., & Yamamoto, S. 2010b, ApJ, 718,L49Sakai, N., Sakai, T., Hirota, T., et al. 2014, Nature, 507, 78 Sánchez-Monge, Á., Beltrán, M. T., Cesaroni, R., et al. 2014, A&A, 569, A11Saul, M., Tothill, N. F. H., & Purcell, C. R. 2015, ApJ, 798, 36Shan, W.-Y., Lu, H.-Z., & Shen, S.-Q. 2012, Phys. Rev. B, 86, 125303Snyder, L. E., Hollis, J. M., Jewell, P. R., Lovas, F. J., & Remijan, A. 2006, ApJ,647, 412Suzuki, H., Yamamoto, S., Ohishi, M., et al. 1992, ApJ, 392, 551Tafalla, M., Myers, P. C., Caselli, P., & Walmsley, C. M. 2004, A&A, 416, 191Tafalla, M., Myers, P. C., Caselli, P., Walmsley, C. M., & Comito, C. 2002, ApJ,569, 815Tafalla, M., Santiago-García, J., Myers, P. C., et al. 2006, A&A, 455, 577Taniguchi, K., Saito, M., Sridharan, T. K., & Minamidani, T. 2018, ApJ, 854,133Turner, B. E. 1971, ApJ, 163, L35Wang, J. Q., Yu, L. F., Zhao, R. B., et al. 2015, Acta Astronomica Sinica, 56, 63Wang, Y., Chang, Q., & Wang, H. 2019, A&A, 622, A185Winnewisser, G., & Walmsley, C. M. 1978, A&A, 70, L37Wu, Y., Zhang, Q., Chen, H., et al. 2005, AJ, 129, 330Wu, Y., Zhang, Q., Yu, W., et al. 2006, A&A, 450, 607Yen, H.-W., Takakuwa, S., Ohashi, N., et al. 2014, ApJ, 793, 1Zhang, X.-Y., Zhu, Q.-F., Li, J., et al. 2017, A&A, 606, A74Zhou, S., Wu, Y., Evans, II, N. J., Fuller, G. A., & Myers, P. C. 1989, ApJ, 346,168Ziurys, L. M. 2006, Proceedings of the National Academy of Science, 103,12274 Article number, page 7 of 16 & Aproofs: manuscript no. ms
Table 1.
Observed transitions and telescope parameters
Molecular Q . Q . Transition freq.(MHz) S i j µ ( D ) E low ( K ) E up ( K ) FWHM(") v chan ( m / s )HC N 43.27 86.22 J = = = = = = = = = = N 147.1 293.8 J = = = = = = N 346.7 693.0 J = = = S 68.0 135.5 J = H 165.5 329.7 N = = / / = = / / CCH 34.42 88.27 J = = = H = N 43.27 86.22 J = Article number, page 8 of 16uefang Wu et al.: A particular carbon-chain-producing region: L1489 starless core
Table 2.
Observed and derived parameters
Species Transitions V LS R1 T MB1 ∆ V R T MB dv N ( S ) (km / s) (mK) (m / s) (K m / s) (10 cm − )HC N J = = = = = = = = = = N J = = = = = = N J = = = S J = C H N = = / / = = / / CCH J = = = H N J = ( ) The uncertainties of V LS R , T MB and ∆ V are obtained from Gaussian fitting. ( ) The uncertainties of R T MB dv are calculated from error transfer formula. ( ) Column densities of species are adopted as average values calculated from di ff erent hyperfine lines weighted by T MB / σ , assumingall are optical thin with an identical excitation temperature, except that of HC N based on HC N J = − CCH whose column density is adopted as the value given by rotation diagram. Uncertaintiesof column densities are calculated from error transfer formula, including the error introduced from the uncertainty of excitationtemperature. ( ) C S parameters at O point are listed.
Table 3.
Dust parameters
Source S S S S S Size T d N H Jy Jy Jy Jy Jy arcsec K 10 cm − L1489 38.07 84.69 102.86 75.58 38.22 169.07 13.8(0.2) 1.0(0.1) ( The 850 µ m flux (5.78 Jy) is combined for the SED of L1489, which is from (Hogerheijde & Sandell 2000). ) Table 4.
Temperatures T ex (HC N) T rot (CH CCH) T d T d (inner-outer)(K) (K) (K) (K)11.5(1.0) 12.6(1.0) 13.8(0.2) 11.7-12.4 ( The T d (inner-outer) is model fitted from the SCUBA 850 µ m map of L1489 (Ford & Shirley 2011). ) Article number, page 9 of 16 & Aproofs: manuscript no. ms Ra D e c Fig. 1. (a): HC N J = = N J = = N J = = / s − / s) of HC N J = = / s. The red contour denotes 3 σ level (0.06 Kkm / s) of HC N wing integration. The filled black square represents O point, and the black triangle represents P point (see text). The IRS is shownby the hexagonal star (Yen et al. 2014). Small black crosses show sampled points of K u band observation.Article number, page 10 of 16uefang Wu et al.: A particular carbon-chain-producing region: L1489 starless core OP IRS
Fig. 2.
Emissions of HC N J = N J = N J = S J = u band. Left: Spectral lines at the P point. Theemission of HC N J = = N J = = N J = N J = S J = σ of integrated emissions. The symbolsare the same as those in Fig. 1(c) . Article number, page 11 of 16 & Aproofs: manuscript no. ms
Fig. 3.
Emissions of the transitions in the 3 mm band. Left: Spectral lines of the P point. Right: Emission intensity contours from 50% to 90% insteps of 10% of the peak value (see Table 2). C H N = = / / H N = = / / H intensitycontours. Red contours represent 3 σ of integrated emissions. The symbols are the same as those in Fig. 1(c). The green and pink filled trianglesrepresent peak position of c-C H J = N J = Fig. 4.
Spectra at O position. Left: Spectra at K u band. Right: Spectra at the 3 mm band detected with the 13.7 m telescope of PMO.Article number, page 13 of 16 & Aproofs: manuscript no. ms up (K)10 N u × Q ( T r o t ) g CH CCHT rot = 12.6±1.0 K ( m)10 S ( J y ) N H = (1.02±0.07)×10 cm T d = 13.80±0.21 K = 1.75L1489 Fig. 5.
Left: Spectra of CH CCH J = = CCH . Right: SED of the L1489 IRS from thePACS 160 µ m and SPIRE wavelengths of Herschel as well as SCUBA 850 µ m. The filled squares represent the input fluxes. The line shows thebest fitting of the gray-body model.Article number, page 14 of 16uefang Wu et al.: A particular carbon-chain-producing region: L1489 starless core H C N H C N H C N C S C H C H C C H −1 N ( X ) / N ( T M C ) L1489L1527Serp S1aL492L1521BL1498L1544
Fig. 6.
Comparison between the CCM column densities of L1489 EMC and typical CCM rich sources, normalized by the values of TMC-1 (Kaifuet al. 2004; Suzuki et al. 1992). The source names are denoted in the lower-right corner with di ff erent colors, including starless cores SerpensSouth 1a (Serp S1a; Li et al. 2016), L492 (Hirota & Yamamoto 2006; Hirota et al. 2009), L1521B (Suzuki et al. 1992; Hirota et al. 2004), L1498(Suzuki et al. 1992; Kuiper et al. 1996) and L1544 (Suzuki et al. 1992) as well as WCCC source L1527 (Sakai et al. 2008). The HC N columndensity of L1544 is derived from that of HC N assuming the ratio of HC n + / HC n + (n = N (C S) ofL1527 is deduced from column densities of C S and the ratio of the C S / C S (Sakai et al. 2008; Suzuki et al. 1992).Article number, page 15 of 16 & Aproofs: manuscript no. ms AU a D e c AU a D e c AU a D e c Fig. 7.
Left: The HC N core (gray-scale) and the NH (1,1) core (blue dashed lines, quoted from Myers et al. (1988) ) are overlaid on the JCMTSCUBA 850 µ m continuum data from JCMT proposal ID M97AN16 (Hogerheijde & Sandell 2000) (cyan contours, evenly stepped from 0.075to 0.75 Jy / beam in log-scale). The gray and blue triangles denote the peaks of the EMC (P point) and NH , respectively. The IRS is also marked onthe map (hexagonal star). The blue and red wings of the CO (3-2) outflow (blue solid and red dotted lines, quoted from Hogerheijde et al. (1998))and the medium infalling lobe (red solid line, quoted from Yen et al. (2014)) were also overlaid on the figure. Middle: As in left panel except cyancontours representing Herschel SPIRE 500 µ m continuum map, from 1.5 Jy / beam to 5 Jy / beam stepped by 0.3 Jy / beam. Right: As in left panelexcept cyan contours representing Herschel SPIRE 250 µ m continuum contours, from 1.5 Jy / beam to 5 Jy / beam stepped by 0.3 Jy //