Hunting for the elusive methylene radical
A. M. Jacob, K. M. Menten, Y. Gong, P. Bergman, M. Tiwari, S. Bruenken, A.O.H. Olofsson
AAstronomy & Astrophysics manuscript no. main_arxiv © ESO 2021January 8, 2021
Hunting for the elusive methylene radical
A. M. Jacob (cid:63) , K. M. Menten , Y. Gong , P. Bergman , M. Tiwari , , S. Brünken , and A.O.H. Olofsson Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany Department of Space, Earth and Environment, Chalmers University of Technology, Onsala Space Observatory, 43992 Onsala,Sweden University of Maryland, Department of Astronomy, College Park, MD 20742-2421, USA Radboud University, Institute for Molecules and Materials, FELIX Laboratory, Toernooiveld 7, 6525 ED Nijmegen, The Nether-landse-mail: [email protected]
Received November 13, 2020; accepted December 21, 2020
ABSTRACT
Context.
The N K a K c = − transitions of ortho-CH between 68 and 71 GHz were first detected toward the Orion-KL andW51 Main star-forming regions. Given their high upper level energies (225 K) above the ground state, they were naturally thought toarise in dense, hot molecular cores near newly formed stars. However, this has not been confirmed by further observations of theselines and their origin has remained unclear. Generally, there is a scarcity of observational data for CH and, while it is an importantcompound in the astrochemical context, its actual occurrence in astronomical sources is poorly constrained. Aims.
In this work, we aim to investigate the nature of the elusive CH emission, address its association with hot cores, and examinealternative possibilities for its origin. Owing to its importance in carbon chemistry, we also extend the search for CH lines byobserving an assortment of regions, guided by the hypothesis that the observed CH emission is likely to arise from the hot gasenvironment of photodissociation regions (PDRs). Methods.
We carried out our observations first using the Kitt Peak 12 m telescope to verify the original detection of CH towarddi ff erent positions in the central region of the Orion Molecular Cloud 1. These were followed-up by deep integrations using thehigher angular resolution of the Onsala 20 m telescope. We also searched for the N K a K c = − transitions of para-CH between440–445 GHz toward the Orion giant molecular cloud complex using the APEX 12 m telescope. We also obtained auxiliary data forcarbon recombination lines with the E ff elsberg 100 m telescope and employing archival far infrared data. Results.
The present study, along with other recent observations of the Orion region reported here, rule out the possibility of anassociation with gas that is both hot and dense. We find that the distribution of the CH emission closely follows that of the [CII]158 µ m emission, while CH is undetected toward the hot core itself. The observations suggest, rather, that its extended emissionarises from hot but dilute layers of PDRs and not from the denser parts of such regions as in the case of the Orion Bar. This hypothesiswas corroborated by comparisons of the observed CH line profiles with those of carbon radio recombination lines (CRRLs), whichare well-known PDR tracers. In addition, we report the detection of the 70 GHz fine- and hyperfine structure components of ortho-CH toward the W51 E, W51 M, W51 N, W49 N, W43, W75 N, DR21, and S140 star-forming regions, and three of the N K a K c = − fine- and hyperfine structure transitions between 68-71 GHz toward W3 IRS5. While we have no information on the spatial distributionof CH in these regions, aside from that in W51, we again see a correspondence between the profiles of CH lines and those of CRRLs.We see a stronger CH emission toward the extended HII region W51 M rather than toward the much more massive and denser W51E and N regions, which strongly supports the origin of CH in extended dilute gas. We also report the non-detection of the 2 − transitions of para-CH toward Orion. Furthermore, using a non-LTE radiative transfer analysis, we can constrain the gas temperaturesand H density to (163 ±
26) K and (3 . ± . × cm − , respectively, for the 68–71 GHz ortho-CH transitions toward W3 IRS5,for which we have a data set of the highest quality. This analysis confirms our hypothesis that CH originates in warm and dilutePDR layers. Our analysis suggests that for the excitation conditions under the physical conditions that prevail in such an environment,these lines are masering, with weak level inversion. The resulting amplification of the lines’ spontaneous emission greatly aids in theirdetection. Key words.
ISM: molecules – ISM: abundances – ISM: clouds – ISM: lines and bands – methods: observational – radiative transfer
1. Introduction
The methylene radical, CH , is of considerable astrophysical in-terest given that it is both produced and destroyed at an earlystage in the sequence of ion-molecule reactions that govern inter-stellar chemistry (see Godard et al. 2014, and references therein).Theoretical models of the di ff use interstellar clouds in the lineof sight (LOS) toward ζ Per by Black et al. (1978) and van (cid:63)
Member of the International Max Planck Research School (IM-PRS) for Astronomy and Astrophysics at the Universities of Bonn andCologne
Dishoeck & Black (1986) as well as equilibrium models of denseclouds in Orion by Prasad & Huntress (1980) have predictedCH to possess abundances similar to, if not greater than, thatof CH. Both species are primarily formed by the dissociative re-combination of the methyl ion CH + with an electron and bothare destroyed in dense clouds via reactions with atomic oxygento form HCO and HCO + , which are chemical species that act asbuilding blocks for more complex interstellar molecules. Vejby-Christensen et al. (1997) computed the complete branching ra-tios of the di ff erent possible reaction products of the dissociativerecombination reaction of CH + and found that CH was the ma- Article number, page 1 of 25 a r X i v : . [ a s t r o - ph . GA ] J a n & A proofs: manuscript no. main_arxiv jor reaction product with a branching ratio of 40%. Additionally,CH has been speculated to play an important role in the photo-dissociation sequence of methane (CH ) in cometary ice man-tles (van Dishoeck et al. 1996). Despite its importance and largepredicted abundances, CH is yet to be detected in comets andonly a handful of detections have been made in the interstellarmedium (ISM), unlike the ubiquitous CH.Even laboratory measurements of this simple radical haveproven to be di ffi cult owing to its reactive nature and to its light-ness and peculiar b-type selection-rules (Michael et al. 2003;Brünken et al. 2004). The latter two characteristics result inwidely spaced energy levels particularly between the energeti-cally lowest rotational states of CH . With only four rotationaltransitions below 1000 GHz (as shown in Fig. 1), the rotationaltransitions of CH are di ffi cult to target in the laboratory, andeven more so with astronomical observations, as they are eitherinaccessible from the ground or lie close to the edges of atmo-spheric windows.CH was first unambiguously detected in the ISM by Holliset al. (1995) who identified the intrinsically strongest fine- andhyperfine-structure (HFS) transitions of the N K a K c = − multiplet of ortho-CH . The frequencies of the detected lines liein the 68–71 GHz spectral range ( ≈ . , which e ff ectively demarcates the radio- from the millime-tre wavelength range. Their detections became possible after ac-curate frequencies had been measured in the laboratory by Lo-vas et al. (1983). These transitions arise from levels with ener-gies of ∼
225 K above the ground state and were observed inemission toward the dense molecular ‘hot cores’ associated withthe Kleinmann-Low Nebula in Orion Molecular Cloud 1 (Orion-KL) and W51 Main.Following this work, several other astronomical searches forCH have been published, but Lyu et al. (2001) reported the ten-tative detection of several of the electronic bands of CH near1410 and 1416 Å toward HD 154368 and ζ Oph in absorption.However, recently Welty et al. (2020) were only able to publishan upper limit for CH absorption near 1397 Å from the rela-tively dense molecular core of the translucent interstellar cloudin the LOS to HD 62542. Their upper limits on the CH col-umn density are six and three times lower, respectively, than thenominal values given by Lyu et al. (2001) for HD 154368 and ζ Oph.A breakthrough came with the detection of far-infrared (FIR)absorption lines from low-lying energy levels of both of CH ’sspin isomeric forms, the ortho- and para-CH , which were dis-covered toward the intense dust continuum emission of theSgrB2 (M) and W49 N high-mass star-forming regions (SFRs)by Polehampton et al. (2005) from the data taken with theLong Wavelength Spectrometer (LWS, Clegg et al. 1996) aboardthe Infrared Space Observatory (ISO); see also Polehamptonet al. (2007). Absorption was not only found at the systemicLSR velocities of the SFRs, but also at the velocities of di ff useand translucent interstellar clouds intervening along the linesof sight. These detections were enabled by accurate frequencymeasurements provided by the Cologne molecular spectroscopylaboratory and astrophysics group (Brünken et al. 2004). Thesesupra-terahertz absorption lines not only represent the first detec-tion of CH in low excitation states, but they also yield reliabletotal CH column densities, N (CH ), of (7 . ± . × cm − ,which are consistent with chemical model predictions for di ff useclouds.In this paper, we report the detection of the N K a K c = − transitions of CH toward nine SFRs – seven of which CH ortho -CH para -CH T [K] GHz
GHz
GHz
GHz - GHz
GHz
Fig. 1.
Ground state rotational energy level diagram of ortho-, and para-CH up to an energy of 350 K, adapted from Brünken et al. (2004).Transitions measured in the laboratory are labelled with their respec-tive transitions indicated by arrows. Highlighted in red and blue are thetransitions that are discussed in this paper. The corresponding fine- andhyperfine structure transitions are not depicted here as they would notbe visible on this scale. is detected for the first time and attempt to address questionsregarding the origin of its emission. In addition, we also reportthe non-detection of the para-CH , N K a K c = − transitionsbetween 440–445 GHz toward the Orion molecular cloud, whichwere observed with the Atacama Pathfinder Experiment (APEX)12 m sub-millimetre telescope.
2. CH spectroscopy For CH ( B ), which has a spin angular momentum, S , of unity,a state with rigid body angular momentum quantum number, N ,splits into three levels with the total angular momentum quantumnumbers, J , which have values of N + N , or N −
1. Each of theserotational levels ( N K a K c displayed in Fig. 1) splits into three fine-structure levels due to spin-spin and spin-rotation interactions.Furthermore, CH contains two identical protons, each with anuclear spin, I H = /
2, and, consequently, has two nuclear-spinisomers: ortho ( I tot =
1) and para ( I tot = I tot of unity, each of the three fine-structurelevels of ortho-CH further splits into three hyperfine structure(HFS) levels with total angular quantum number, F , of J − J ,and J + J (cid:44)
0, while para-CH states with I tot =
3. Ortho-CH emission at 4.3 mm – A summary ofprior results The 4 − rotational transitions of ortho-CH lie in a rarelyobserved spectral window, close to the 53–66 GHz atmosphericO band and comprise three main fine-structure components J = → , →
3, and 3 → Article number, page 2 of 25. Jacob et al.: Hunting the elusive Methylene radical respectively (which show additional HFS splitting). The corre-sponding frequencies and spectroscopic parameters of all thestudied HFS transitions are summarised in Table 1. In additionto the HFS lines listed, with regard to the ortho-CH rotationaltransitions, for every J → J − ∆ F =
0. The EinsteinA-coe ffi cients for these lines are smaller than those of the lineslisted in Table 1 with ∆ F = ∆ J = ≈ . While these lines were covered in our band pass, nonewere detected; for root mean square (rms) noise levels, see Ap-pendix A. Detected thus far in emission toward hot cores, thefine-structure lines of o-CH and their corresponding HFS linestabulated in Table 1 are naively expected to be weakly populatedgiven the physical conditions that prevail in these regions. In-trigued by the fact that these transitions are observed in emission,Dagdigian & Lique (2018) carried out simple non-LTE radiativetransfer calculations using the one-dimensional escape probabil-ity code, RADEX (van der Tak et al. 2007a). The negative ex-citation temperatures that they obtained suggest that the 4 en-ergy level is selectively enhanced by population inversion andthat all three fine-structure transitions show weak maser emis-sion. However, questions regarding the source of their emissionor the association of these lines with the hot cores could not bedefinitively answered.The Orion Kleinmann-Low (KL) nebula, one of the two re-gions in which Hollis et al. (1995) first detected ortho-CH (here-after o-CH ), is the densest part of Orion Molecular Cloud 1(OMC-1). In itself, it is a highly complex region, which encom-passes: (1) the eponymous ‘hot core’, that is, a compact and verydense, hot >
120 K) region that is surrounded by more extendeddense molecular gas; (2) Orion South (here, Orion S), anotherhot core that, in contrast to the more famous one in KL (Zapataet al. 2011), harbours an embedded young stellar object (YSO)that was first identified as the NH emission peak S6 (Batria et al.1983) and by SiO emission, indicating outflow activity (Ziurys& Friberg 1987), which conclusively proves the presence of anembedded YSO, probably of intermediate mass. In a complexarrangement, Orion S lies in front of an extension of the H II region and the photodissociation region (PDR) associated withthe Orion Nebula (M42) (Mangum et al. 1993); (3) A high den-sity part of this PDR presents the prominent Orion Bar (Walm-sley et al. 2000). To make things complicated, the o-CH spec-tra toward Orion-KL obtained by Hollis et al. (1995) using theNRAO Kitt Peak (KP) 12 m radio telescope in Arizona, with afull width at half maximum (FWHM) beam width of 86 (cid:48)(cid:48) , coversseveral of these components, leaving the origin of the observedCH emission unclear. Given the high energies above the groundstate of these lines’ (upper-level energies of ≈
224 K), it was anatural conclusion to attribute the hot core to the source of theiremission in Orion, whose molecular material was known to becharacterised by temperatures >
120 K or even higher values (seeHermsen et al. 1985, and references therein.).In a series of attempts to elucidate the location and size ofthe CH emission in Orion, we compared the KP 12 m tele-scope results with data taken with two larger telescopes, namely,the Institut de Radioastronomie Millimétrique (IRAM) 30 mand the Green Bank telescope (GBT) operated by the NRAO.We scale the emission expected to be observed with the di ff er- The CDMS contains a full listing of the CH transitions under: https://cdms.astro.uni-koeln.de/cgi-bin/cdmssearch?file=c014501.cat The National Radio Astronomy Observatory is operated by Associ-ated Universities, Inc., under contract by the National Science Founda-tion. ent telescopes by the inverse of the beam filling factor, that is,( θ + θ ) /θ . For the Orion-KL hot core, we estimate a FWHMsource size, θ S , of 12 (cid:48)(cid:48) , from the Very Large Array (VLA) imag-ing of the NH J , K = (3 ,
3) line (Pauls et al. 1983), whose en-ergy levels are 127 K above the ground. For the FWHM beamdiameters, θ B , of the KP 12 m, and the IRAM 30 m telescopesand the 100 m GBT, we assume 86 (cid:48)(cid:48) , 35 (cid:48)(cid:48) and 12 (cid:48)(cid:48) , respectively at ∼
70 GHz. Compared to the KP 12 m telescope, this should resultin 5.5 and 26 times higher main-beam brightness temperatures, T mb , for the IRAM 30 m telescope and the GBT, respectively.In 2004, we carried out observations of the 70.68 GHz fine-structure component of CH using the IRAM 30 m telescopetoward the Orion-KL region. Remarkably, we did not detect theo-CH J = → T mb scale, smoothed to a common velocityresolution of 0.66 km s − . If the o-CH emission was indeedshown to arise from the famous hot core, then we would expectthe line to have been detected with a T mb ∼
140 mK with theIRAM 30 m telescope.Moreover, the CH transitions also remained undetected inthe 67–93.6 GHz spectral line survey presented by Frayer et al.(2015) toward Orion-KL using the GBT, at an rms noise level of ∼
61 mK at 68 GHz and 36 mK at 70 GHz. The 4 mm receiveron the GBT (FWHM ∼ (cid:48)(cid:48) ) was pointed directly at the positionof the Orion-KL hot core. Scaling here also the KP 12 m tele-scope data by the beam filling factor for a compact source witha FWHM of 12 (cid:48)(cid:48) in both the KP 12 m telescope and GBT, weexpect the T mb measured by the GBT to be ∼ .
65 K for the J = → at 68 GHz, that is, 26 timeshigher than that measured using the KP 12 m telescope. Thenon-detections of these CH transitions using the IRAM 30 mtelescope and the GBT suggest that o-CH is simply not presentwithin the dense Orion-KL hot core itself and that, instead, itsemission arises from a more di ff use (and presumably also warm)surrounding medium. This is not surprising as chemical models(Black et al. 1978; Prasad & Huntress 1980; Lee et al. 1996)have previously predicted high abundances of o-CH in both dif-fuse regions as well as in intermediate density gas layers at theedges of dense clouds. This motivated our search for hot o-CH in Orion among regions outside of the hot core.The remainder of this paper is structured as follows. InSect. 4, we describe all our observations of CH , as well asthose of carbon radio recombination lines (CRRLs). We presentthe data in Sect. 5. Here, we start out by quite extensively de-scribing our attempts to clarify the origin of the 68–71 GHz N K a K c = − o-CH emission lines in several diverse envi-ronments spread over ∼ . N K a K c = − sub-millimetre wavelength transitions of para-CH (hereafter p-CH ) that have frequencies between 440 and445 GHz. After this, we present the results of our observationsof the 68–71 GHz o-CH line toward other sources, which alsocontain PDRs, and compare them with the CRRL data. This isfollowed-up by our analysis of the CH data in Sect. 6, wherewe describe our chemistry modelling and radiative transfer anal-ysis. Finally, in Sect. 7 we discuss the main conclusions drawnfrom our results and provide a brief summary of the observa-tional state of the art and future prospects of studies of interstel-lar CH in Sect. 8. Article number, page 3 of 25 & A proofs: manuscript no. main_arxiv
Table 1.
Spectroscopic parameters of the rotational transitions of CH studied in this work. Transition Frequency A E S ij E u J (cid:48) − J (cid:48)(cid:48) F (cid:48) − F (cid:48)(cid:48) [MHz] 1 × − [s − ] [K] N K a K c = − (para-CH )3-4 – 444825.666 59.8 0.928 155.972-3 – 439960.991 64.3 1.353 156.271-2 – 444913.930 70.1 0.631 155.85 N K a K c = − (ortho-CH )5-4 6-5 68371.278 0.216 2.3295-4 ** ** ** Notes.
The spectroscopic data - frequencies, Einstein A -coe ffi cients ( A E ) and line strengths ( S ij ) were taken from Ozeki & Saito (1995) and theCologne Database for Molecular Spectroscopy (Müller et al. 2005). ( ** ) Indicates the central HFS transition, which was used to set the velocityscale in the analysis.
4. Observations
Motivated by the above considerations, we searched for o-CH in hot media with intermediate densities near the envelopes ofhot cores, that is, PDRs, which form at the interface between H II regions and dense molecular clouds and are characterised by thedensity of the cloud and the strength of the far-ultraviolet (FUV)(6 eV < h ν < . emission arisesfrom PDRs, we first re-examined the fine-structure transitions ofo-CH at 68 and 70 GHz and their associated HFS transitionsusing the KP 12 m telescope toward a number of positions inthe OMC-1. Then we followed-up with observations carried outwith the new receiver covering the 4 mm band of the Onsala20 m telescope (Belitsky et al. 2015).The majority of the additional sources toward which we car-ried out our search (and not located in Orion) represent well-known H II regions and giant molecular cloud (GMC) coresshowing active star formation. We targeted positions correspond-ing to peaks identified by a previous mapping of other contin-uum emission and molecular lines. The observed positions co-incide with bona fide H II regions and the substantial distances( ≥ ff set from the fully ionised gas and are also covered inour 52 (cid:48)(cid:48) FWHM single-dish beam. Our study is further supple-mented with observations of p-CH at 444 GHz made using theAPEX 12 m telescope and of CRRLs made using the E ff elsberg100 m telescope. In the following sections, we describe the tech-nical aspects of these observations. The observations of Orion-KL were carried out between Octo-ber and November 2005 (project id: 5029) and between Januaryand April 2006 using the 3 mm receiver of the KP 12 m tele-scope. The dual channel single sideband system was operatedin dual polarisation mode, using the Millimetre Auto Correla-tor (MAC) as backend, which provided a spectral resolution of391 kHz over 8192 channels, spanning an e ff ective bandwidthof 300 MHz. Thus, two frequency setups were required, cen- tred on 68.370 and 70.679 GHz to cover each of the individual5 − − a posi-tion between Orion-KL and S (here referred to as KL / S), the so-called radical-ion peak (RIP) located 4 (cid:48) north of Orion-KL (Un-gerechts et al. 1997), and the Orion Bar, a neighbouring PDR.This Orion Bar position was neither covered by the beam of theKP 12 m, nor that of the IRAM 30 m telescopes in previous ob-servations. Moreover, the 68-71 GHz CH transitions were alsonot covered by the line survey carried out by Cuadrado et al.(2016) between 80–360 GHz using the IRAM 30 m telescope to-ward the same position. In our observations, the ∼ (cid:48)(cid:48) beam ofthe KP 12 m telescope was centred on the source coordinates tab-ulated in Table 2. The spectra were calibrated using a main beame ffi ciency of 0.64. The resultant spectra were then subsequentlyprocessed using the GILDAS-CLASS software (Pety 2005) andup to a second order polynomial baseline was removed. In April and May 2019, we observed the N K a K c = − transi-tions of o-CH using the 4 mm receiver on the Onsala 20 m tele-scope (Walker et al. 2016) (project id: O2018b-07) and followed-up with further observations in January and February 2020. Thenew 4 mm receiver, equipped with a cooled dual-polarisationhigh electron mobility transistor (HEMT) amplifier, was tunedto a frequency of 69.52 GHz, such that all three fine-structurecomponents (with a maximum separation of 2.3 GHz) and theirrespective HFS components could be observed simultaneously,while leaving enough baseline on either side of the 4 GHz IFbandpass. The FWHM beam width at this frequency (69 GHz)was measured to be 52 (cid:48)(cid:48) . The observations were carried out indual beam switch mode with a beam throw of 10. (cid:48)
5. Using the We pointed at a position that is at an o ff set of ( − , + (cid:48)(cid:48) relative tothat used by Tahani et al. (2016) for their line survey, which is negligiblegiven the KP 12 m’s 86 (cid:48)(cid:48) FWHM beam. Software package developed by IRAM, see for more information regarding GILDAS pack-ages.Article number, page 4 of 25. Jacob et al.: Hunting the elusive Methylene radical
Table 2.
Summary of o-CH observations using the KP 12 m, and Onsala 20 m telescopes toward the Orion pointing positions and their corre-sponding line parameters. Source α J2000 δ J2000
Line υ sys υ LSR ∆ υ T mb a rms b S / N c (S / N) tot d [MHz] [km s − ] [km s − ] [km s − ] [mK] [mK]KP 12 m TelescopeOrion-KL 05:35:14.10 -05:22:26.54 70678.633 3.0–6.0 13.4(0.6) 5.2(0.6) 12.7(3.2) 6.5 2.0 4.370679.543 9.5(0.6) 5.2(0.6) 16.6(4.2) 6.5 2.570680.720 4.5(0.6) 5.2(0.6) 21.7(5.4) 6.5 3.3Orion S 05:35:13.10 -05:23:56.00 70678.633 6.5–8 12.4(0.4) 5.2(0.5) 20.1(3.2) 7.0 2.8 6.170679.543 8.5(0.4) 5.2(0.5) 26.4(4.2) 7.0 3.870680.720 3.5(0.4) 5.2(0.5) 34.4(5.4) 7.0 4.968371.278 -13.8(0.4) 4.9(1.0) 35.2(9.6) 7.1 5.068375.875 * / S 05:35:07.60 -05:23:11.00 70678.633 6–12 16.4(0.7) 3.5(0.5) 20.1(4.6) 7.2 2.7 3.370679.543 12.5(0.7) 3.5(0.5) 26.4(6.0) 7.2 3.670680.720 7.5(0.7) 3.5(0.5) 34.4(6.8) 7.2 4.7Orion RIP 05:35:15.80 -05:19:00.50 70678.633 8–10 12.8(0.9) 3.1(0.7) 19.9(3.6) 6.2 3.2 4.070679.543 9.0(0.9) 3.1(0.7) 26.1(4.7) 6.2 4.270680.720 4.0(0.9) 3.1(0.7) 34.0(6.1) 6.2 5.5Orion Bar 05:35:22.80 -05:25:01.00 70678.633 9–10 15.4(0.3) 2.6(0.6) 21.1(6.0) 11.5 1.8 2.870679.543 11.5(0.3) 2.6(0.6) 27.7(8.0) 11.5 2.470680.720 6.5(0.3) 2.6(0.6) 36.1(10.3) 11.5 3.168371.278 -11.4(0.3) 2.1(0.5) 33.1(11.6) 7.8 4.268375.875 * / S (1) 05:35:16.96 -05:22:02.7 70679.633 6–12 19.4(1.8) 4.5(1.2) 23.8(7.5) 14.2 1.7 2.270679.543 15.5(1.8) 4.5(1.2) 31.3(9.7) 14.2 2.270680.720 10.5(1.8) 4.5(1.2) 40.8(12.8) 14.2 2.8Orion-KL / S (6) 05:35:24.96 -05:22:32.7 70679.633 6–12 16.3(1.3) 4.7(0.9) 23.9(5.1) 13.8 1.7 3.170679.543 12.5(1.3) 4.7(0.9) 31.4(6.6) 13.8 2.370680.720 7.5(1.3) 4.7(0.9) 40.8(8.8) 13.8 3.0Orion Bar (2) 05:35:22.80 -05:25:01.0 70679.633 9–10 15.4(1.0) 4.7(1.1) 28.8(14.6) 25.3 1.1 2.370679.543 11.5(1.0) 4.7(1.1) 37.9(19.1) 25.3 1.570680.720 6.5(1.0) 4.7(1.1) 49.3(24.7) 25.3 1.9Orion Bar (5) 05:35:20.81 -05:25:17.1 70679.633 9–10 15.4(0.8) 5.0(1.0) 28.8(8.4) 15.0 1.9 3.870679.543 11.5(0.8) 5.0(1.0) 37.9(11.4) 15.0 2.570680.720 6.5(0.8) 5.0(1.0) 49.3(14.5) 15.0 3.3
Notes. ( a ) Peak main-beam brightness temperature derived from the integrated intensity of the detected line features. ( b ) The rms noise level on the T mb scale, quoted for a spectral resolution of 0.85 km s − and 0.97 km s − for observations made using the KP 12 m, and Onsala 20 m telescopes,respectively. ( c ) The signal-to-noise ratio (S / N) with respect to the peak line temperature. ( d ) The S / N of the integrated intensity calculated from thelisted rms noise levels and line widths for the 70 GHz CH line. ( * ) Represents the HFS component used to determine the LSR velocity axes of thespectra shown in Fig. 2. For each HFS line, the velocities of the other HFS components are listed as they appear in the spectra shown in this figure.
References.
For the radial velocities, see Gong et al. (2015) and references therein. − . Timely pointing and focus ac-curacy checks were performed by observing nearby stellar SiO( (cid:51) = , J = →
1) masers.The intensity calibration was done every 12 minutes usingthe standard chopper-wheel method, whereby the second orderchopper-wheel correction term was calculated and applied fol-lowing Ulich & Haas (1976). We express the intensity scale ofour spectra in units of T mb , by assuming a main beam e ffi ciencyof 0.55 (on average), while the velocity scale is given with re-spect to the local standard of rest. Similarly to the KP 12 m data,the calibrated spectra obtained using the Onsala 20 m telescopewere further analysed using the CLASS software. However, we see signatures of a regular standing-wave pattern with a fre-quency of ∼
16 MHz, which are likely associated with the enclos-ing radome structure whose reflective properties worsen at thelower end of the 4 mm band. This is particularly so for the obser-vations toward Orion, whose transit at the Onsala 20 m telescopeis much lower in comparison to that at the KP 12 m telescope.We corrected for contributions from the standing-wave featuresby using a standard standing-wave removal method based on aFast Fourier Transform analysis. The resulting spectra are thenbox-smoothed to velocity bins of ∼ − and polynomialbaselines up to the third order were subtracted. Article number, page 5 of 25 & A proofs: manuscript no. main_arxiv
Table 2.
Continued.: Summary of o-CH observations using the Onsala 20 m telescope and derived line parameters toward the other sourcespresented in this study. Source α J2000 δ J2000
Line υ sys υ LSR ∆ υ T mb a rms b S / N c (S / N) tot d [MHz] [km s − ] [km s − ] [km s − ] [mK] [mK]W3 IRS5 02:25:40.5 62:05:52 70678.633 -39.0[1] -35.1(0.5) 7.3(0.9) 22.9(6.4) 9.8 2.3 5.070679.543 -39.0(0.5) 7.3(0.9) 30.0(8.4) 9.8 3.170680.720 -44.0(0.5) 7.3(0.9) 39.1(9.3) 9.8 4.068371.278 -60.4(0.3) 5.4(0.8) 21.6(4.0) 4.4 4.968375.875 * -40.1(0.3) 6.2(0.6) 25.0(3.3) 4.5 5.668380.973 -19.0(0.3) 5.5(0.5) 27.5(3.4) 4.5 6.269007.179 e -59.3(0.4) 8.7(0.9) 8.2(2.7) 3.0 2.869014.202 * -39.0(0.4) 6.9(1.3) 13.3(2.8) 3.0 4.469019.187 -7.5(0.4) 9.7(2.2) 9.9(2.9) 3.0 3.3W51 E 19:23:44.0 14:30:30 70679.633 + + + + + + + + + + + + + + + + + + + + − . + + + + + − . Notes. ( a ) Peak temperature derived from the integrated intensity of the detected line features. ( b ) The rms noise level on T mb scale, quoted for aspectral resolution of 0.97 km s − for observations made using the Onsala 20 m telescope. ( c ) The signal-to-noise ratio (S / N) with respect to thepeak line temperature. ( d ) The S / N of the integrated intensity calculated from the listed rms noise levels and line widths for the 70 GHz CH line. ( e ) The reported intensity for this component accounts for contributions from the blended NS J , F = / , / → / , / J , F = / , / → / , / ( * ) Represents the HFS component used to determine the LSR velocity axes of the spectra shown in Fig. 8. For each HFS line, the velocities of theother HFS components are listed as they appear in the spectra shown in this figure.
References.
For the radial velocities: [1] Imai et al. (2000); [2] Parsons et al. (2012); [3] Jackson & Kraemer (1994); [4] Bally et al. (2010);[5] Dickel et al. (1978); [6] Bally et al. (2002).
The fine-structure components of the N K a K c = − transi-tion of p-CH were observed in 2013 August , using the highfrequency channel of the sideband separating (2SB), dual fre-quency band First Light APEX Submillimetre Heterodyne Re-ceiver, FLASH (hereafter FLASH-460, Heyminck et al. 2006)on the APEX 12 m sub-millimetre telescope (Güsten et al.2006). The FWHM beam size at 443 GHz is 14 (cid:48)(cid:48) . The band-pass was selected such that we covered the 444 GHz transitions Project id: M-091.F-0040-2013 APEX is a collaboration between the Max-Planck-Institut fur Ra-dioastronomie, the European Southern Observatory, and the OnsalaSpace Observatory. in the upper sideband alongside the CO J = → ffi ciency of 0.95 and a main beam e ffi ciency of0.60. Polynomial baselines up to a second order were removedand the subsequently obtained spectrum was box-smoothed tochannel widths of 1 km s − . Each pointing position was inte- Article number, page 6 of 25. Jacob et al.: Hunting the elusive Methylene radical grated on, for a total time of ∼
23 mins, except for the Orion Sposition, toward which we carried out deeper integration , for atotal time of 12.7 hours. We performed CRRL measurements (project id: 08-19) towardthose targets with successful detections of CH , in position-switching mode with the S20mm receiver of the 100-m telescopeat E ff elsberg, Germany , on 2019 5–6 July and 22–23 August.The S20mm receiver is a double-beam and dual-polarisation re-ceiver operating in the frequency range between 12–18 GHz.This range contains ∆ n =
1, ‘ α ’, radio recombination lines fromH, He and C with principal quantum numbers, n , between 80 and72; see Appendix C. Frequencies of RRLs from all three speciescan be calculated following the prescriptions of Lilley & Palmer(1968), or simply retrieved from the Splatalogue database .For our analysis, we only use data from the central beam ofthe S20 receiver. The FFTS (e.g. Klein et al. 2012) data serveas backend, each of which consists of 65536 channels. For theobservations toward the di ff erent Orion positions and W3 IRS5,we used a total bandwidth of 300 MHz with a channel widthof 4.6 kHz, corresponding to a velocity spacing of 0.09 km s − at 15 GHz. Our observations of the other targets utilised a totalbandwidth of 5 GHz and a channel width of 38.1 kHz, corre-sponding to a velocity spacing of 0.76 km s − at 15 GHz. The fo-cus was adjusted using observations of strong continuum sourcesat the beginning of each observing session. Pointing observa-tions were carried out roughly every two hours toward strongcontinuum sources nearby. The pointing accuracy was found beless than 5 (cid:48)(cid:48) . NGC 7027 was used as the flux calibrator, and theflux calibration accuracy is estimated to be within ∼ (cid:48)(cid:48) at 15 GHz. This angular resolution allows for ameaningful comparison with the CH transitions observed us-ing the Onsala telescope at ∼
69 GHz (with a FWHM of 52 (cid:48)(cid:48) ).The main beam e ffi ciency is about 0.65 and the typical systemtemperature is about 15 K. The data reduction was once againperformed using the CLASS software.
5. Results
The lines of the HFS triplet of the J = → overlap with one another because of their closefrequency spacing, ∆ ν , of 0.88 and 1.18 MHz, correspondingto 3.7 and 5.0 km s − , respectively. This blending of the indi-vidual components with one another broadens the observed pro-file. While this greatly aids in the detection of this component,it does not reveal accurate line properties, such as the intrin-sic line width. On the other hand, with a frequency separationof 9.6 and 12 MHz, respectively, the individual HFS compo-nents of the J = → J = → near68 GHz and 69 GHz, respectively, are well resolved. However,the F = → E u = . J , F = / , / → / , / Project id: M-091.F-0045-2013 The 100-m telescope at E ff elsberg is operated by the Max-Planck-Institut für Radioastronomie (MPIFR) on behalf of the Max-PlanckGesellschaft (MPG). https://splatalogue.online// fore, the HFS-resolved o-CH line profiles are modelled by si-multaneously fitting individual Gaussian profiles to the observedHFS lines using the CLASS software. The relative contributionof the emission from the NS transition present in the strongestHFS component of the J = → F = / → / J = → , we mod-elled the observations by using the line widths derived from the J = → ff ected byblending and contamination, and thereby we were able to revealthe true shape of the CH emission profile. While the thus de-termined line width can be used to model the line profiles of theother fine structure lines, this line is unfortunately not detectedtoward a majority of our sources above a 3 σ noise level. The rmsnoise levels we obtained for the di ff erent sources are listed inAppendix A. For those sources toward which the J = → J = → by minimising the mean square error betweenthe modelled fit and the observations over several iterations cov-ering a range of line widths typically between 2 and 10 km s − insteps of 0.23 km s − . This ‘empirical’ model simultaneously fitsthe three HFS lines using Gaussian profiles with: (1) positionsdetermined by the velocity separation of the HFS lines relativeto the 70.679543 GHz component; (2) one common line width;and (3) individual line intensities estimated by scaling the peaktemperature by the line strength of each HFS line. Since the typ-ical positions are well-known for a given source and the line in-tensities can be constrained from the observations, the line widthis the only free parameter. Varying the line width in each itera-tion and comparing the mean square error between the observedline profile and that of the combined HFS fit, the scheme con-verges for the line width returning the least error. The errors inthe line parameters are determined from the covariance matrixand depend on the rms noise of the system. There are additionalsystematic uncertainties in the line width, which are caused bythe low signal-to-noise ratio (S / N) in these spectra. A summaryof the derived line parameters is given in Table 2. in Orion The spectra observed toward the di ff erent Orion positions withthe KP 12 m telescope are displayed in Fig. 2. We detect the J = → − transitionof CH near 70 GHz toward all the Orion pointing positions atsystemic velocities between 8 and 10 km s − . A second observa-tional setup was used toward the Orion S and Bar positions, withthe band centred on the J = → (cid:38) σ level.In the spectrum of the 68 GHz component toward Orion S, wealso detect part of the K = J = → CCH near 68.3649 GHz.While the measured line intensity toward Orion-KL is con-sistent with that of the original detection by Hollis et al. (1995),we find the line intensities derived toward the other Orion posi-tions to be of very comparable, if not greater, strengths than thestrongest emission being seen toward the Orion S position. Thesuccessful detection and higher CH intensities observed toward This analysis was carried out using Python packages numpy andscipy (Harris et al. 2020). Article number, page 7 of 25 & A proofs: manuscript no. main_arxiv regions outside of the hot core once again brings into questionthe association of the CH emission with the hot core. Moreover,we detect CH emission toward an intermediate position be-tween Orion-KL and S, which we name Orion KL / S and towardwhich we would not expect any CH emission if it were confinedto the hot core. The detection of this high-lying CH transitiontoward all the Orion positions with comparable line strengths, in-dicates an extended emission component of CH arising from arather dilute, but hot ISM component and also suggests that thismolecule is simply not found within the hot core. The compati-bility of the intrinsic velocity of the observed CH emission to-ward Orion KL with that of the extended ridge component, υ LSR = − (see Gong et al. 2015, and references therein),which represents the typical velocities of the ambient gas presentin this region and further hints at the extended nature of the CH emission. Since these observations were once again carried outusing the KP 12 m telescope with a large beam size, the resultingsignal may certainly contain contributions from neighbouring re-gions located within the beam.In an attempt to address the question of where the CH emis-sion arises from, we studied the emission characteristics of thismolecule at several positions within the Orion complex, usingthe Onsala 20 m telescope, which has a beam size of 54 (cid:48)(cid:48) at69 GHz. These results are discussed in Sect. 5.3. For conve-nience, the pointing positions chosen for our study are groupedinto two, namely, the Orion KL / S and Orion Bar regions, re-spectively. The di ff erent positions are classified based on theirproximity to either the hot core, Orion-KL, or to the PDR, theOrion Bar. emission The Orion KL/S region
In addition to the nominal Orion-KL and Orion S positions usedin our KP observations (hereafter referred to as Orion KL / S posi-tions (2) and (5)), we also observe four positions correspondingto the massive O7V star θ Ori C, position (4), which is respon-sible for the ionisation of the Orion Nebula, [C II ] peak positions(1) and (6) selected from Pabst et al. (2019) and a position be-tween the Orion-KL and S regions, position (3). In Fig. 3, wedisplay the calibrated and baseline-subtracted spectra observedtoward the di ff erent Orion KL / S positions. We detect CH emis-sion toward the two [C II ] peak positions marked (1) and (6), ata ∼ σ level with an average rms noise of 14 mK for the twospectra. However, we see no clear signal toward the other OrionKL / S positions, above an rms noise level of 22 mK on average.For the positions toward which we do not clearly detect CH , westack and then average the spectra (taken three at a time) in orderto reduce the spectral noise and investigate our suspicion that theCH emission is extended. The resulting spectra are displayed inthe bottom panel of Fig. 3. Upon scaling the fit results obtainedtoward positions (1) and (6), we see that all four independentcombinations of the stacked and averaged profiles hint at thepresence of a signal at velocities close to the systemic velocity ofthe source. The strongest CH signals are observed in the profilesthat result from the combination of positions (2, 3, 5) and (3, 4,5). This suggests that the CH emission is weakest at positions(2) and (4) which correspond to the Orion-KL hot core and the θ Ori C positions, respectively. Cross-correlating the intensitiesintegrated over a velocity range from −
10 to +
20 km s − (whichis roughly the velocity interval over which we expect to observeCH signatures) obtained toward the individual, and combinedsets of pointing positions amongst each other, we find that po- sitions (2) and (4) show negative correlations toward almost allother sources. The negative and / or weak correlations (tending tozero) shown by these components (see Fig. 4) suggest that thereis no association between them and reveal that CH is likely notpresent within the hot core. The correlations we present werecomputed using the Pearson product moment correlation coef-ficient. This coe ffi cient describes the strength of the linear re-lationship between each pair of spectra by using the standarddeviation of each data set and the covariance between them. Theunderlying assumption made in this calculation is that the datafollow a Gaussian distribution. For comparison, we also com-puted the correlation strengths based on non-parametric statisticsby using the Spearman correlation coe ffi cient. While the relativestrengths produced by both methods are di ff erent, they both re-produce the same monotonic trends in their correlations. Sincethe aim of this analysis is to simply distinguish the nature of thecorrelation, namely positive or negative, in all subsequent analy-ses that make use of correlation coe ffi cients, we use the standardPearson coe ffi cient.The detection of CH toward the hot core initially discussedin Hollis et al. (1995) and also in this work, can now be explainedthrough its successful detection in the Orion KL / S (1) position.O ff set from the nominal Orion KL position by (40.2 (cid:48)(cid:48) , 24.8 (cid:48)(cid:48) ), apart of Orion KL / S (1) is covered by the KP 12 m beam. Thisresolves the observational discrepancies between the KP 12 mtelescope, the GBT, and the IRAM 30 m telescope (detailed inSect. 3) and indicates that the CH emission is associated withgas layers similar to those traced by [C II ]. Sensitivity is of courseanother important factor. As discussed earlier, given its extendedemission, we would expect to observe CH with a similar linestrength using the di ff erent telescopes. The non-detections usingthe IRAM 30 m telescope and GBT therefore exclude the pos-sibility that the CH emission originates in compact regions likehot cores. The Orion Bar region
In addition to the Orion Bar position previously observed usingthe KP 12 m telescope, which corresponds to the peak of theCO emission [hereafter known as Orion Bar position (2)], wecarried out observations toward four additional positions. Theycorrespond to positions at or near the emission peaks of HCN,CF + , C H, and CO + and are labelled (1), (3), (4), and (5), respec-tively, as displayed in Fig. 5. The positions were selected on thebasis of previous studies by Neufeld et al. (2006), Stoerzer et al.(1995), Cuadrado et al. (2015), and Nagy et al. (2015b) and ref-erences therein. CH emission is detected at a 2.3 and 4 σ leveltoward the Orion Bar nominal position (2) and the C H emis-sion peak at position (5) at an rms noise level of 25 and 15 mK,respectively. We do not detect any appreciable signal from theother Orion Bar positions, even toward positions (1) and (3) atrms noise levels down to 17 mK.Similar to the pointing positions toward the Orion KL / S re-gion for which we do not detect CH emission, we stack and av-erage the positions with non-detections in the Orion Bar region.The combined profiles of the independent pairs between posi-tions (1), (3), and (4), as well as that obtained when consideringall three together, are displayed in the bottom panel of Fig. 5.Scaling the fit parameters obtained from position (2), we findthat there is a weak indication of CH emission present in all thecombinations. Carrying out a correlation analysis amongst thedi ff erent Orion Bar positions we do not see any anti-correlations.However, from the correlation matrix presented in Fig. 6 we infer Article number, page 8 of 25. Jacob et al.: Hunting the elusive Methylene radical
R. A. (J2000) D e c . ( J ) [K kms ]50 0 50 LSR T m b J = 4 370679 MHz
Orion RIP [ K ] [kms ] T m b [ K ] J = 4 370679 MHz
Orion KL/S
J = 4 370679 MHz
Orion S
50 0 50
LSR
J = 5 468375 MHz
Orion S [kms ] CH CCH T m b J = 4 370679 MHz
Orion KL [ K ] J = 4 370679 MHz
Orion Bar
50 0 50
LSR
J = 5 468375 MHz
Orion Bar [kms ] Fig. 2.
Integrated intensity map of the [CII] 158 µ m emission overlaid with integrated intensity contours (orange) from 600 to 1000 K km s − insteps of 100 K km s − toward OMC-1, adapted from Pabst et al. (2019). The positions observed using the KP 12 m telescope are marked by bluecrosses, while the blue circles represent the KP beam. The corresponding calibrated and baseline subtracted o-CH spectra are displayed alongsidethe map. The individual fits to the HFS components are displayed by dotted blue curves while solid red curves displays the combined fit. Thepositions and relative intensities of the HFS components are marked below the spectra. The dashed blue line indicates the systemic velocities ofthe sources which lie between + +
11 km s − . In light blue we also mark the CH CCH lines covered in the 68375 MHz spectrum towardOrion S. The velocity scale is set by the F (cid:48) − F (cid:48)(cid:48) = J – J − that position (3) has the weakest correlation coe ffi cients particu-larly with that of positions (1), (4), and (5).The observed CH emission in the Orion Bar suggests thatthe molecule’s abundance decreases as we move away from theionisation front (near the H II region) and toward the molecularclouds deeper within the PDR. Observations of HF emission byKavak et al. (2019) across the Orion Bar centred near the CO + peak reveals a similar morphology. These authors were able toshow that the bulk of the HF emission peaked in a region sepa-rating the H and [C II ] emission from the molecular emission inthe denser clumps or close to the ionisation front. A direct com-parison between the CH and HF emissions is di ffi cult becauseboth sets of observations were carried out toward di ff erent posi-tions in the Bar. However, since emission from both species hasbeen observed near the Orion Bar CO + emission peak, we cancompare the CH line profile with that observed for HF at thisposition (this corresponds to position (2) in Kavak et al. (2019)and position (5) in this work). Both species show peak emissionat velocities and line widths that are consistent within the errorbars, between 10–10.5 km s − and 4.4–5 km s − , respectively.From their line emission survey over the entire range of fre-quencies o ff ered by the Herschel Heterodyne Instrument for theFar Infrared (HIFI) under the HEXOS HIFI key guaranteed timeproject, Nagy et al. (2015a) inferred typical line widths between2 and 3 km s − toward the Orion Bar CO + peak, with the excep-tion of a handful of species including HF. The broader line widthof HF was reconciled as being due to its association with di ff usegas present in the inter-clump regions of Orion Bar. Therefore, it is conceivable that o-CH is also likely to arise from a similarcloud population tracing dilute but hot ( T kin =
120 K) gas, un-like most of the other species studied by Nagy et al. (2015a) andothers, including Cuadrado et al. (2015, 2017), which trace gasdensities between 10 and 10 cm − . However, it is di ffi cult toextend such a comparison toward the other sources in our studybecause HF is typically observed in absorption toward the enve-lope of these molecular clouds even showing several absorptioncomponents along the LOS. in Orion The N K a K c = − transitions of p-CH between 440 and445 GHz that have upper level energies of 156 K. Lying within asub-millimetre window, they are accessible from high-mountainsites but have gone undetected thus far. Observations of thesep-CH lines could potentially aid our understanding of CH ’sexcitation. We do not detect any sign of the 444 GHz p-CH transitions toward the di ff erent Orion positions given in Sect. 4.3above an rms noise level of 77 mK on average, for a spectral res-olution of 1 km s − . Moreover, even after a deeper integrationtoward the Orion S position, which resulted in an rms level of8 mK, we did not detect any signatures of p-CH . The Einstein A -coe ffi cients and hence, the critical densities of the 444 GHzp-CH lines are two orders of magnitude larger than that ofthe 70 GHz o-CH lines. With critical densities on the order of2 × cm − the non-detection of these lines is no surprise and Article number, page 9 of 25 & A proofs: manuscript no. main_arxiv
R. A. (J2000) D e c . ( J ) Orion KL/S [K kms ] 0.000.050.10 J = 4 370679 MHz
Orion KL/S (1) T m b J = 4 370679 MHz
Orion KL/S (2) [ K ]
50 0 50
LSR
J = 4 370679 MHz
Orion KL/S (3) [kms ] LSR [kms ] J = 4 370679 MHz
Orion KL/S (6)
LSR [kms ] T m b J = 4 370679 MHz
Orion KL/S (4) [ K ]
50 0 50
LSR
J = 4 370679 MHz
Orion KL/S (5) [kms ]
50 0 50
LSR T m b J = 4 370679 MHz (2, 3, 4) [kms ] [ K ]
50 0 50
LSR
J = 4 370679 MHz (2, 3, 5) [kms ]
50 0 50
LSR
J = 4 370679 MHz (2, 4, 5) [kms ]
50 0 50
LSR
J = 4 370679 MHz (3, 4, 5) [kms ] Orion KL/S : Stacked Averaged Spectra
Fig. 3.
Top: Same as Fig. 2 but for o-CH observations made using the Onsala 20 m telescope toward the Orion KL / S region. Blue circles representthe beam size of the Onsala 20 m telescope centred at the di ff erent pointing positions and the dotted white circle marks the KP beam at theOrion-KL and Orion S positions marked in Fig. 2. Bottom: Stacked and averaged o-CH spectra obtained by combining the Orion KL / S pointingpositions (2), (3), (4), and (5), taken three at a time. The line fits are scaled using the line parameters obtained toward Orion KL / S positions (1)and (6). ( ) ( ) ( ) ( ) ( ) ( ) ( , , ) ( , , ) ( , , ) ( , , ) (1)(2)(3)(4)(5)(6)(2,3,4)(2,3,5)(2,4,5)(3,4,5) -0.340.42 -0.150.12 -0.19 0.260.12 -0.28 0.14 0.380.41 -0.068 0.47 0.078 0.190.12 0.37 0.66 0.65 0.15 0.290.14 0.3 0.61 0.32 0.63 0.38 0.73-0.041 0.25 0.16 0.71 0.71 0.13 0.67 0.750.3 -0.3 0.63 0.76 0.75 0.34 0.66 0.74 0.75 C o rr e l a t i o n c o e ff i c i e n t Orion KL S6
Fig. 4.
Pearson product-moment correlation coe ffi cients between theintegrated o-CH intensities, observed toward the individual and com-bined pointing positions in the Orion KL / S region. consistent with our finding that CH exclusively resides in hot,but low-density regions. in other sources The results of our observations discussed thus far point to theorigin of the o-CH emission in regions of intermediate gas den-sities in the envelopes of hot cores, probing PDR layers, ratherthan the hot cores themselves. In order to confirm the associa-tion of the observed o-CH emission with PDRs, we have alsosearched for o-CH emission in: (1) other well-known SFRs thatharbour PDRs, (2) (proto-) planetary nebulae (PNe) that are sur-rounded by molecule-rich envelopes that resemble the compo-sition of PDRs and (3) supernovae remnants (SNRs). Their co-ordinates are listed in Tables 2, A.1, and A.1, alongside theirassumed centroid LSR velocities and, for each line group, therms noise levels.We successfully detected the blended J = → in emission toward W51 M, N, and E, W49 N, W3 IRS5,W43, W75 N, DR21, and S140 at the systemic velocities of thesesources. The resulting spectra are displayed in Figs. 7–9. In thefollowing sections, we discuss the observed characteristics ofthis line blend for select sources in more detail. W51
Given that W51 Main (M) was one of the original targets towardwhich Hollis et al. (1995) first detected CH , we re-observedthis position in order to first verify their detection and then car-ried out observations toward two luminous condensations har- Article number, page 10 of 25. Jacob et al.: Hunting the elusive Methylene radical
R. A. (J2000) D e c . ( J ) Orion Bar [K kms ] 0.050.000.050.10 J = 4 370679 MHz
Orion Bar (2) T m b J = 4 370679 MHz
Orion Bar (4) [ K ] LSR T m b [ K ] J = 4 370679 MHz
Orion Bar (1) [kms ]
50 0 50
LSR
J = 4 370679 MHz
Orion Bar (3) [kms ]
50 0 50
LSR
J = 4 370679 MHz
Orion Bar (5) [kms ]
50 0 50
LSR T m b J = 4 370679 MHz (1, 3) [kms ] [ K ]
50 0 50
LSR
J = 4 370679 MHz (1, 4) [kms ]
50 0
LSR
J = 4 370679 MHz (3, 4) [kms ]
50 0 50
LSR
J = 4 370679 MHz (1, 3, 4) [kms ] Orion Bar : Stacked Averaged Spectra
Fig. 5.
Top: Same as Fig. 3 but for o-CH observations made using the Onsala 20 m telescope toward the Orion Bar region. The blue circlesrepresent the beam size of the Onsala 20 m telescope centred at the di ff erent pointing positions and the dotted white circle marks the KP beam atthe Orion Bar position marked in Fig. 2. Bottom: Stacked and averaged o-CH spectra obtained by combining the Orion KL / S pointing positions(1), (3), and (4), two at a time and all together. The line fits are scaled using the line parameters obtained toward Orion Bar positions (2) and (5). ( ) ( ) ( ) ( ) ( ) ( , ) ( , ) ( , ) ( , , ) (1)(2)(3)(4)(5)(1,3)(1,4)(3,4)(1,3,4) C o rr e l a t i o n c o e ff i c i e n t Orion Bar
Fig. 6.
Pearson product-moment correlation coe ffi cients between theintegrated o-CH intensities, observed toward the individual and com-bined pointing positions in the Orion Bar region. boring high-mass YSOs present in this region, W51 North (N),also known as W51 IRS 2, and W51(E). The W51 cloud complex, lying in the Sagittarius spiral armat a distance of 5.4 kpc from the Sun (Sato et al. 2010), is oneof the best-studied SFRs in our Galaxy. W51 E and N are thetwo active and presumably youngest centres of activity, host-ing the ultracompact (UC) H II regions e1–e8 and d, respec-tively (Ginsburg 2017, and references therein), infrared and sub-millimetre continuum emission (Thronson & Harper 1979; Ja ff eet al. 1984), H O and OH masers (Genzel et al. 1981), and knotsof hot NH ( J , K ) = (3,3) line-emitting gas (Ho et al. 1983; Goddiet al. 2015; Ginsburg 2017), all of which are signposts of ac-tive star formation within cores with masses of (cid:38) M (cid:12) each.W51 N and E are separated by a projected distance of 1. (cid:48)(cid:48) J = → − transition of o-CH near70 GHz toward all three positions at the source intrinsic veloci-ties between 57 and 62 km s − (see Fig. 7). While the emissionprofiles toward each of the three positions are comparable, thestrongest emission arises from the W51 M region, toward whichHollis et al. (1995) pointed the KP 12 m telescope. This is notsurprising, given that the radio emission (and that from its PDR,too) has a larger angular size than that from E and N, resultingin a larger beam filling factor. Therefore, the fact that we seestronger CH emission toward the extended H II region W51 M Article number, page 11 of 25 & A proofs: manuscript no. main_arxiv than toward the much more massive and denser W51 E and Nregions, strongly supports an origin of CH in extended dilutegas.The observed line profile of the 70 GHz component towardW51 M is comparable to that reported by Hollis et al. (1995)however, we do not detect distinct emission from the HFS linescorresponding to the J = → J = → ff erent W51 pointing positions byaligning their frequency scales. Due to the contamination fromstrong NS emission, it is di ffi cult to disentangle the HFS fea-tures corresponding to the 69 GHz transition. However, the HFSlines of the 68 GHz component are weakly visible. In order togauge the true nature of this emission, we further stacked each ofthe individual HFS lines by aligning their velocity scales. Thisstacking exercise revealed a 3.5 σ detection of this line with awidth of 9.1 km s − and a peak line temperature of 23 mK. Theline width determined from this 68 GHz component is consis-tent with the value of the intrinsic line width derived from theiterative HFS decomposition of the 70 GHz lines. The spectrumresulting from our stacking analysis is displayed alongside the70 GHz lines in Fig. 7. This highlights the fact that the non-detection of the 68 GHz and 69 GHz transitions of CH in ourobservations is primarily due to a sensitivity issue. W3 IRS5
The W3 IRS5 cluster system has a well known double IR sourceat the centre of an embedded cluster of a few hundred low massstars (Megeath et al. 1996). Located in the W3-Main region at adistance of 2 . × L (cid:12) (Campbell et al. 1995), itis considered to be at the early stages of star formation. We de-tect the blended HFS transitions of o-CH near 70 GHz towardthis region in emission, centred around -40 km s − (with a S / N > ff erent W51 positions, we did not achieve noise levelsthat are low enough to clearly detect the 68 GHz, and 69 GHzfine-structure lines. However, by stacking our data with a deepintegration spectrum of a bandpass covering 67.3–69.8 GHz, ob-tained by one team member during the course of a di ff erent studyusing the OSO 20 m telescope toward this source, we were ableto detect both the 68 GHz and 69 GHz transitions of o-CH andtheir respective HFS lines as well.We find the HFS transitions of the 68 GHz fine-structureline to be well-resolved with peak temperatures between 21 and27 mK and line widths of 5.3–6 km s − , which are compara-ble with one another, as well as the HFS-stacked 68 GHz com-ponent, toward the combined W51 positions. As discussed inSec. 5.1, the strongest HFS component of the 69 GHz line blendswith the F = / → / F = / → / A -coe ffi cients and upper leveldegeneracies, we can decompose the relative contribution of theNS line from that of o-CH under the assumption that they areoptically thin for conditions of LTE. The o-CH spectra for boththe 68 GHz, and 69 GHz transitions are displayed in Fig. 8 alongwith the modelled fit to the contamination from NS. Other sources
In addition to the results presented in the above sections, wesuccessfully detected o-CH emission from the HFS blended J = → ≥ σ level. The observed spectra are dis-played in Fig. 9 and the results are tabulated in Table 2.At first glance, it appears surprising that the observed lineintensities in the nearby Orion region, which is a distance of ≈
400 pc (Menten et al. 2007; Kounkel et al. 2017), are compara-ble to the values we find for the other more distant regions, whichare between 1.4 kpc (DR21) to 11 kpc (W49 N) away from theEarth. If all sources (including Orion) were unresolved and hadan identical intrinsic size and line luminosity, then ultimately allof their emission would be detected in our beam and their mea-sured intensities would scale with the beam-filling factor. In thatcase, emission from Orion would be by far the strongest. In con-trast, our results indicate that the CH emission from Orion isvery extended and has a very low surface brightness. This meansthat in contrast to the other more distant sources, in the case ofOrion, our beam only samples a small portion of the CH emis-sion, resulting in a low intensity in some PDR positions, or evena non-detection in others, although CH may be present overmuch of the volume of the PDR. The detection of o-CH toward SFRs strongly suggests that CH resides in all cases, as in Orion, in warm intermediate-densityregions, namely PDRs, which surround the denser, fully molec-ular material harbouring embedded YSOs. Indeed, the HFS-decomposed line profiles of the observed CH spectra show LSRvelocities and line widths that are similar to those of previ-ously observed prime PDR tracers, namely low frequency CR-RLs and neutral atomic carbon lines; see, for example, Heileset al. (1996); Wyrowski et al. (1997); Roshi et al. (2006) andJakob et al. (2007).Having spatial distributions, line widths, and radial velocitiesconsistent with an origin in the neutral gas close to the C + / C / COtransition layer, CRRLs are particularly useful tools for probingthe physical conditions and kinematics of these regions (Hoang-Binh & Walmsley 1974; Natta et al. 1994; Salas et al. 2019).Since the properties derived from the observed CRRL line pro-files reflect the physical conditions of the PDR, coupling our ob-servations of o-CH with ancillary CRRL data will help us toconstrain the origins of the observed o-CH emission.In the following analysis, we compare the profiles of HFS-stacked 70 GHz lines with those observed for CRRLs. Becauseof the small velocity (frequency) separation between the HFSlines corresponding to the 70 GHz fine-structure transitions, theyappear to be blended, which makes it di ffi cult to simply stackthem. For this reason, we use the modelled results from the in-dividual HFS fits for stacking. The expected rms noise of thestacked and averaged spectrum is then added back to the mod-elled HFS-stacked profile. The di ff erent steps involved in thisexercise are detailed and illustrated in Appendix B.In the following, as a basis for a CH / CRRL comparison, weuse a PDR model to explore the abundance versus visual extinc-tion ( A v ) profiles of CH , ionised and neutral carbon (C + andC , respectively), and other species. We then compare the modelresults with observational constraints derived from the line pro-files.Our models were created using a simple Python-based PDRcode, PyPDR (Bruderer 2019). The code computes chemical Article number, page 12 of 25. Jacob et al.: Hunting the elusive Methylene radical
R. A. (J2000) D e c . ( J ) W51E W51M W51N [Jy] LSR T m b J = 4 370679 MHz
W51N [km/s] [ K ] T m b J = 4 370679 MHz
W51M [ K ]
100 50 0 50 100 150 200
LSR T m b J = 4 370679 MHz
W51E [ K ] [kms ] LSR T m b W51 stacked [ K ] [kms ] Fig. 7.
Left: Spectra of the o-CH J = → ∼ +
60 km s − . Centre: Ku-band image (14.5 GHz, tracing ionised gas) of the W51 region obtained using acombination of JVLA B and D arrays taken from Ginsburg et al. (2016), is used to indicate the di ff erent Onsala pointing positions. The distancebetween the di ff erent positions is greater than half the beam size of 27 (cid:48)(cid:48) at 69 GHz. Right: The position- and HFS-stacked spectrum of the 68 GHzCH lines. The strong emission at (cid:51) >
77 km s − represents contamination from a CH CCH line (see Sect. 5.2). abundances by evolving chemical rate equations iteratively, byutilising a pure gas-phase, time-dependent chemical network(except for H ), containing 30 species, including CH . The in-put conditions used for the model are UV radiation field, G , inHabing units, the total gas density, n H , and the cosmic-ray ioni-sation rate, ζ H . We studied two sets of models with gas densities n H of 10 and 10 cm − , each of which was exposed to a UVradiation field of 1 and 10 . All four models were exposed to thesame cosmic-ray ionisation rates, fixed at ζ H = . × − s − ,which corresponds to the typical value expected in di ff use andtranslucent clouds (Indriolo et al. 2015; Neufeld & Wolfire 2016;Jacob et al. 2020). The resulting (normalised) abundances of rel-evant carbon-chain species are displayed in Fig. 10 as a functionof A v . In both models with G =
1, the CH abundance peaksin the transition layer from C + to C with its distribution peak-ing at A v (cid:46)
1. The models with a higher UV radiation field showthe CH abundance distribution to follow more closely that ofneutral atomic carbon. From both sets of modelled results, it isclear that CH traces gas layers between the dissociation front(which marks the transition from H → H ) and the molecularcloud, tracing gas layers where C + → C and CO is not the mainreservoir of carbon. Overall, the modelled results are consistentwith our premise that both CH and CRRLs are probing similarcloud layers.In Appendix C, we present a complete summary of the ob-served CRRLs; however, as an example, we display the observedrecombination line spectrum toward W3 IRS5 in Fig. 11. Com-paring the narrow line profile of the CRRL with that of thecorresponding broader ( ∆ υ (cid:38)
30 km s − ) H, and He RRLs, it isclear (as has been known for a long time) that the observed CR-RLs do not arise from the hot ionised gas of the H II region but,rather, from the periphery of these regions, next to the neutralgas, namely, in PDRs. We separate the narrow line profiles ofthe CRRL transitions from those of the broader He RRLs, with which they are blended, by subtracting a Gaussian fit to He RRLfrom the observed spectra. The residual line profile representsthe relative contribution of the CRRL which is then used in ouranalysis. As mentioned in Sect. 4, toward some of the sources,the observational setup we used covers several RRLs with prin-cipal quantum numbers, n , ranging from 72 to 80 for ∆ n = ff erent HRRLs covered in oursetups for each individual source. On average, we find the linestrengths to vary by 23 % at most. Figure 11 shows that the cen-troid velocities and line widths of the CRRLs match the valuesof the HFS stacked CH line emission. This is in agreement withthe gas-phase chemistry revealed by our simple PDR models andstrongly suggests that the two species trace the same gas.It has been shown by Salgado et al. (2017), that the linewidths of high frequency CRRLs (with n < n > / or radiation broadening. Therefore, considering thee ff ect of line broadening as being only due to the random thermalmotion of particles in the gas and non-thermal e ff ects or turbu-lence, the observed line widths can be expressed as follows, ∆ υ = (cid:113) ∆ υ + ∆ υ = (cid:114) k B T kin m C + (cid:104) ∆ υ nth (cid:105) (1)where, T kin is the gas temperature, k B is the Boltzmann constant, m C is the mass of the carbon atom, and (cid:104) ∆ υ nth (cid:105) rms is the rootmean-square measure of the turbulent velocities. Furthermore,based on the premise that the CRRLs and o-CH lines trace thesame gas layers, we assume that both species are impacted bythe same turbulent flows and will hence have the same turbu-lent widths, ( ∆ υ nth ) CRRL = ( ∆ υ nth ) CH . Additionally, since the Article number, page 13 of 25 & A proofs: manuscript no. main_arxiv T m b J = 4 370679 MHz
W3 IRS5 [ K ] T m b J = 5 468375 MHz
W3IRS5 [ K ]
150 100 50 0 50
LSR T m b J = 3 269010 MHz
NSW3IRS5 [ K ] [km/s] Fig. 8.
Spectra showing the HFS transitions corresponding to the J = → J = → J = → J , F = , → , HFS line at 69017.995 MHz ismarked in light blue. The relative contribution of the NS feature is mod-elled using the nearby NS F = / → / F (cid:48) − F (cid:48)(cid:48) = J – J − observed line width of the CH emission is comparable to thatof the CRRL (see Fig. 11), when re-arranging Eq. 1 to equate thenon-thermal components, it is clear that the o-CH gas temper-ature is simply proportional to that of the CRRL, scaled by theratio of their masses.From CRRL data alone, it would be di ffi cult to accuratelydetermine the physical conditions of the C + region, for example,because of uncertainties involved in their excitation or, practi-cally speaking, simply because of the low intensities of theselines and their blending with the stronger He RRLs. These is-sues can be overcome by comparing the line intensities of CR-RLs with those of the FIR fine-structure line of ionised carbonat 158 µ m. This approach has been employed by several stud-ies to constrain the electron density, and temperature of C + lay-ers. Typically, the temperatures cover a range of values from1000 K near the H II region to ∼
100 K at the outer boundariesof this C + layer. Modelling this layer toward the W3 region,Sorochenko & Tsivilev (2000) computed a value for the kinetictemperature, T kin , of at most 200 K and an electron density of n e =
54 cm − . The PDR structure used in their analysis wasadapted from Howe et al. (1991) and does not assume a ho-mogeneous distribution of material but rather that it consists ofdense clumps ( n H ∼ cm − ) embedded in a dilute medium( n H ≤
300 cm − ) at T k ≥
100 K in order to be consistent with ob-servations. Therefore, for the specific case of W3 IRS5 we canconstrain the gas temperature to a value of at most 233 K usingEq. 1.
6. Discussion We perform non-LTE radiative transfer calculations using thestatistical equilibrium radiative transfer code RADEX (van derTak et al. 2007b), for a uniform expanding sphere geometry un-der the large velocity gradient (LVG) approximation. The codecomputes level populations, line intensities, excitation tempera-tures, and optical depths as a function of the physical conditionsspecified as input, based on the escape probability formalism.Assuming that H is the primary collisional partner of CH inthe ISM, we carry out our non-LTE analysis by using rate co-e ffi cients recently computed by Dagdigian & Lique (2018) forcollisions between CH and Helium from which we obtain o-CH –H collisional rate coe ffi cients by scaling the rates by fac-tor of 1.4, for all the fine-structure transitions among the lowest22 energy levels of o-CH . By adopting a background temper-ature of 2.73 K, with a fixed line width as estimated from theintrinsic widths of the observed Onsala spectra, we run a grid ofmodels with varying physical conditions, with the aim to con-strain the gas densities, n H , and kinetic temperatures, T kin , tovalues that are consistent with the observed o-CH emission orits upper limits. The models were computed over a temperature-density grid of size 500 × T kin values between 20–300 K(constrained by the collisional data) and n H values in the rangeof 10–10 cm − . Given that we were only able to clearly detectall three sets of o-CH HFS lines toward W3 IRS5, the non-LTEanalysis is carried out specifically for this source. Using absorp-tion spectroscopy, Polehampton et al. (2005) were able to (fromtheir column density measurements) determine a [CH] / [CH ] ra-tio through observations of both CH and CH transitions near150 µ m and 107.7 / µ m, respectively, using the ISO-LWS.These authors obtained a [CH] / [CH ] ratio of 2.7 ± ratio of 1.6, and a value closer to 3.7 for anortho-to-para CH ratio of 3 for the (systemic) +
64 km / s veloc-ity component toward Sgr B2(M), values that are consistent withresults obtained by Viti et al. (2000) – whose models addition-ally take into account grain-surface chemistry. This correspondsto o-CH abundances with respect to H between 9 . × − and1 . × − when scaled using the [CH] / [H ] ratio of 3.5 × − asdetermined by She ff er et al. (2008). By using CH column den-sities determined toward W3 IRS5 by Wiesemeyer et al. (2018)from this radical’s 150 µ m ground-state transition observed us-ing the GREAT instrument on board the Stratospheric Observa-tory For Infrared Astronomy (SOFIA) (Young et al. 2012) andthe above [CH] / [CH ] ratio, we can constrain the column den-sities of o-CH to about (4 . ± . × cm − for ortho-to-para CH ratios of between 1.6 and 3. We therefore ran modelsin the temperature-density plane for fixed values of N (o-CH )at 3 × , 5 × and 7 × cm − , the results of whichare displayed in Fig. 12. The radiative transfer analysis is fur-ther simplified by assuming a beam filling factor of unity for anextended o-CH cloud. From the results of each column den-sity model, we see that the di ff erent o-CH lines trace similartemperature and density conditions as indicated by their con-tours. We constrain these results by comparing them with thedistribution of the line intensity ratio between the 70 GHz, and68 GHz lines which are both free from contamination, in thetemperature-density plane. The physical conditions that prevailin these regions are constrained based on the behaviour of χ across this parameter space. The χ value is computed across Article number, page 14 of 25. Jacob et al.: Hunting the elusive Methylene radical
100 50 0 50
LSR T m b J = 4 370679 MHz
W49N [ K ] [kms ]
100 50 0 50 100
LSR T m b J = 4 370679 MHz
W75 N [kms ] [ K ]
50 0 50 100 150
LSR T m b J = 4 370679 MHz
W43 [ K ] [kms ]
100 50 0 50 100
LSR T m b [ K ] J = 4 370679 MHz
DR21 [kms ]
100 50 0 50
LSR T m b J = 4 370679 MHz
S140 [kms ] [ K ] Fig. 9.
Spectra of the o-CH J = → A V N o r m a li s e d a b un d a n c e H H C + CH +3 CCH CH CO (n = 10 cm , G = 1) A V (n = 10 cm , G = 1 × 10 ) H H C + CH +3 C CH CH CO
H H (n = 10 cm , G = 1) C + CH +3 CCH CH CO (n = 10 cm , G = 1 × 10 ) H H C + CH +3 C CH CH CO
Fig. 10.
Variation of the gas-phase abundancesof fundamental species in carbon chemistry,normalised by their respective peak abundancesas a function of visual extinction, A v . The cloudparameters for the di ff erent models are as fol-lows: n H = cm − , G = n H = cm − , G = × , n H = cm − , G = , and n H = cm − , G = × (bottom-right). All the models presented hereuse a constant cosmic-ray ionisation rate of ζ H = . × − s − . The pink and grey shadedregions in each panel highlight the transitionlayer between H-H and C + -C, respectively. the entire grid as follows, χ = (cid:88) (cid:0) T r,obs − T r,mod (cid:1) /σ T r,obs , (2) χ = χ / ( n − , (3)where T r,obs and T r,mod represent the observed and modelled linebrightness temperatures on the T mb scale, σ T r,obs represents un- certainties in the observed o-CH spectra and n is the num-ber of degrees of freedom. The χ values were fit at a 99.9%probability for two degrees of freedom. For each column den-sity model, we find that the χ values show more variationswith T kin than with n H . In Table 3, we summarise the estimatedrange of gas temperatures and densities for the di ff erent mod- Article number, page 15 of 25 & A proofs: manuscript no. main_arxiv
50 0 50 100 150 200 250
LSR T m b W3IRS5 [ K ] H(75)He(75)C(75) [kms ]
80 60 40 20 0
LSR T m b W3IRS5 [ K ] [kms ] stacked68GHz CH stackedC(75)CI / 125 Fig. 11.
Left: Observed H α , He α , and C α transitions for n =
75 toward W3 IRS5. The LSR velocity scale is given with respect to the C α line at15.82 GHz. The pink dashed curve represents the fit to the He RRL and the shaded purple region confines the CRRL. Right: Decomposed CRRLprofile in blue alongside the 68 GHz and 70 GHz HFS-stacked CH line profiles displayed by the dashed orange and dashed-dotted violet curves,respectively. The stacked profile of the 70 GHz CH transition was obtained from the HFS decomposition model. The hatched grey regions displaythe line profile of the P − P transition of C I at 492.160 GHz scaled down by a factor of 125 on the T mb scale. els. Across the di ff erent models we derive T kin between (150–200) K and n H ∼ × cm − consistent with not only previouschemical models (Lee et al. 1996) but also our hypothesis thato-CH must arise in a hot but dilute media. The derived valuesof temperature are also consistent with that determined from thecomparison with CRRLs, as discussed in Sect. 5.6. Furthermore,with critical densities of the order of 10 cm − , the derived den-sities reveal that the observed 68–71 GHz o-CH emission linesarise from sub-thermally excited gas.We expect the non-LTE analysis toward the W51 pointingpositions to reproduce physical conditions similar to those de-rived toward W3 IRS5 since the line intensities of both the70 GHz and 68 GHz components of CH , and their correspond-ing line ratios are comparable toward both sources. The resultsof this analysis confirm our suspicions and clearly addresses whyCH is not widely detected in the Orion Bar PDR. From the non-LTE analysis presented here, we would expect the abundance ofCH to peak in regions with H gas densities of ∼ × cm − and the Orion Bar PDR, as discussed earlier, simply traces gasof higher densities ( > cm − ). The dense Orion Bar is a spa-tially limited region in the much larger, on average, more dif-fuse PDR associated with the Orion Nebula, for which a den-sity of 10 cm − is a representative value (Cuadrado et al. 2015;Nagy et al. 2015a). Because of the Orion Nebula’s proximity( ≈
400 pc) this larger scale PDR is well resolved, even in our arcminute size beams. However, the other regions toward which wedetect CH emission are much further away and our observationssample the total extent of their large, low-density PDRs. This hasalready been discussed at the end of Sect. 5.5 and here findssupport by our non-LTE calculations. Furthermore, the nomi-nal hydrogen nucleus density we derive is a factor of ∼ − estimated by Welty et al.(2020), who were only able to obtain upper limit for the CH transition near 1397 Å in the translucent cloud along the LOStoward HD 62542.As for the excitation of these lines, similar to the results ob-tained by Dagdigian & Lique (2018), we find that over the verywide range of gas densities we modelled ( n H = cm − ),our calculations produce negative opacities ( ∼ − − ) or line in-version at an excitation temperature of − .
36 K. We thus find allthree of the fine-structure lines of the 4 − transition of o-CH seen in emission to be weak masers. From our RADEX models, Table 3.
Summary of RADEX results.
Model N (o-CH ) log ( n H / cm − ) T kin [cm − ] [K]I 3 × . + . − . + − II 5 × . + . − . + − II 7 × . + . − . + − we compute a ∼
1% inversion in the population of the 4 − fine structure transitions, where the percentage of population in-version for a two-level system is given by ( n u − n l ) / ( n u + n l ) and n u , and n l are the upper and lower energy level populations. Sincethe models produce masing conditions even in the absence ofstrong, external radiation fields and without line overlap consid-erations, it is clear that the observed masing e ff ect in these high-lying lines is a robust phenomenon which preferentially popu-lates the N K a K c = level over that of the N K a K c = level. Theobserved emission spectra may have contributions from the ex-tended continuum background radiation in these regions as wellas collisional pumping e ff ects. While the weak level-inversionobserved in these lines explains why they are detectable in thefirst place, the degree of population inversion itself is greatly de-pendent on the collisional rate coe ffi cients.We exclusively detect CH within PDR cloud layers associ-ated with H II regions and not in those associated with (P)PNeand SNRs, a few of which were also observed by us. The lackof detectable amounts of CH in these types of objects can beattributed to the fact that CH arises from dilute PDR layers, asimplied by our non-LTE radiative transfer analysis. The PDRssurrounding (P)PNe are dense regions ( ∼ − cm − ) asthey form in the compressed inner layers of the remnant circum-stellar envelopes of AGB stars (see for example Cox et al. 2002).This argument explaining the non-detection of CH in (P)PNeowing to the elevated densities of their PDRs also holds true forthe case of the SNRs for which upper limits are presented inthis study. Our Onsala 20 m beam was pointed toward a densemolecular clump in IC 443 residing in the interaction zone ofthe SNR with a molecular cloud that has gas densities as high as10 cm − (Dickman et al. 1992), while for the dense molecularknots detected in high- J CO lines toward Cas A, gas densi-
Article number, page 16 of 25. Jacob et al.: Hunting the elusive Methylene radical
100 200 300 T kin [K] l o g ( n H / c m ) N(CH ) = 3 × 10 cm
100 200 300 T kin [K] N(CH ) = 5 × 10 cm
100 200 300 T kin [K] N(CH ) = 7 × 10 cm RatioT mb
70 GHzT mb
68 GHzT mb
69 GHz
Fig. 12.
RADEX non-LTE modelling of o-CH toward W3 IRS5. The solid and dashed- red, blue and orange curves represent the observed peakmain-beam brightness temperatures of the 70 GHz, 68 GHz and 69 GHz, respectively, and their uncertainties. We note that for the 69 GHz o-CH lines being contaminated by NS, we only represent the uncontaminated modelled intensity and an upper limit. The black dashed and dotted linesmark the limits of the brightness temperature ratio between the 70 and 68 GHz lines. The grey shaded region characterises χ values of < × cm − (left),5 × cm − (centre), and 7 × cm − (right). ties of 10 − cm − have been determined (Wallström et al. 2013).Using the physical conditions derived from these models andthe rate coe ffi cients calculated between p -CH –H , which arealso based on calculations by Dagdigian & Lique (2018), we pre-dict the brightness temperatures and excitation temperature forthe 444 GHz p-CH transitions. The model was run for a rangeof gas densities n (H ) = − cm − with a fixed CH col-umn density and gas temperature of N (p-CH ) = × cm − (in the limit where N (p-CH ) = N (o-CH )) and T kin =
163 K,respectively. The modelled results are displayed in Fig. 13. Upuntil the critical density of 2 × cm − is reached, the bright-ness temperatures reproduced by the models are low and evenslightly negative, particularly for the range of densities derivedfrom the 70 GHz o-CH lines. The excitation temperature atthese frequencies is also small and close to the background ra-diation temperature, which, for these models, is governed by thecosmic microwave background at 2.73 K. With excitation tem-peratures below 10 K up to densities of 10 cm − , the modelssuggest that it is highly unlikely to observe detectable amountsof p-CH at 444 GHz and in absorption, if any at all. We further compare the observed CH and CRRLs withthat of the neutral atomic carbon, C I , P − P transition at492.160 GHz for those sources toward which this transition isavailable, using ancillary data from the Herschel / HIFI archive .These observations, published in Gerin et al. (2015), were pro-cured using the HIFI band 1a which provides a FWHM beamwidth of 43.1 (cid:48)(cid:48) at these frequencies, which is comparable to thebeam size of our CH and CRRL observations. We find that theline ratio between the C I , and CH
70 GHz emission to be almosta constant, with a value of 125:1 toward W3 IRS5, W51 M, andW49 N, (see Fig. 11 and top panel of Fig. C.4). Furthermore, wesee that toward W51 M and W49 N the C I emission shows twocomponents along the LOS and, perhaps, the second componentmay also be present in CH but has gone undetected because our See, http://archives.esac.esa.int/hsa/whsa/ log( n H /cm ) T m b [ K ] ×10 p CH log( n H /cm ) T e x [ K ] p CH Fig. 13.
Main-beam brightness temperature ( T mb ) and excitation tem-perature ( T ex ) for the 2 − fine-structure lines of p-CH . The insetpanels expand on the ( T mb ) and ( T ex ) values for gas densities between10–10 cm − . observations do not attain the sensitivity necessary to detect thisweaker component. Article number, page 17 of 25 & A proofs: manuscript no. main_arxiv A v [ C I ] / [ C H ] G o = 10G o = 10 G o = 10 G o = 10 G o = 10 Fig. 14.
Variation in the [C I ] / [CH ] abundance ratio as a function of A v . The di ff erent curves represent models at di ff erent values of G o , aslabelled. The dashed blue curve marks the observed [C I ] / [CH ] intensityratio of 125. The similarity in the line intensity ratio between the emis-sion of the C I and the 70 GHz CH component toward thesethree sources strongly suggests that CH is likely to be formedunder similar physical conditions in these regions, in spite oftheir di ff erent levels of star-formation. Using the constraints onthe physical conditions derived in Sect. 6.1, we revise the chemi-cal models presented in Sect. 5.6. We once again ran the PyPDRcode, this time for the gas density and temperature values derivedfrom our non-LTE analysis. In order to reproduce the observed[C I ] / [CH ] line intensity ratio of 125, we carried out the analy-sis for di ff erent models with varied values of G o at a fixed dusttemperature of 50 K. None of the models were capable of repro-ducing the observed [C I ] / [CH ] line intensity ratio (see Fig. 14).The model that most closely matched the observations gives us a[C I ] / [CH ] ratio of 268, and was that with a G o = in Habingunits for values of A v between 0 and 1. However, overall themodelled line ratios are almost always overestimated; or rather,the CH abundance is underestimated. This could either be be-cause the underlying chemical network for CH used in thesemodels is incomplete or because the PDR model itself is toosimplistic or a combination of both. Perhaps a more robust PDRmodel is required, namely, one which includes dust illuminationby both an internal as well as an external heating source withinthe PDR structure.
7. Summary and conclusions
In this work, we present observations of the 4 − transi-tions of o-CH made with both the KP 12 m and Onsala 20 mtelescopes. The former observations were prompted by the non-detection of the 68–71 GHz CH lines toward the Orion-KL hotcore using the IRAM 30 m telescope and the GBT. Our KP ob-servations were therefore used, firstly, to verify the original de-tection of CH toward the Orion-KL, and secondly, to addressif the CH emission is extended by investigating nearby posi-tions in OMC-1 and its related H II region / PDR. These observa-tions confirm the detection of CH toward Orion-KL and revealhigher intensities of CH toward PDR positions. The interpre-tation of these results clearly indicates that the CH emission isextended but does not constrain their source of origin. We clarifythe nature of the CH emission using the smaller angular beamsize of the Onsala 20 m telescope. These observations show thatthe distribution of CH toward the Orion KL-S and Orion Bar regions coincides with peaks of the [C II ] emission. This is con-sistent with the results of simple PDR models which show thatthe CH abundance peaks in gas layers where ionised carbon re-combines to form neutral atomic carbon. As for the observations,we find that the line profiles of the CH transitions are akin tothose of CRRLs, which are well-known tracers of PDRs. Basedon CH ’s association with PDRs, we expanded our target sam-ple to well known H II regions, which naturally have associatedPDRs, and successfully detected the 70 GHz fine-structure lineand its HFS components for the first time toward W51 E, W51 N,W49 N, W43, W75 N, DR21, and S140 as well as all three of the68–71 GHz fine- and HFS components toward W3 IRS5. Stack-ing each of the HFS components of the 68 GHz transition in apreviously averaged spectrum of the di ff erent W51 positions re-vealed a 3.5 σ detection of the 68 GHz CH line. Carrying out anon-LTE analysis, we constrained the physical conditions of theCH clouds to gas densities and temperatures of ∼ × cm − and 150–200 K, respectively. The physical conditions probed areconsistent with those that prevail in translucent clouds and arein agreement with the results of previous theoretical models byLee et al. (1996), who show the fractional abundances of CH topeak at intermediate densities. While our non-LTE analysis wasfocused on W3 IRS5, the HFS-stacked and averaged pointingpositions toward W51 show similar line intensities and thereforeprobe similar gas conditions. It is only with further, highly sen-sitive observations of the 68 GHz component of CH can weverify if the physical properties of the gas probed by CH areubiquitous across sources. Our analysis of the excitation condi-tions reveal that these lines are weakly masing with low negativeexcitation temperatures. The weak level inversion amplifies theemission of these lines which may otherwise have remained un-detectable. We also report the non-detection of the 2 –3 tran-sitions of p-CH toward Orion S at a 3 σ upper limit of 24 mKand at 0.23 K toward the other positions in Orion.Our results establish the 4 − transition of CH as tracingthe hot but low density component of the ISM present in PDRlayers, with its abundance peaking at the edges of dense clouds.As an essential intermediary in interstellar carbon chemistry, it isimportant to carry out further observations of this molecule andparticularly of its energetically lower transitions. We hope thatthe results of our study will help guide future searches for thisradical in other sources at FIR and sub-millimetre wavelengths.
8. Open questions and outlook
The present work provides strong observational evidence forCH in PDRs of relatively low density (several times 10 cm − ).This is supported by our radiative transfer modelling, whichsuggests temperatures >
150 K for the CH -bearing material.While such densities are also found in translucent interstellarclouds, the temperatures in these clouds are thought to be muchlower, 10–50(?) K (Snow & McCall 2006). Evidence for CH intranslucent, and maybe even di ff use interstellar clouds can alsobe found in the spectra of the FIR o- and p-CH ground-statetransitions toward Sgr B2. In these, Polehampton et al. (2005)find, absorption not only at Sgr B2’s systemic velocity but alsofrom clouds along the LOS to the Galactic centre region. Al-though these ISO / LWS data have a low spectral resolution, thefitted velocity components closely correspond to those of numer-ous mm-wavelength absorption lines that arise from such inter-vening clouds, whose physical conditions are still poorly con-strained, but which are thought to consist of translucent or dif-fuse interstellar material (Thiel et al. 2019); however, for statisti-cal reasons they are unlikely to be PDRs. This raises the question
Article number, page 18 of 25. Jacob et al.: Hunting the elusive Methylene radical whether CH can attain substantial abundance not only in PDRs,but in translucent (or even di ff use) clouds as well. The availableevidence from electronic transitions discussed in Sect. 1 is sofar inconclusive, but the procurement and analysis of further UVabsorption data would be most interesting.As for extra-galactic observations, Muller et al. (2011) do notdetect o-CH in the line survey they carried out with the Aus-tralian Telescope Compact Array of the 7 mm window, corre-sponding to rest frequencies of 57–94 GHz along two sight linesthrough the red-shift z = . cm − and 80 K, respectively, conditions that are neither di-lute enough nor hot enough to excite the CH lines in question.Alternatively, their survey is still too shallow to detect the CH lines.It is di ffi cult to meaningfully constrain the CH abundancejust with HFS lines assigned to a single rotational level ( J =
5) –while emission from higher frequency rotationally excited sub-millimetre CH transitions will remain undetectable because oftheir high critical densities. To compound this di ffi culty, the onlyaccessible lines that are likely to be excited; the 69–71 GHz o-CH lines, extensively discussed here, are ubiquitously weak andboosted into detectability by inversion, which makes their inter-pretation not quite so straightforward. Moreover, the validity ofthe conclusions drawn from these lines is heavily dependent onthe quality of the collisional rate coe ffi cients used in the radia-tive transfer modelling, which are currently only available forthe CH –He system (and were scaled by us to apply to colli-sions with H ). Recently, van der Tak et al. (2020) have dis-cussed the current status and future plans for the Leiden Atomicand Molecular Database (LAMDA). In this paper, they review,among others, the use of collisional rate coe ffi cients for colli-sions of molecules with He and the usual practice of scaling toestimate H rate coe ffi cients. For the case of CH they assign alow accuracy (factor of ∼ , amongst other molecules, has ‘to be the object of new scat-tering studies considering H as a projectile’. Calculating CH –H rate coe ffi cients would certainly be a significant computa-tional project. Given that in a practical sense, the complete inter-pretation of CH emission in the ISM critically depends on thesecollisional rate coe ffi cients, it is one that should be considered.Apart from the 68–71 GHz lines, we are only left with theUV and the FIR resonance lines. These also present a bleak pic-ture as the former still await a convincing identification in theISM. With regard to the FIR lines, unfortunately, neither the o-and nor the p-CH ground state lines at 127.6 and 107.7 µ m,respectively, were covered by the HIFI instrument on Herschel.Moreover, none of the modules of the GREAT instrument onSOFIA covers their wavelengths at present. Even if GREAT wasable to do this (and it may do so in the future), detecting theselines might be di ffi cult given their low opacities observed by ISO(with the low LWS resolution) of just 0.02 to 0.05 toward one ofstrongest FIR sources in the sky, Sgr B2 (whereas for the 149 µ mCH ground state, lines opacities of ∼ . in an astrochemical context remains elusive. Acknowledgements.
The authors acknowledge support from Onsala Space Ob-servatory for the provisioning of its facilities / observational support. The OnsalaSpace Observatory national research infrastructure is funded through SwedishResearch Council grant No 2017-00648. We would like to thank the KP 12 mtelescope operators for their help during the on-site and remote observations,and Abby Hedden for her invaluable assistance with the observations as well as Arnaud Belloche and Ed Polehampton. We thank Paul Dagdigian for a discus-sion on the collision rates. We would also like to thank the anonymous refereefor a careful review of the article and valuable input. The authors would like toexpress their gratitude to the developers of the many C ++ and Python libraries,made available as open-source software, in particular this research has madeuse of the NumPy (Harris et al. 2020), SciPy (Jones et al. 2001) and matplotlib(Hunter 2007) packages. References
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Table A.1.
Onsala 20 m telescope non-detections.Source α J2000 δ J2000 υ LSR
Line rms * [kms − ] [GHz] [mK]Orion 05:35:16.96 -05:22:02.7 + / S (1) 69 33Orion 05:35:14.28 -05:22:27.5 + / S (2) 68 5369 48Orion 05:35:14.28 -05:23:16.5 + / S (3) 68 5569 43Orion 05:35:16.46 -05:23:22.7 + / S (4) 68 6169 37Orion 05:35:13.10 -05:23:56.0 + / S (5) 68 4069 30Orion 05:35:24.96 -05:22:32.7 + / S (6) 69 25Orion 05:35:25.30 -05:24:34.0 + + + + + + + + + + + Appendix A: CH non-detections In this appendix we list the rms noise levels for the sources to-ward which we do not detect any o-CH . Appendix B: CH
70 GHz HFS stacking
As mentioned in Sec. 5.5 and 5.6, we stack the HFS-decomposedmodel fits of each of the individual HFS components of the70 GHz transition of CH , prior to comparing its line profilewith the observed line profiles of the CRRLs. Furthermore, inorder to carry out a more objective comparison we incorporatedthe expected noise of the (stacked and averaged) 70 GHz spectraby including an additive Gaussian noise term. As an example,the di ff erent steps that are a part of this exercise are illustratedfor the 70 GHz spectrum observed toward W3 IRS5 in Fig. B.1. Table A.1.
Continued.Source α J2000 δ J2000 υ LSR
Line rms * [kms − ] [GHz] [mK]W3(OH) 02:27:04.10 61:52:22.00 -46.0 70 768 1369 8S233 05:35:51.19 35:44:12.90 -17.5 70 1068 1769 15S235 05:40:53.30 35:41:49.00 -57.0 70 1968 3769 30CRL618 04:42:53.62 36:06:53.30 -26.0 70 6068 10069 80CRL2688 21:02:18.27 36:41:37.00 -32.0 70 1268 2069 18NGC7023 21:01:36.90 68:09:48.00 2.5 70 1268 2469 22NGC7027 21:07:01.75 42:14:10.00 8.9 70 668 1169 9NGC7538 23:13:37.20 61:30:00.00 -64.9 70 1968 3869 28Cas A 23:23:24.90 58:50:03.30 -2600.0 70 868 1569 12IC 443 06:18:02.70 22:39:03.00 -4.4 70 1068 1669 14 Notes.
Columns are: (left to right) source name, J2000 coordinates ofobserved position, assumed LSR velocity, line ID, and 1 σ rms noiselevel. ( * ) On the T mb scale the rms noise is quoted for a spectral reso-lution of 0.97 km s − . All the sources are SFRs, except for NGC7027which is a planetary nebula, CRL618 and CRL2688, which are proto-planetary nebulae and Cas A and IC 443, which are supernova remnants.NGC7023 is a reflection nebula with a PDR. Appendix C: Observed recombination lines
In this appendix we present the observed spectra of the di ff er-ent RRLs. As mentioned in Sect. 4, there are two observing se-tups used for these observations. The first setting observed onlythe H-, He-, and CRRLs for a principal quantum number of 75,near 15 GHz. This setup was used to carry out observations to-ward the di ff erent Orion positions and W3 IRS5 (displayed inFig. 11). The second setup was used for the remainder of thesources toward which o-CH was observed, utilising a 5 GHzbandwidth from 12.5-17.5 GHz, which covered RRLs with prin-cipal quantum numbers between 80 and 72. Subsequently, forthese sources, the RRL profiles for the di ff erent quantum num-bers are stacked and averaged; based on these profiles, the con-tributions of the He RRL (modelled using a Gaussian fit) is sub-tracted to yield the CRRL profile used in the analysis. Article number, page 21 of 25 & A proofs: manuscript no. main_arxiv T m b [ K ] W3IRS5
Observed spectrumhfs fitsCombined fit T m b [ K ] NoiseShifted hfs fitshfs stack
80 60 40 20 0
LSR T m b [ K ] [kms ] Stacked spectrum+ noise
Fig. B.1.
Top: Observed o-CH
70 GHz spectrum towards W3 IRS5(black) alongside the HFS decomposed fits (dotted blue) and the com-bined fit (red). Middle: The HFS stacked profile (red) with the fit to eachHFS component shifted to the systemic velocity of the source (dottedblue) and the Gaussian noise profile (black). Bottom: Resultant stackedspectrum with the addition of noise (black).
Table C.1.
Best-fit Gaussian parameters for the CRRLs.
Source υ ∆ υ T mb [km s − ] [km s − ] [mK]Orion KL / S (2) 8.8(0.1) 3.7(0.2) 140.0(5.4)Orion KL / S (3) 9.3(0.1) 4.8(0.3) 74.4(3.3)Orion KL / S (5) 8.1(0.1) 4.6(0.3) 64.7(3.3)Orion Bar (1) 10.1(0.1) 3.7(0.2) 239.4(8.5)Orion Bar (3) 10.3(0.2) 4.6(0.5) 118.8(9.1)Orion Bar (4) 10.6(0.1) 2.6(0.1) 304.6(11.1)W3 IRS5 -38.4(0.1) 6.8(0.4) 131.4(6.6)W51 M 59.8(1.0) 6.1(1.8) 69.0(3.7)W49N 5.6(0.6) 6.7(1.5) 79.0(3.7)DR21 -8.2(2.6) 7.7(1.5) 60.5(6.1)W43 92.4(0.6) 2.4(1.0) 18.6(3.6)
Article number, page 22 of 25. Jacob et al.: Hunting the elusive Methylene radical
50 0 50 100 150 200 250
LSR [kms ] T m b [ K ] Orion KL/S(2)
H(75)He(75)C(75)
50 0 50 100 150 200 250
LSR [kms ] T m b [ K ] Orion KL/S(3)
H(75)He(75)C(75)
50 0 50 100 150 200 250
LSR [kms ] T m b [ K ] Orion KL/S(5)
H(75)He(75)C(75)
10 0 10 20 30
LSR [kms ] T m b [ K ] Orion KL/S(2)
C(75)
10 0 10 20 30
LSR [kms ] T m b [ K ] Orion KL/S(3)
C(75)
10 0 10 20 30
LSR [kms ] T m b [ K ] Orion KL/S(5)
C(75)
Fig. C.1.
Top: H-, He-, and CRRLs toward di ff erent Orion KL / S OSO pointing positions, with the Gaussian fit to the He-RRL displayed by thedashed magenta curve. Bottom: Resulting CRRL with the Gaussian fit to the profile displayed in magenta. The velocity scale is given with respectto the CRRL.
50 0 50 100 150 200 250
LSR [kms ] T m b [ K ] Orion Bar (1)
H(75)He(75)C(75)
50 0 50 100 150 200 250
LSR [kms ] T m b [ K ] Orion Bar (3)
H(75)He(75)C(75)
50 0 50 100 150 200 250
LSR [kms ] T m b [ K ] Orion Bar (4)
H(75)He(75)C(75)
10 0 10 20 30
LSR [kms ] T m b [ K ] Orion Bar (1)
C(75)
10 0 10 20 30
LSR [kms ] T m b [ K ] Orion Bar (3)
C(75)
10 0 10 20 30
LSR [kms ] T m b [ K ] Orion Bar (4)
C(75)
Fig. C.2.
Same as Fig. C.1, but toward the Orion Bar OSO pointing positions. Article number, page 23 of 25 & A proofs: manuscript no. main_arxiv
LSR T m b W51M
HHeCn = 80n = 72 [kms ] [ K ]
25 50 75 100 125
LSR T m b W51M
HeC [ K ] [kms ] n = 80n = 72 0.00.51.0 T m b W51M
HeC [ K ]
25 50 75 100 125
LSR T m b [ K ] [kms ] Resultant C
LSR T m b W49N
HHeCn = 80n = 72 [kms ] [ K ]
25 0 25 50 75
LSR T m b W49N
HeCn = 80n = 72 [kms ] [ K ] T m b W49N
HeC [ K ]
25 0 25 50 75
LSR T m b [ K ] [kms ] Resultant C50 0 50 100 150 200
LSR T m b DR21
HHeCn = 80n = 72 [ K ] [kms ]
50 25 0 25 50
LSR T m b DR21
HeC [ K ] [kms ] n = 80n = 72 0.00.20.4 T m b DR21 [ K ] HeC50 25 0 25 50
LSR T m b [ K ] [kms ] Resultant C50 100 150 200 250
LSR T m b W43
HHeCn = 79n = 72 [ K ] [kms ]
50 75 100 125 150
LSR T m b W43
HeCn = 79n = 72 [ K ] [kms ] T m b W43
HeC [ K ]
75 100 125 150
LSR T m b [ K ] [kms ] Resultant C
Fig. C.3.
Left: H-, He-, and CRRLs observed with the E ff elsberg 100 m telescope for α lines with principal quantum number, n , from 72 (top)to 80 (bottom) (except for W43) in steps of ∆ n =
1. The positions of the di ff erent lines are marked by dashed dark blue lines. The spectra areuniformly o ff set by 0.5 K on the T mb . Centre: Zoomed in view of the He-, and CRRL. The o ff sets are not uniform but varied for the ease of viewingwith a typical o ff set of 0.2 K on the T mb scale. Right: Average He-, and CRRL profile (top) with the Gaussian fit to the He RRL presented by thedashed dark blue curve. The resulting spectrum of the CRRL is displayed below with its fit given by the dashed magenta Gaussian. For all spectra,the velocity scale is appropriate for the CRRL. The di ff erent rows (from top to bottom) present these results for W51 M, W49 N, DR21, and W43,respectively.Article number, page 24 of 25. Jacob et al.: Hunting the elusive Methylene radical
20 40 60 80 100
LSR T m b W51M [ K ] [kms ] stackedCCI / 125
20 0 20 40
LSR T m b W49N [ K ] [kms ] stackedCCI / 125
40 20 0 20
LSR T m b DR21 [ K ] [kms ] stackedC
60 80 100 120
LSR T m b W43 [ K ] [kms ] stackedC Fig. C.4.
Decomposed CRRL profile in blue alongside the 70 GHz HFS stacked CH line profile displayed by the dashed-dotted violet curvetoward (clock-wise from top) W51 M, W49 N, DR21, and W43, respectively. The stacked profile of the 70 GHz CH transition was obtained fromthe HFS decomposition model. The hatched grey regions in the W51 M and W49 N spectra display the line profiles of the P − P transition ofC I at 492.160 GHz scaled down by a factor of 125 on the T mbmb