Massive 70 micron quiet clumps I: evidence of embedded low/intermediate-mass star formation activity
A. Traficante, G.A. Fuller, N. Billot, A. Duarte-Cabral, M. Merello, S. Molinari, N. Peretto, E. Schisano
aa r X i v : . [ a s t r o - ph . GA ] J un Mon. Not. R. Astron. Soc. , 000–000 (0000) Printed 8 October 2018 (MN L A TEX style file v2.2)
Massive 70 µ m quiet clumps I: evidence of embeddedlow / intermediate-mass star formation activity A. Traficante , ⋆ , G.A. Fuller , N. Billot , A. Duarte-Cabral , M. Merello ,S. Molinari , N. Peretto and E. Schisano Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, University of Manchester, Oxford Road, Manchester M13 9PL, UK IAPS - INAF, via Fosso del Cavaliere, 100, I-00133 Roma, Italy Observatoire astronomique de l’universit´e de Geneve, Chemin des Maillettes, 51, CH-1290 Versoix, Suisse School of Physics and Astronomy, Cardi ff University, Queens Buildings, The Parade, Cardi ff CF24 3AA, UK
ABSTRACT
Massive clumps, prior to the formation of any visible protostars, are the best candidates tosearch for the elusive massive starless cores. In this work we investigate the dust and gasproperties of massive clumps selected to be 70 µ m quiet, therefore good starless candidates.Our sample of 18 clumps has masses 300 . M . ⊙ , radius 0 . R .
00 pc, surfacedensities Σ > .
05 g cm − and luminosity / mass ratio L / M .
3. We show that half of these70 µ m quiet clumps embed faint 24 µ m sources. Comparison with GLIMPSE counterpartsshows that 5 clumps embed young stars of intermediate stellar mass up to ≃ . ⊙ . Westudy the clump dynamics with observations of N H + (1 − −
0) and HCO + (1 − . . ˙M . . × − M ⊙ yr − , comparable with values found in high-mass protostellar regions, and free-fall time ofthe order of t f f ≃ × yr. The only appreciable di ff erence we find between objects withand without embedded 24 µ m sources is that the infall rate appears to increase from 24 µ mdark to 24 µ m bright objects. We conclude that all 70 µ m quiet objects have similar propertieson clump scales, independently of the presence of an embedded protostar. Based on our datawe speculate that the majority, if not all of these clumps may already embed faint, low-massprotostellar cores. If these clumps are to form massive stars, this must occur after the formationof these lower mass stars. Key words:
Stars – stars: formation – stars: kinematics and dynamics – stars: massive – Re-solved and unresolved sources as a function of wavelength – radio lines: stars – submillimetre:stars – Physical Data and Processes – line: profiles
Massive stars play a crucial role in the formation and gas en-richment of the hosting Galaxy, and yet the formation mech-anism of these extreme objects is unclear (Beuther et al. 2007;Zinnecker & Yorke 2007; Tan et al. 2014). The massive star for-mation begins in molecular clouds with su ffi cient density to formmassive objects (Tan et al. 2014). Some of these regions havebeen detected in absorption against the strong 8 and 24 µ mbackground, the so-called IRDCs (Perault et al. 1996; Carey et al.1998). Massive stars form in the densest part of the natal molec-ular cloud, within condensations that are called clumps (Blitz1993; Zinnecker & Yorke 2007; Tan et al. 2014), objects with size ≃ . − ⋆ e-mail:alessio.trafi[email protected] Σ = .
05 g cm − (Urquhart et al. 2014), and with no signatures ofon-going star formation activity are ideal massive starless clumpcandidates.These extremely young clumps may be the precursors of mas-sive starless cores, an initial condition required in core accretionmodels of star formation (Tan et al. 2014). It is however not welldetermined if massive clumps embed massive starless cores, or ifthey fragment into a number of low-mass cores. In a sample of9 high-mass infared-quiet cores in Cygnus X, Duarte-Cabral et al.(2013) found that 8 out of 9 of these cores are driving outflows,therefore must be protostellar. The remaining one has only a tenta-tive outflow detection, and could potentially be in a prestellar phase.Starless clump candidates are hard to find, in particular fortheir short lifetime, of the order t ≃ yr (Motte et al. 2007). Theidentification of a statistically significant number of these candi-dates in the Galaxy requires unbiased surveys of the Galactic Planeat wavelengths which allows us to trace the cold dust envelopes c (cid:13) A. Traficante et al. of these star forming regions, which emit principally in the far-infrared (FIR) / sub-mm.The ATLASGAL survey (Schuller et al. 2009) observed awide portion of the I and IV quadrant of the Galactic plane at870 µ m and produced a survey of starless clumps in the region10 ◦ l ◦ , | b | ◦ (Tackenberg et al. 2012). This survey iden-tified 210 starless clumps, but only 14 which may form stars moremassive than 20 M ⊙ . The search for young massive cluster (YMC)precursors in the range 20 ◦ > l > ◦ , combining ATLASGALdata with methanol emission, found only 7 potential YMC candi-dates (Urquhart et al. 2013). In the characterization of the proper-ties of cluster progenitors combining the MALT90 (Jackson et al.2013) and the ATLASGAL surveys, Contreras et al. (2017) identi-fied 24 over 1244 sources as potential starless candidates and only1 clump with properties consistent with a YMC precursor. A recentsearch for young massive cluster progenitors in the Galactic cen-ter using ATLASGAL and the H O southern Galactic plane sur-vey (HOPS) found that all YMC candidates are already formingstars (Longmore et al. 2017). These results are in agreement withthe finding of Ginsburg et al. (2012) using the Bolocam GalacticPlane survey (BGPS, Aguirre et al. 2011). These authors searchedfor massive clumps in the first quadrant and found that none wase ff ectively starless. More recently Svoboda et al. (2016) identifiedover 2000 starless clump candidates in the entire BGPS survey anda lack of candidates with masses in excess of 10 M ⊙ .A major contribution to the field comes from the Herschel sur-vey of the Galactic Plane, Hi-GAL (Molinari et al. 2010). Hi-GALobserved the entire Plane in the wavelength range 70 λ µ mallowing a direct estimation of temperature, mass and luminosity ofthe clumps with known distances. In particular, since the presenceof a 70 µ m source is interpreted as a signpost of protostellar activ-ity (Dunham et al. 2008), the Hi-GAL survey allows a characteri-zation of hundreds of 70 µ m quiet clumps, i.e. starless clump can-didates (Veneziani et al. 2013; Elia et al. 2013; Elia & et al. 2017).A first search of starless clumps with Hi-GAL was carried out byVeneziani et al. (2013). These authors analyzed Hi-GAL scienceverification data taken in two 2 ◦ × ◦ wide regions centered on l = ◦ and l = ◦ and found hundreds of starless clump candi-dates. In a ≃ ◦ wide region of the outer Galaxy Elia et al. (2013)identified 688 starless clumps, the majority of them gravitationallybound sources.Combining the Hi-GAL data with the comprehensive cat-alogue of IRDCs of Peretto & Fuller (2009), Traficante et al.(2015a) have performed a survey of starless and protostellar clumpsassociated with IRDCs with known distances ( ≃ ◦ l ◦ . These authors found 667 starlessclump candidates with masses up to 10 M ⊙ , ≃
240 of which withsurface density Σ > .
05 g cm − , so potentially forming massivestars.In this paper we present a detailed study of a sample of70 µ m quiet clumps mostly extracted from the Traficante et al.(2015a) catalogue for which we made follow-up observations inthe dense molecular tracers N H + (1 − −
0) and HCO + (1 −
0) with IRAM 30m telescope . A second paper in this series(Traficante et al. 2017, submitted; hereafter, Paper II) is dedicatedto the study of the properties of the non-thermal motions of theseclumps.This paper is divided as follow: in Section 2 we present the IRAM is supported by INSU / CNRS (France), MPG (Germany) and IGN(Spain). observations of the dust continuum and the line emission; in Sec-tion 3 we describe the photometry steps that we follow to obtainthe fluxes of these clumps combining the Hi-GAL, ATLASGALand BGPS datasets. In Section 4 we analyze the spectral energydistribution (SED) of the clumps and derive the main properties oftheir dust emission. In the same Section we also analyze the Mid-infrared (MIR) counterparts to identify clumps with faint 24 µ memission and GLIMPSE counterparts. In Section 5 we analyze thespectra of the dense gas tracers and we derive gas column densitiesand abundances. In Section 6 we identify clumps with evidence ofinfalling motions and explore the relations between dust and densegas tracer properties, comparing the properties of the 24 µ m darkand 24 µ m bright sources; In Section 7 we summarize our results. A sample of 17 starless clump candidates has been selected fromthe Traficante et al. (2015a) as objects with Σ > .
05 g cm − ,mass M >
300 M ⊙ , bolometric luminosity over envelope mass ra-tio L / M . <
15 K), indica-tive of very young stage of evolution (see Section 4), no (or faint)emission at 70 µ m after visual inspection of each source and nocounterparts in the MSX and WISE catalogues in correspondenceof the Herschel dust column density peak. In addition, we checkedfor di ff erent masers emission associated with these clumps, as theyare an indication of on-going star formation activity. We searchedin the methanol multibeam survey (MMB, Green et al. 2009) andfound no Class II CH OH masers in the sources of our sample(Breen et al. 2015); from the MMB survey we also searched forhydroxyl (OH) masers at 6035 MHz (Avison et al. 2016), a transi-tion often associated with high-mass star forming regions, and alsofound no associations. We searched for CH OH and OH masers us-ing also the Arecibo surveys of Olmi et al. (2014), which is moresensitive than the MMB survey and it is targeted to identify weakmasers associated with Hi-GAL high-mass objects. We found noCH OH masers at distances less then 100 ′′ from the source cen-troids. We found one source, 34.131 + ≃ σ above the r.m.s. of the observa-tions made with the Arecibo telescope, Olmi et al. 2014). Finally,we checked several surveys of water masers in the first quadrant(Merello & et al. 2017, and references therein) and found that onlyone source, 23.271-0.263, has a H O maser association (at ≃ ′′ from the source centroid), identified in the survey of Svoboda et al.(2016). Note that the source 18.787-0.286 is classified as starless inTraficante et al. (2015a) catalogue and has no maser associations,although it shows a 70 µ m counterpart. The 70 µ m source howeveris faint, with a peak emission of ≃
60 mJy / pixel compared to abackground of ≃
130 mJy / pixel. The clump follows all the otherselection criteria and has very low dust temperature (T = µ m quiet clumps observed in theGalaxy. The dust continuum properties of the clumps have been evaluatedfrom the Hi-GAL fluxes at 160, 250 350 and 500 µ m. We com-bined these data with fluxes at 870 µ m from the ATLASGAL sur- c (cid:13) , 000–000 assive 70 µ m quiet clumps I vey (Schuller et al. 2009) and at 1.1 mm from the BGPS survey(Aguirre et al. 2011).The Hi-GAL survey (Molinari et al. 2010) observed the wholeGalactic plane ( | b | ◦ , and following the Galactic warp) atwavelengths of 70, 160, 250, 350 and 500 µ m using both PACS(Gri ffi n et al. 2010) and SPIRE (Poglitsch et al. 2010) instrumentsin parallel mode. The nominal Hi-GAL spatial resolution is ≃ [5 , . , , , . ′′ at [70, 160, 250, 350, 500] µ m respectively.However, due to the fast scan speed mode adopted in the par-allel mode, the 70 and 160 µ m beams are degraded down to ≃ . ′′ and 13 . ′′ respectively. The sensitivity is ≃ [27 , , , , / sr at [70, 160, 250, 350, 500] µ m respectively (Traficante et al.2011). The data reduction follows the standard Hi-GAL data re-duction pipeline (Traficante et al. 2011), and the final maps havebeen corrected following the weighted-GLS procedure described inPiazzo et al. (2011). The maps have been calibrated in comparisonwith IRAS and Planck data as described in Bernard et al. (2010).The ATLASGAL survey (Schuller et al. 2009) covers a totalof ≃
420 square degrees of the Galactic Plane in the longituderange − ◦ l ◦ and has been carried out with the LABOCAinstrument installed in the APEX 12m telescope. The survey has aspatial resolution of ≃ . ′′ and a sensitivity of ≃
70 mJy / beam inthe | l | ◦ longitude region (Csengeri et al. 2014).The BGPS survey has covered ≃
170 square degrees of theinner Galaxy in the range − . ◦ l . ◦ , | b | ◦ andhas mapped the emission at 1.1 mm with a spatial resolution of33 ′′ and a sensitivity of 30-100 mJy / beam (Aguirre et al. 2011;Ginsburg et al. 2013).We made dedicated photometry measurements for each clumpdirectly on the maps instead of using the existent catalogues in or-der to minimize the uncertainties arising from the combination ofdi ff erent surveys. The method is described in in Section 3. We searched for sources associated with each clump in the mid-infrared (MIR) using the Spitzer surveys of the Galactic Plane at24 µ m (MIPSGAL, Carey et al. 2009), and in the range 3 . − µ m(GLIMPSE, Benjamin et al. 2003). The MIPSGAL sensitivity is ≃ / beam, while the GLIMPSE sensitivity is [0.5,0.5,2.0,5.0]at [3.6,4.5,5.8,8.0] µ m respectively (Carey et al. 2009, and refer-ences therein). Details of the data reduction are in Benjamin et al.(2003) and (Carey et al. 2009) for GLIMPSE and MIPSGAL re-spectively. Molecular line data were acquired at the IRAM 30m telescope inJune 2014 under the project 034-14. The observations have beencarried out with the On the Fly observing mode to map a 2 ′ × ′ wide region which covers the entire extension of each clump. O ff -positions has been chosen within 30 ′ from the source centroids andchecked with single pointings to verify that they were emission-free. The EMIR receiver at 3 mm was tuned at the N H + (1–0)central frequency (93.17346 GHz). This tuning includes the simul-taneous observations of the HNC (1 −
0) and HCO + −
0) emissionlines. The VESPA backend was tuned at the maximum spectralresolution, 20 kHz ( ≃ .
06 km / s), to resolve the N H + hyperfinecomponents and covers only the N H + (1 −
0) emission line. TheFast Fourier Transform Spectrometer (FTS) was tuned to cover awider range of frequencies and was used to trace the HNC and HCO + emission with a spectral resolution of 50 kHz ( ≃ .
17 km / s).The system temperature varied in the range 92 T sys
162 K.The data have been reduced with the standard GILDAS CLASS software. The average sensitivity per channel of the reduced spec-tra has been evaluated after smoothing the data to ≃ . / s andmeasuring the r.m.s. in 20 emission-free channels for each source.The 1-sigma r.m.s. per ≃ . / s channel varies in the range0 . σ .
32 K. The beam FWHM at this frequency is ≃ ′′ . The far infrared (FIR) dust photometry at wavelengths 160 λ µ m has been done using Hyper , an enhanced aperture photom-etry algorithm specifically designed for crowded regions, blendedsources and multi-wavelength analysis (Traficante et al. 2015b).The photometry process is the same adopted in Traficante et al.(2015a). For each source, a 2d-Gaussian fit at 250 µ m defines theclumps. The fit can vary to encompass a region of at least 1 FWHMat 250 µ m (18 ′′ ) and can be up to twice the 250 µ m FWHM ineach direction, to avoid the contribution from underlying filamen-tary structures. The FWHMs of the Gaussian fit define the apertureradius. This definition of the aperture region includes at least one500 µ m beam. The aperture region is used to estimate the flux at160, 250, 350 and 500 µ m for Hi-GAL and it is also used to esti-mate the flux of the ATLASGAL and BGPS counterparts directlyfrom the maps. With this choice we estimate the flux consistentlyat all wavelengths. The Hi-GAL fluxes of the clumps in the Traficante et al. (2015a)catalogue have been re-evaluated with
Hyper parameters tunedspecifically for each clump, in order to maximize the photom-etry accuracy of these highly confused regions. For 6 sources,the adapted photometry coincides with the photometry of theTraficante et al. (2015a) catalogue. For 9 clumps we perform a dif-ferent source deblending with respect to the source catalogue. In5 cases we deblended more sources to account for faint sourcesnot identified in Traficante et al. (2015a), and for the other 4 caseswe did not include any companion subtraction since the regionsare highly confused and the background estimation dominates theemission surrounding the clump. Due to the complexity of the localbackground emission associated with each source, in most casesthe fit reaches the maximum allowed size (FWHM = ′′ ) alongone direction. However, the Hyper fit converges for all sourcesbut one, 23.271-0.263. We manually forced the source aperturefor 28.792 + + Hyper fit converges but the region is heavily confused and we forced theaperture to be circular. In 23.271-0.263 the automatic fit did notconverge and we manually force the aperture region in order to en-compass at least a ≃ ′′ region in one direction. The fluxes we esti-mate di ff er for ≃
25% on average with the fluxes in Traficante et al.(2015a).The Hi-GAL fluxes have been corrected for both aperture andcolour corrections as described in Traficante et al. (2015a). For thecolour correction we consider a clump temperature of T =
11 K, theaverage temperature of the clumps (see Section 4). c (cid:13) , 000–000 A. Traficante et al.
The coordinates and photometry for all the clumps are in Ta-ble 1.
We evaluated the 870 µ m fluxes from the ATLASGAL calibratedmap for each source. For consistency, we compare our photome-try with the fluxes presented in the ATLASGAL compact sourcescatalogue (Csengeri et al. 2014). This catalogue contains ≃ ff erence in the flux estimation of thesesources between the Hyper integrated fluxes and the fluxes in theATLASGAL catalogue is ≃ ff erent ap-proaches used to evaluate the flux. The extraction method adoptedby Csengeri et al. (2014), a multi-scale wavelet filtering of the largescale structures, preserves the compact dust condensations but fil-ters out the emission arising from scales larger than ≃ ′′ . Theaperture size chosen for our clumps is up to 72 ′′ , so it is likely thatpart of the flux in the ATLASGAL catalogue has been filtered out.To test for this e ff ect and also to check the reliability of Hyper onthe ATLASGAL maps we compare the
Hyper photometry with thephotometry of the ATLASGAL catalogue in a random region of theGalactic Plane. We chose a 3 degree wide region, 21 ◦ l ◦ ,and we measure with Hyper the integrated flux in a circular regionof radius R = ′′ and R = ′′ , similar to the median of the sourcesize of the ATLASGAL sources (27 ′′ , Csengeri et al. 2014) and ofour 18 clumps (30.8 ′′ ) respectively. We identified 122 sources incommon in the region 21 ◦ l ◦ . The peak fluxes, less sensitiveto the chosen algorithm, are in excellent agreement between Hyper and the ATLASGAL catalogue (Figure 1, upper panel). The aver-age flux di ff erence is only ≃ ′′ and slightly di ff erent using an aper-ture radius of 30 ′′ (Figure 1, central and lower panel, respectively),with an average di ff erence of ≃
17% and ≃
32% respectively. Thesetests show that the
Hyper photometry on the ATLASGAL maps isreliable and that part of the large-scale flux may be filtered out inthe ATLASGAL catalogue. The source fluxes at 870 µ m are in Ta-ble 1. The 1100 µ m flux has been evaluated from the BGPS maps, whichcover 15 out of 18 sources. Three sources are not covered by theBGPS survey (22.53-0.192, 25.609 + Hyper photometry for a factor 1.46, the suggestedaperture correction for an aperture radius of 20 ′′ (Aguirre et al.2011), since the aperture radii vary between 20 ′′ and 35 ′′ (see Table1). Ten of the fifteen sources have been identified in the most re-cent version of the catalogue, the BGPSv2.1, which contains 8594compact sources (Ginsburg et al. 2013). The Hyper fluxes of these10 sources in common are in good agreement with the fluxes inthe BGPS catalogue, with a mean di ff erence of ≃ Hyper photometry and the BGPS cataloguehas been also showed in Traficante et al. (2015b), with mean fluxdi ff erences of ≃ Hyper fluxes at 1.1 mm are in Table 1.
One of the major sources of uncertainties in the flux estimationarises from the background emission associated with each clump.
Figure 1.
Flux comparison between the ATLASGAL catalogue ofCsengeri et al. (2014) and the
Hyper photometry for 122 sources identifiedin the Galactic region 21 ◦ l ◦ . Upper panel : peak flux comparison.The agreement is excellent, with a mean di ff erence of ≃ = x line. Central panel : integrated fluxcomparison using an aperture of 25 ′′ for the Hyper photometry. The blueline is the y = x line. Lower panel : the same of the central panel but using anaperture radius of 30 ′′ . The background identification and removal is particularly criti-cal for Hi-GAL data, where the cold dust emission associatedwith the background structures has the peak of its emission (e.g.Peretto et al. 2010). The uncertainties in the background estimationare less important in the ATLASGAL and BGPS data, since mostof the extended emission is filtered out in these ground-based ex-periments. Furthermore, these clumps are not isolated but foundin proximity of other contaminating sources.
Hyper does a 2d- c (cid:13)000
Hyper does a 2d- c (cid:13)000 , 000–000 assive 70 µ m quiet clumps I Clump Source RA Dec F µ m F µ m F µ m F µ m F µ m F µ m FWHM min
FWHM max
PA Deblend( ◦ ) ( ◦ ) (Jy) (Jy) (Jy) (Jy) (Jy) (Jy) ( ′′ ) ( ′′ ) ( ◦ )15.631-0.377 1 18:20:29:1 -15:31:26 1.21 3.98 4.66 2.68 0.61 0.29 28.37 36.00 126.59 018.787-0.286 1 18:26:15.3 -12:41:33 9.50 32.76 28.67 15.41 3.40 1.17 29.58 36.00 223.89 019.281-0.387 1 18:27:33.9 -12:18:17 11.38 25.53 21.38 8.38 2.23 0.88 36.00 36.00 90.0 122.53-0.192 + + + + + + These sources are not covered by the BGPS observations.
Table 1.
Photometry results of the 18 clumps studied in this work. Col.1: Clump name; Col.2: source id number as in Traficante et al. (2015a) catalogue; Cols.3 −
4: Coordinates of the clump centroids obtained from the Gaussian fit done at 250 µ m; Cols. 5 −
10: Clump fluxes at 160, 250, 350, 500, 870 and 1100 µ mrespectively; Cols. 11 −
13: minimum, maximum FWHMs and PA of the 2d-Gaussian fit. When the fit gives FWHM min = FWHM max the source is circular andthe PA if fixed to 90 ◦ ; Col. 14: deblend parameter. 1 means that one (or more) source companion has been deblended before measuring the clump flux.Survey Wavelength Sensitivity FWHM( µ m) (MJy / sr) ( ′′ )MIPSGAL 24 2.7 adapted from Carey et al. (2009) adapted from Traficante et al. (2011) adapted from Csengeri et al. (2014) adapted from Ginsburg et al. (2013) Table 2.
Sensitivity and FWHMs comparison between the surveys used inthis work: MIPSGAL, Hi-GAL, ATLASGAL and the BGPS.
Gaussian modeling of the source companions which are then sub-tracted prior to the flux evaluation. As indicated in Table 1, weperformed the source deblending in seven clumps. Finally, anothersource of uncertainty arises from the comparison of surveys withdi ff erent sensitivities and spatial resolutions, as showed in Table 2,which could potentially a ff ect the background estimation and theflux estimation of the clumps. Also, in ground-based experimentspart of the extended emission may be filtered-out, resulting in a po-tential underestimation of the source fluxes. We have assumed anerror on the Hi-GAL fluxes of 20% (as in Traficante et al. 2015a)and of 40% on the ATLASGAL and BGPS fluxes to account forthe uncertainties arising from the combination of di ff erent surveys,namely due to the di ff erent beam responses and filtering associatedwith space-based and ground-based surveys. We adopted a single-temperature greybody model to fit the sourceSEDs. The model assumes that the temperature gradient across theclump is small due to the absence of significance internal sourcesand strong interstellar radiation fields outside the clump. However,the error associated with the mass estimation assuming a single-temperature greybody model instead of solving a complete radia-tive transfer model is negligible for starless clumps (Wilcock et al.2011).The source flux S ν at frequency ν is: S ν = M κ d νν β B ν ( T ) Ω (1)where M is the source mass, d the heliocentric distance, κ thedust mass absorption coe ffi cient at reference frequency ν = − and assuming a gas / dust mass ratio of 100(Preibisch et al. 1993). B ν ( T ) is the blackbody value at temperature T and frequency ν , and Ω is the solid angle of the source. Thefree parameters of the fit are mass and temperature. We fixed thespectral index β for all the sources to β = .
0, appropriate for dense,cold clumps (e.g. Sadavoy et al. 2013).We derived the physical parameters of each clump using Hi-GAL plus ATLASGAL and (when available) BGPS fluxes. Thefits to the SEDs have been performed with the mpfit routine(Markwardt 2009). The SEDs are shown in Figure 2 and the phys-ical parameters are summarized in Table 3.The mean temperature of these clumps is < T > = . ± . < T > ≃ . c (cid:13) , 000–000 A. Traficante et al. greybody model does not take into account this variation, howeverthe very low temperatures we measure are compatible with a coldcentral core and relatively low contribution of any external radia-tion field.The mass of the clumps covers the range 269 M ⊙ , with a bolometric luminosity (evaluated in the range 24 λ µ m) 18 L
669 L ⊙ . The mean mass is ≃ . × M ⊙ for a mean bolometric luminosity of ≃
200 L ⊙ . The averageL / M ratio, a good indicator of the evolution of massive regions(Molinari et al. 2016; Cesaroni et al. 2015), is only < L / M > ≃ . ≃ / M of starless clump can-didates (1.1, Traficante et al. 2015a). L / M << − M diagramagainst the sample of starless clump candidates in Traficante et al.(2015a). The green tracks are the Molinari et al. (2008) evolution-ary tracks for single high-mass cores. This model describes the evo-lution of massive cores from the beginning of the star formationprocess prior to the formation of a zero-age main sequence star inthe L − M diagram, following the McKee & Tan (2003) accretionmodel. The high-mass cores follow a two-phases model. Duringthe initial accretion phase the luminosity increases sustained by thecollapse and the mass slightly decreases due to outflows (the anal-ogous of a Class 0 object in the low-mass regime). The sources fol-low an almost vertical path in the diagram, up to a turnover point(the analogous of a class I object) after which the sources followan almost horizontal path. This second phase begins with the for-mation of an HII region, the luminosity remains roughly constantwhile the mass is expelled through radiation and molecular out-flows. The dotted line in Figure 3 corresponds to the best fit ofthe analogous of Class 0 objects for massive cores (Molinari et al.2008). This model assumes that the accretion rate onto the centralstar increases with time. For comparison, we also show the em-pirical border between Class 0 and Class I sources (discussed in,e.g. Andr´e, Ward-Thompson & Barsony 2000; Duarte-Cabral et al.2013), which instead considers a decreasing accretion rate (blue-dotted line). In both cases, the clumps distribution lie well belowthe Class 0 regime, in a region characterized by extremely youngobjects.
In order to investigate if our clumps are likely going to form mas-sive stars we first explore their mass-radius relationship. An em-pirical high-mass star formation threshold in this diagram has beenproposed by Kau ff mann & Pillai (2010, KP), which identified aspotentially high-mass star forming regions all the clusters withM(r) >
870 M ⊙ (r / pc) . . Recently, Baldeschi et al. (2017) ana-lyzed the bias in the estimation of the physical parameters of mas-sive clumps and found a relationship M(r) > ⊙ (r / pc) . ,which is more stringent then the KP threshold. As shown in Fig-ure 4, following the KP criterion all but three clumps are abovethe threshold, 15.631-0.377, 30.357-0.837 and 32.006-0.51. Con-versely, following the Baldeschi et al. (2017) criterion 7 clumpsmay not form high-mass stars: 15.631-0.377, 30.357-0.837 and32.006-0.51 plus 19.281-0.387, 25.982-0.056, 28.792 + + Σ t belowwhich clumps may not form massive stars. The value of Σ t is still debated. Tan et al. (2014) assumes 0 . Σ t − as the rangeof values for the threshold, while e.g. Urquhart et al. (2014) identi-fied Σ t = .
05 g cm − based on the analysis of massive clumps inATLASGAL. In our sample there are no sources with Σ .
05g cm − , and 7 sources with Σ . − (Table 3). In Fig-ure 4 we also show lines of constant surface density. The clumps15.631-0.377, 30.357-0.837 and 32.006-0.51 are the sources withthe lowest mass and surface density. Four sources, 18.787-0.286,24.013 + + Σ > . − and M > × M ⊙ . These clumps are among the most mas-sive 70 µ m quiet clumps observed to date and they will potentiallyform stars with mass comparable with the most massive protostarsobserved in the Galaxy (Peretto et al. 2013; Avison et al. 2015).Combining these two criterion the majority of these clumpswill likely produce massive stars. Although 70 µ m quiet clumps are good candidates to be star-less (Dunham et al. 2008; Giannini et al. 2012; Veneziani et al.2013), it may happen that some of these clumps already embed24 µ m sources (Elia & et al. 2017), identified thanks to the bet-ter MIPSGAL sensitivity compared with Hi-GAL (Table 2). Inthe Gutermuth & Heyer (2015) 24 µ m source catalogue, 5 clumpshave indeed a 24 µ m counterpart. Two of them (18.787-0.286 and30.454-0.135) however are likely to be foreground sources, as theyare the only to have a 2MASS counterpart and they are not locatedat the Hi-GAL column density peak.In order to look for faint 24 µ m sources not identified in theGutermuth & Heyer (2015) catalogue, we visually inspected theMIPSGAL counterparts of each clump. This inspection reveals thatin 50% of the sample at least one 24 µ m counterpart is presentwithin a 250 µ m beam centered in the clump centroid. Four sources(18.787-0.286, 24.013 + + µ m counterpart within the clump region.The complete list of sources identified by eye, with their positions,is in Table 4.To investigate the properties of these counterparts we per-formed a dedicated photometry on the MIPSGAL maps with theAperture Photometry Tool package . We estimate the photometryin a circular region with a radius of 3.5 ′′ , which includes ≃ µ m beam (5.9 ′′ ) in the aperture, and applied the correspond-ing aperture correction of 2.78, as suggested in the MIPS instru-ment handbook . The background has been estimated as the medianvalue of the emission measured in a circular annulus surroundingeach source. These 24 µ m counterparts are very faint, with fluxes F in the range 4 . F .
31 mJy. The source photometryis in Table 4.To further characterize and classify these sources we look forcounterparts in the GLIMPSE survey (Benjamin et al. 2003). Eightsources have GLIMPSE counterparts within a radius of 3 ′′ fromthe 24 µ m centroid (the MIPSGAL beam), associated with 5 dif-ferent clumps, showed in Table 5. One source, 19.287-0.386, has 2GLIMPSE sources associated with the 24 µ m counterpart. All but 1source have GLIMPSE counterparts at all the four IRAC bands: 3.6,4.5, 5.8 and 8.0 µ m. 18.787-0.286 2 does not have a counterpart at http://irsa.ipac.caltech.edu/data/SPITZER/docs/mips/mipsinstrumenthandbook/50/ c (cid:13)000
31 mJy. The source photometryis in Table 4.To further characterize and classify these sources we look forcounterparts in the GLIMPSE survey (Benjamin et al. 2003). Eightsources have GLIMPSE counterparts within a radius of 3 ′′ fromthe 24 µ m centroid (the MIPSGAL beam), associated with 5 dif-ferent clumps, showed in Table 5. One source, 19.287-0.386, has 2GLIMPSE sources associated with the 24 µ m counterpart. All but 1source have GLIMPSE counterparts at all the four IRAC bands: 3.6,4.5, 5.8 and 8.0 µ m. 18.787-0.286 2 does not have a counterpart at http://irsa.ipac.caltech.edu/data/SPITZER/docs/mips/mipsinstrumenthandbook/50/ c (cid:13)000 , 000–000 assive 70 µ m quiet clumps I Figure 2.
SED fitting for the 18 clumps studied in this work.c (cid:13) , 000–000
A. Traficante et al.
Clump Mass Luminosity Radius Σ Temperature χ Distance(M ⊙ ) (L ⊙ ) (pc) (g cm − ) (K) (kpc)15.631-0.377 269(79) 18 0.54(0.05) 0.06(0.01) 9.6(0.4) 0.60 3.47(0.35)18.787-0.286 1915(550) 206 0.69(0.07) 0.27(0.06) 10.3(0.4) 2.96 4.36(0.44)19.281-0.387 701(206) 123 0.67(0.07) 0.10(0.02) 11.4(0.5) 1.61 3.82(0.38)22.53-0.192 1579(488) 349 0.80(0.08) 0.16(0.04) 11.8(0.6) 2.01 5.77(0.58)22.756-0.284 655(194) 136 0.55(0.06) 0.14(0.03) 11.8(0.5) 0.67 4.43(0.44)23.271-0.263 997(297) 285 0.72(0.07) 0.13(0.03) 12.3(0.6) 4.76 5.21(0.52)24.013 + + + + + Table 3.
Clumps properties derived from the SEDs. Col.1: Clump name; Col.2: Clumps mass with errors estimated from the SED fit and assuming a fluxuncertainties of 20% for Hi-GAL fluxes (Traficante et al. 2015a) and of 40% on ATLASGAL and BGPS fluxes, plus an uncertainties on distance estimationof 10%; Col.3: Bolometric luminosity estimated integrating the flux in the range 160-1100 µ m; Col. 4: Source radius derived from the geometric mean of theFWHMs in Table 1; Col. 5: Temperatures and associated uncertainties; Col. 6: χ of the SED fits done using Hi-GAL, ATLASGAL and, where available,BGPS datasets; Col.7: Source distance taken from Traficante et al. (2015a). Figure 3.
Bolometric luminosity vs. envelope mass. The red points are theclumps presented in this work. The grey points are all the starless clumpcandidates identified in Traficante et al. (2015a). The four tracks correspondto the evolutionary tracks for massive cores with final masses of 8, 13, 18and 28 M ⊙ , from left to right, respectively. The green-dotted line is the fit toClass 0 objects as in Molinari et al. (2008), L ∝ M . . The blue-dotted lineis the empirical border between individual Class 0 and Class I protostel-lar objects discussed in Duarte-Cabral et al. (2013, and references therein),L ∝ M . . µ m. We use the GLIMPSE fluxes to classify these clumps ac-cording to the prescriptions of Lada (1987) and Gutermuth et al.(2009). The Lada (1987) classification scheme is based on theslope α of spectral index in the IRAC bands: 0 α − α − α − S [3.6][4.5] YSO classification (Phase 2). The two classifi-cation schemes agree well and we found 6 Class I sources and 3Class II sources. 18.787-0.286 2 has a Class I and Class II source Figure 4.
Mass vs. radius distribution of the 70 µ m quiet clumps. The lightgreen dotted line delimits the empirical KP threshold for high-mass starformation in IRDCs. All but three sources lie above the threshold. The darkgreen dotted line is the Baldeschi et al. (2017) threshold. Seven clumps arebelow this more stringent threshold. The grey asterisk marks the positionof the clump embedding the massive protostar in SDC335 (Peretto et al.2013). associated with the same 24 µ m counterpart. The classification is inTable 5. At least five clumps embed Class I or Class II sources butare 70 µ m quiet in the Hi-GAL maps.The MIR fluxes can be used to estimate the properties ofthe central stars with the Robitaille et al. (2006) SED fitting tool.This tool computes radiative transfer models of young stellar ob-jects in a range of masses and evolutionary stages to model a cen-tral star, an accretion disk and an envelope (Robitaille et al. 2006;Robitaille et al. 2007). The tool provides several hundreds of mod-els, each one describing a set of physical parameters with a χ value c (cid:13) , 000–000 assive 70 µ m quiet clumps I Clump 24 µ m RA 24 µ m Dec 24 µ m flux( ◦ ) ( ◦ ) (mJy)18.787-0.286 1 + + + + + + There is also a foreground 24 µ m source with a 2MASS counterpart and a 24 µ m flux of 560mJy (Gutermuth & Heyer 2015). There is also a foreground 24 µ m source with a 2MASS counterpart and a 24 µ m flux of 82mJy (Gutermuth & Heyer 2015). Table 4.
Emission of the 24 µ m source for the clumps with a faint 24 µ mcounterpart identified. Col.1: Clump name; Col. 2-3: Coordinates of the24 µ m source counterpart; Col. 6: Integrated flux of the 24 µ m source. Theflux has been estimated in a radius of 3.5 ′′ and corrected for an aperturecorrection of 2.78. describing the goodness of the fit. The grid of models for eachsource are shown in Figure 5. In order to obtain a representativevalue for the physical parameters, we average the fit results for allmodels with χ
5. The representative values are obtained with aweighted mean, with the weight being the inverse of the χ value,similar to the procedure adopted in Grave & Kumar (2009). If thenumber of models with χ χ valueswere always above 5, we average the results of the best 50 models.The weighted values of the extinction, column density and massof the central star for each source are in Table 5. All these sourcesare highly extincted, with A V in the range 25 . A V . . Σ . − applyingthe conversion factor described in Bohlin, Savage & Drake (1978).The best fit models are compatible with stars of intermediate mass,in the range 2 . M ∗ . ⊙ .The absence of a 70 µ m counterpart in a FIR / submm brightclump is not a good indicator that the clumps are starless. Some ofthese clumps embed already formed intermediate mass stars, mostof which may still be in the process of accreting.In the next Sections we analyze the data obtained from thedense gas tracers and we explore if these data may help to distin-guish between clumps with and without 24 µ m counterparts. We detected N H + (1 − −
0) ad HCO + (1 −
0) emissionin all our clumps. We excluded from the line analysis the clump24.528-0.136 as it shows absorption features due to the contamina-tion of the o ff -position and 30.454-0.135 because the spectra showat least two components along the line of sight with similar inten-sities but we cannot separate the two sources in the dust continuumdata. The spectra of the 16 clumps are presented in Appendix B. H + The N H + (1 −
0) spectra were fitted using the CLASS hfs task,which takes into account all the hyperfine components. Total opti-cal depth τ tot , central N H + velocity v LSR and velocity dispersion σ are estimated directly from the fit, following the prescription ofthe CLASS manual. These are shown in Table 6. The total opti-cal depth varies in the range 0 . τ tot .
2. In four clumps theuncertainties associated with the estimation of the optical depth, σ τ , are such that τ tot σ τ . In two of these clumps, 28.178-0.091and 28.792 + τ tot = .
10 which is the lowest al-lowed value permitted in the hfs routine. For these four clumps weevaluated the r.m.s. of the spectrum in CLASS before and after thesubtraction of the line fit. These values di ff er for less than 10%, sug-gesting that the fits are well constrained, so we accepted the param-eters estimated from the fit. The optical depth of the main compo-nent can be recovered from τ tot as τ main = r i ∗ τ tot , with r i the relativeintensity of the main hyperfine component (0.259). The majorityof the clumps for which the fit converges are moderately opticallythin with τ main .
1. The average value is < τ main > = .
61, in agree-ment with the average optical depth of quiescent clumps associatedwith IRDCs (Sanhueza et al. 2012). Two sources have τ main > H + (1 −
0) optical depth decreasesas the clump evolves from a quiescent to an active phase (from τ quiescent = . τ active = .
5, Sanhueza et al. 2012), to increaseagain when the protostellar core forms an active HII region identi-fied by a bright 8 µ m emission ( τ red = .
8, Sanhueza et al. 2012).Lower optical depths in protostellar cores with respect to prestellarcores are also observed in low-mass star forming regions with com-parable column densities (Crapsi et al. 2005; Emprechtinger et al.2009).The excitation temperature T ex is derived assuming LTE:T ex = T ln(A − +
1) with (2) A = T b T (1 − e − τ ) + e T / T bg − = h ν/ k , T bg = . b = T ∗ A F ef f / B ef f . F ef f and B ef f are the telescope forward and beam e ffi ciency re-spectively. We assume F ef f = .
98 and B ef f = .
78 (Rygl et al.2013).T ex varies in the range 3 . . T ex . . k =
10 K, similar tothe average dust temperature of these sources, all but three clumps(28.178-0.091, 28.792 + + ex . T k . However, we assumed a filling factor of 1 for all theobservations which may overestimate the region of N H + emissionand underestimate T ex .The N H + column density has been derived as in Caselli et al.(2002) and Pineda et al. (2013), with a dipole magnetic moment µ D = . B = . H + column densities are in Table 7. The average N H + columndensity is 9.2 × cm − , with a maximum of ≃ . × cm − in 24.013 + H + column density estimates are similarto what is observed in other star forming regions (Sanhueza et al.2012; Rygl et al. 2013).The N H + abundance with respect to the H , X (N H + ), thelatter derived from the dust surface density and assuming a meanmolecular weight of 2.33, is in the range 0 . × − X (N H + ) . × − (Table 7), with an average value of < X (N H + ) > = . × − . These values are likely to be underestimated as we c (cid:13) , 000–000 A. Traficante et al. λ ( µ m) -12 -11 -10 λ F λ ( e r g s / c m / s ) λ ( µ m) -13 -12 -11 λ F λ ( e r g s / c m / s ) λ ( µ m) -12 -11 -10 λ F λ ( e r g s / c m / s ) λ ( µ m) -13 -12 -11 -10 λ F λ ( e r g s / c m / s ) λ ( µ m) -12 -11 -10 λ F λ ( e r g s / c m / s ) λ ( µ m) -13 -12 -11 -10 λ F λ ( e r g s / c m / s ) λ ( µ m) -12 -11 -10 λ F λ ( e r g s / c m / s ) λ ( µ m) -13 -12 -11 -10 λ F λ ( e r g s / c m / s ) Figure 5.
Grid of models for each clump with an embedded MIR source as obtained from the Robitaille et al. (2006) SED fitter tool. The black curve in eachpanel is the best-fit model. are assuming a filling factor of 1, however they are in agreementwith the findings in massive dense cores / clumps (Pirogov et al.2007; Rygl et al. 2013) and in massive clumps associated withIRDCs (Ragan et al. 2006; Sanhueza et al. 2012). The source withthe highest N H + abundance, X (N H + ) ≃ − , is 32.006-0.51. Thisis one of the sources with the lowest mass surface density, is belowthe KP threshold for the formation of massive stars (see Section4.1) and does not embed any 24 µ m source. However 15.631-0.377,one of the other source below the KP threshold and with no 24 µ mcounterparts, has a N H + abundance close to the average. There is no clear indication that X (N H + ) is di ff erent between sources withand without a 24 µ m counterpart. + The HNC and HCO + optical depth and gas column density cannotbe directly estimated from the data as we have only observed theHNC and HCO + (1 −
0) line, which we expect to be optically thickwithin these cold, dense regions (e.g Sanhueza et al. 2012). c (cid:13) , 000–000 assive 70 µ m quiet clumps I Figure 6.
Distribution of A V for the best-fit models of each source embedded in clumps with MIR counterparts. The blue vertical line is the weighted averagemean with weight equal to the χ value of each fit. We have estimated the HCO + and HNC column densitieswith RADEX (van der Tak et al. 2007) in a similar fashion toPeretto et al. (2013). We have well constrained measurements ofthe dust column density, gas temperature (assumed equal to the dusttemperature) and velocity dispersion (from N H + emission) for ourclumps, and the only unknown variable is the gas column density.We run RADEX iteratively assuming di ff erent values of the HCO + and HNC column densities until the evaluated radiation tempera-ture matched the measured peak temperature. We consider this tem-perature as the temperature of the main peak in the HCO + and HNCspectra. RADEX allows also the estimation of the optical depth ofthe lines. As showed in Table 7, both the HCO + and HNC lines are,within the uncertainties, optically thick in all clumps. Therefore, the measured temperature is a lower limit for the temperature ofthe main peak in both HCO + and HNC spectra, which gives a lowerlimit to the estimated gas column density. The HCO + and HNC col-umn densities with the uncertainties and the abundances relative tothe dust and to the N H + are in Table 7. Although with the strongcaveat that this procedure gives only an estimate of the gas columndensity using a single optically thick line observation, the valueswe found are comparable to those found in massive clumps (e.g.Miettinen 2014; Zhang et al. 2016). We do not find any significantdi ff erences in the HNC or HCO + abundances between clumps thathost a 24 µ m counterpart and clumps without any MIR counter-parts. Due to the uncertainties in these measurements however, wecannot give definitive conclusions on the observed trends. c (cid:13) , 000–000 A. Traficante et al.
Figure 7.
Same of Figure 6, but for the mass of central stars. + spectra Optically thick line profiles can be used to identify signatures of dy-namical activity in star forming regions. The di ff erence in velocitybetween the brightest peaks of an optically thick line and an opti-cally thin line can be used to compute the the skewness parameter δ v (Mardones et al. 1997), defined as δ v = v thick − v thin ∆ v thin (4)where v thick is the LSR velocity of the brightest HCO + or HNCpeak, v thin and ∆ v thin are respectively the LSR velocity and FWHM of N H + assumed as an optically thin line. We estimate v thick fitting2 Gaussians to each HCO + and HNC spectrum with the mpfitfun IDL routine (Markwardt 2009). In clumps without a well defineddouble peak in the HCO + or HNC spectra, v thick was estimatedwith a single Gaussian. The skewness parameters are in Table 8.Mardones et al. (1997) define | δ v | > .
25 as a significant detectionof skewness.A positive skewness parameter indicates a red asymmetry inthe spectrum, which could be interpreted as signature of outflowsactivity (e.g. Peretto, Andre & Belloche 2006). A negative δ v con-versely is indicative of blue-asymmetric spectrum, a signature ofinfall motions in both HCO + (Fuller, Williams & Sridharan 2005)and HNC (Kirk et al. 2013) spectra. Simulations of infalling high- c (cid:13) , 000–000 assive 70 µ m quiet clumps I Clump Class A V Σ M ∗ χ Models(mag) (g cm − ) (M ⊙ )18.787-0.286 1 I 56.4 ± ± ± − ± ± ± − I-II 66.5 ± ± ± − ± ± ± − + ± ± ± − + ± ± ± − ± ± ± − ± ± ± − µ m counterpart. Table 5.
Physical parameters and classification of the central star in each clump with NIR / MIR counterparts. Col 1: Clump name; Col. 2: Source classifi-cation obtained as described in the text; Col. 3: Average magnitude; Col. 4: Average column density, obtained from A V using the conversion described inBohlin, Savage & Drake (1978); Col. 5: Average mass of the central star; Col.6: χ range of the models used to estimate the weighted parameters as describedin the text; Col. 6: Number of models used to estimate the physical parameters. The average values and uncertainties associated with the estimation of A V ,column density and mass of the central star have been evaluated as the weighted mean and variance of the distributions with weights equal to the inverse ofthe χ value of each fit.Clump v LSR σ T ex τ tot (km s − ) (km s − ) (K)15.631-0.377 40.2(0.1) 0.30(0.01) 3.41(0.70) 3.75(0.73)18.787-0.286 65.7(0.1) 1.07(0.01) 5.20(0.39) 1.86(0.13)19.281-0.387 53.6(0.1) 0.47(0.03) 3.15(0.80) 5.33(1.17)22.53-0.192 76.3(0.1) 1.25(0.01) 5.01(0.13) 1.22(0.03)22.756-0.284 105.1(0.1) 0.95(0.01) 4.25(0.52) 1.62(0.19)23.271-0.263 82.3(0.1) 0.94(0.01) 8.32(2.03) 0.62(0.15)24.013 + + + + + Table 6. N H + parameters. Col. 1: Clump name; Col. 2: N H + central ve-locity; Col. 3: Velocity dispersion measured as 1 / (8ln2) / × FWHM of thehyperfine fit; Col. 4: Excitation temperature; Col. 5: Total optical depth. mass star forming regions showed however that red-asymmetricHCO + (1 −
0) spectra may be observed also in absence of outflowsactivity (Smith et al. 2013). Outflows are more reliably traced bylooking for high-velocity wing emission away from the systemicvelocity of the cloud. Since this is outside the scope of this paper,for the purpose of this work, we simply interpret a significant valueof δ v in the HCO + spectra as an indication of significant dynam-ical activity, regardless of their origin. This is the case for all but2 clumps. Five HNC spectra have | δ v | .
25, an indication thatHCO + emission traces more dynamically active gas as further ex-plored in Paper II. Note that the clump 25.982-0.056 has symmet-ric HCO + and HNC line profiles with peaks shifted from the N H + central velocity. This clumps has a skewness parameter higher thanin some asymmetric, double-peaked spectra (e.g. 28.178-0.091, seespectra in Appendix B).In the following, we restrict the analysis to clumps with asym-metric line profiles and significant blue-shifted peaks (i.e. with δ v − . + and HNC spectra, or both.The HCO + spectrum of 22.756-0.284 has the blue- and red-shiftedpeaks of the same intensity. The HNC spectra on the contrary has ablue-shifted peak, compatible with infall signatures. Three of theseclumps (22.53-0.192, 24.013 + The spectra of the seven clumps with blue-shifted peaks can beused to determine their infall velocities v in and mass accretion rates˙M. We calculated the infall velocities following the “two layers”model of Myers et al. (1996). According to this model, v in is:v in = σ v red − v blue ln + e (T blue − T dip ) / T dip + e (T red − T dip ) / T dip . (5)v red and v blue are the velocities of the red and blue peaks respec-tively, and T red and T blue their main beam temperatures. T dip isthe main beam temperature of the valley between the two peaks.The parameters for each source are in Table 9. These parametershave been obtained fitting 2 Gaussians to either the HCO + or HNCspectrum, as indicated in Table 9, with the mpfitfun IDL routine(Markwardt 2009). The spectra with the Gaussian fits are in Figure8. The infall velocities are listed in Table 10. They are in therange 0 . v in .
37 km s − , with an average v in = − .Infall velocities of massive star forming regions are in the range0 . v in , mass − (Fuller, Williams & Sridharan 2005), andsimilar velocities have been observed in massive collapsing clouds(Kirk et al. 2013; Peretto et al. 2013).The infall velocities allow us to evaluate the mass accre-tion rate ˙M = π R n H µ m H v in (Myers et al. 1996), where m H isthe hydrogen mass and n H the volume density obtained from thedust mass and assuming spherical clumps for simplicity. The ac-cretion rates we obtain are 0 . ˙M . × − M ⊙ / yr c (cid:13) , 000–000 A. Traficante et al.
Clump N(H ) N(N H + ) N(HCO + ) N(HNC) X (N H + ) X (HCO + ) X (HNC) τ (HCO + ) τ (HNC)(10 cm − ) (10 cm − ) (10 cm − ) (10 cm − ) (10 − ) (10 − ) (10 − )15.631-0.377 1.58(0.26) 4.97(0.55) 2.00 + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + + . − . + . − . + . − . + . − . + . − . + . − . + + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + . − . + + . − . + . − . + . − . + . − . + . − . + . − . Table 7. N H + and HCO + column density and abundances of our clumps. Col. 1: Clump name; Col. 2: H column density derived from the parameters inTable 3; Col. 3: N H + column density; Col. 4-5: HCO + and HCN column densities estimated iterating RADEX as described in Section 5.2; Col. 6-8: N H + ,HCO + and HNC abundances with respect to H ; Col. 9-10: Optical depth of HCO + and HNC spectra as obtained from the RADEX run.Clump δ v (HCO + ) δ v (HNC)15.631-0.837 -0.06 0.0618.787-0.286 0.71 0.5819.281-0.387 0.70 0.7522.53-0.192 -0.28 -0.1822.756-0.284 1.25 -0.5023.271-0.263 0.28 0.3124.013 + + + + + Table 8.
Skewness parameter δ v evaluated in each clump according to thedefinition of Mardones et al. (1997). Col. 1: clump name; Col. 2-3: skew-ness parameter evaluated for the HCO + and HNC spectra respectively. (see Table 10). These values are comparable with the predictedaccretion rates onto massive protostellar cores (McKee & Tan2003) and with values observed in high-mass star forming regions(Fuller, Williams & Sridharan 2005; Rygl et al. 2013; Peretto et al.2013) and individual protostellar sources (Duarte-Cabral et al.2013).The free-fall time t f f = (3 π/ (32G n H )) / , with G gravitationalconstant and n H gas column density, is 2 . − . × yr (see Ta-ble 10), significantly higher than the massive starless clump candi-dates lifetime ( . yr, Motte et al. 2007; Tackenberg et al. 2012;Svoboda et al. 2016) but consistent with the accretion timescales ofDuarte-Cabral et al. (2013). Within 1 free-fall time, assuming con-stant accretion rate equal to the value measured today, the clumpsaccrete a mass 31 . M accr . .
637 M ⊙ . A clump such as 22.53-0.192 already embeds a core with a central star of ≃ ⊙ and has the potential to accrete mass comparable with or even higher thanthe most massive core in SDC335 (Peretto et al. 2013; Avison et al.2015) within one free-fall time.Two clumps have infall signatures but no visible 24 µ m coun-terparts, 22.756-0.284 and 25.609 + t ∼ yr, the expected lifetimeof infrared-quiet high-mass protostars, Motte et al. 2007). This isan upper limit as the clump accretion rate measured today mayhave increased since the start of the collapse, and the clump mayfragment in several protostellar cores. We obtain M < . ⊙ and M < . ⊙ for 22.756-0.284 and 25.609 + µ m. µ m QUIETCLUMPS In this Section we compare various evolutionary indicators (L / Mratio, dust temperature, surface density, linewidth, N H + abundanceand mass accretion rate) to look for di ff erences between clumpswith or without a 24 µ m source. We first divide our 16 clumpswith well defined gas emission spectra in two groups: clumps with-out 24 µ m counterparts ( N
24, 8 clumps) and clumps with a 24 µ mcounterpart ( Y
24, 8 clumps). We further divide the first group intwo sub-samples: clumps with low values of the skewness pa-rameter (2 clumps, N L ), which may be considered as the lessevolved, and clumps with a significant value of the skewness pa-rameter ( N S ). The properties for the three groups are summa-rized in Table 11. c (cid:13)000
24, 8 clumps). We further divide the first group intwo sub-samples: clumps with low values of the skewness pa-rameter (2 clumps, N L ), which may be considered as the lessevolved, and clumps with a significant value of the skewness pa-rameter ( N S ). The properties for the three groups are summa-rized in Table 11. c (cid:13)000 , 000–000 assive 70 µ m quiet clumps I Clump v red v blue T red T blue T dip Line(km s − ) (km s − ) K K K22.53-0.192 79.86(0.03) 75.43(0.02) 0.58( < + + < + + + + < + Table 9.
Parameters used to estimate infall velocities adopting the Myers et al. (1996) model. The parameters have been derived fitting two Gaussians at eachHCO + spectrum showing infall signatures with the mpfitfun routine (Markwardt 2009). The uncertainties comes from the fit for all but T dip for which weassume the same uncertainties of the corresponding T red and T blue . Col. 1: Clump name; Cols. 2 −
3: velocities of the red- and blue-shifted peak respectively;Cols. 4 −
5: temperatures of the red- and blue-shifted peak respectively; Col. 6: temperature of the dip between the red- and blue-shifed peaks. Col. 7: linespectrum used to fit the Gaussians. Clump Infall vel. Accr. rate t f f M f f µ m(km s − ) (10 − M ⊙ yr − ) (10 yr) (M ⊙ )22.53-0.192 0.34( < < + < + < < < + < Table 10.
Infall parameters of the seven clumps with blue-shifted spectra. Col.1: Clump name; Col. 2: Infall velocity; Col. 3: Mass accretion rate derived fromthe infall velocity; Col. 4: Estimated free-fall time; Col. 5: Mass accreted within 1 free-fall time; Col. 6: Presence (or absence) of a 24 µ m counterpart. • L / M ratio : The first indicator, the L / M ratio, is a well-identifiedindicator of clumps evolution (Molinari et al. 2008; Molinari et al.2016; Cesaroni et al. 2015). The average values of the three groupsare L / M N L = . ± .
07, L / M N S = . ± .
08 and L / M Y = . ± .
07. The L / M ratio is very low in each group and, withinthe dispersion of the measurements, they all exhibit a very similarvalue. • Dust temperature : The dust temperature is also thought toincrease as the clump evolves, and the inner cores warm-up thedust envelope. The average temperatures of the three groups areT N L = . ± .
7, T N S = . ± . Y = . ± . / M indicator, although on average the N L clumpsare slightly colder than the N S and Y
24 clumps, within the dis-persion there is no clear indication of a trend among these threegroups. The observed 24 µ m sources may be too young to signifi-cantly alter the properties of the surrounding dust on clump scales. • Surface density : A trend of increasing surface density frommore quiescent to more evolved clumps has been observed inprevious surveys of star forming clumps (e.g. Urquhart et al.2014; Svoboda et al. 2016), although it is not well established(e.g. Rathborne et al. 2010). Combining a large sample of mas-sive clumps in the Galaxy taken from the Hi-GAL survey,Merello & et al. (2017) showed that there is no evidence of increas-ing Σ with the clumps evolution. Here, we find Σ N L = . ± . Σ N S = . ± .
06 and Σ Y = . ± .
07 g cm − . The N L clumps have the lowest values of surface densities on average, butthere is no di ff erences between N L and Y
24 clumps. • Velocity dispersion : An indicator of the evolution of mas-sive clumps and cores that can be derived from the gas prop- erties is the expected increase of the linewidth as the regionevolves (Smith et al. 2013). Average values for the three groups are σ N L = . ± . σ N S = . ± .
27 and σ Y = . ± . • Σ vs. σ : In Figure 9 we show the relation between gas veloc-ity dispersion and mass surface density of our clumps to explore ifthere is an evolutionary trend. The Pearson correlation coe ffi cient issignificant, 0.71. However the clumps with a 24 µ m source do notoccupy a specific locus of points in this plot, supporting the hypoth-esis that this correlation may be due to the dynamical propertiesof the star forming regions, and not with the clumps evolution, assuggested in some star formation models (Ballesteros-Paredes et al.2011). • N H + abundance : There is evidence that the N H + abundanceincreases as the clump evolves (Sanhueza et al. 2012). The N H + abundance for the three groups are respectively X (N H + ) N L = . ± . × − , X (N H + ) N S = . ± . × − and X (N H + ) Y = . ± . × − . Again, there is a weak indi-cation that the abundance of the N L clumps is lower than in theother two groups, but consistent within the dispersion. • Mass accretion rate : In some star formation models the ac-cretion rate is expected to increase with time (e.g. McKee & Tan2003). We have this information available for only 7 clumps, 2 N S and 5 Y
24 clumps. The mean values are ˙M N S = . ± . × − and ˙M Y = . ± . × − M ⊙ yr − respectively.Despite using only 7 clumps,this is suggestive that the mass accre-tion rate is higher in clumps with a detectable 24 µ m source than inregions with still no observable inner cores (but dynamically activeat the clump scale). If the accretion rate is increasing with time, c (cid:13) , 000–000 A. Traficante et al.
Figure 8.
Bue-shifted spectra used to estimate the infall parameters. The red line is the result of the IDL mpfitfun routine. The blue-dotted vertical lines arein correspondence of the systemic velocity of the clump determined by the N H + fit.Clump L / M T Σ σ X (N H + ) ˙Mgroup (L ⊙ / M ⊙ ) (K) (g cm − ) (km s − ) (10 − ) (10 − M ⊙ yr − ) N L ± ± ± ± ± − N S ± ± ± ± ± ± Y
24 0.20 ± ± ± ± ± ± Table 11.
Average values of various parameters and the associated dispersion for the three classes of 70 µ m quiet clumps: objects with no 24 µ m counterpartsand low skewness parameter ( N L ); clumps with no 24 µ m counterparts but significant value of the skewness parameter ( N S ); clumps with 24 µ mcounterparts ( Y / M ratio; Col. 3: dust temperature; Col. 4: surface density; Col. 5: velocity dispersion; Col. 6: N H + abundance relative to the H ; Col. 7: mass accretion rate. c (cid:13)000
Average values of various parameters and the associated dispersion for the three classes of 70 µ m quiet clumps: objects with no 24 µ m counterpartsand low skewness parameter ( N L ); clumps with no 24 µ m counterparts but significant value of the skewness parameter ( N S ); clumps with 24 µ mcounterparts ( Y / M ratio; Col. 3: dust temperature; Col. 4: surface density; Col. 5: velocity dispersion; Col. 6: N H + abundance relative to the H ; Col. 7: mass accretion rate. c (cid:13)000 , 000–000 assive 70 µ m quiet clumps I Figure 9. Σ vs. σ distribution of the clumps. Red crosses indicate sourceswith no 24 µ m counterpart. Black crosses mark sources with a faint 24 µ mcounterpart. Green and blue circles mark clumps with infall signatures andred-shifted HCO + spectra respectively. this increase must be very rapid based on these (few) points, as allthe other indicators do not yet show evidences of evolution amongthese groups of clumps.The best candidates to embed massive pre-stellar cores are thetwo N L clumps, 15.631-0-377 and 28.792-0.141. These clumpshave on average slightly di ff erent values of the evolutionary indica-tors compared to the other two groups. However, the values arecompatible among the three groups within the dispersion of themeasurements.These results suggest that 70 µ m quiet clumps are all at a verysimilar (and very early) stage of evolution. The early rise of a vis-ible 24 µ m source does not alter the properties of star forming re-gions at the clump scales. We investigated the gas and dust properties of a sample of 18 mas-sive clumps selected to be in a very early stage of massive starformation and 70 µ m quiet.The dust properties have been constrained combining datafrom the Hi-GAL, ATLASGAL and BGPS surveys. The clumpshave mass of 1 . × M ⊙ on average with 2 clumps that exceed2 × M ⊙ , and mass surface densities Σ > .
05 g cm − . Basedon the analysis of the mass surface density and the KP criterion toidentify high-mass stars precursor in IRDCs, the majority of theseclumps have the potential to form high-mass stars. The dust tem-peratures are T <
13 K, lower than the average dust temperatures ofstarless clump candidates (T ≃
15 K Traficante et al. 2015a). Theluminosity is on average ≃ × L ⊙ with L / M ≃ .
17, signifi-cantly lower than the L / M ratio below which clumps are thought tobe quiescent (L / M =
1, Molinari et al. 2016). These values in clumpsselected to be 70 µ m quiet suggest that these massive clumps are atearliest stages of star formation.The inspection of the 24 µ m maps shows that half ofthese clumps have at least one faint 24 µ m counterpart. Eight24 µ m sources associated with 5 di ff erent clumps have at leastone GLIMPSE counterpart. We used the SED fitter tool ofRobitaille et al. (2006) to get an estimate of the properties of theseMIR sources and found that they all have central stars deeply em-bedded in the clumps with 25 . A V .
93 mag. These are sources with masses 2 . . M ∗ . . ⊙ and the equivalent of low-massClass I and Class II sources.The gas dynamics has been studied analyzing the emissionof the dense gas tracers N H + (1 − −
0) and HCO + (1 −
0) in the 16 clumps for which we have well defined spec-tral line emission. The N H + emission is moderately optically thin( < τ > main = . + spectra have been used to identify infall signatures.Seven clumps have blue-shifted spectra with skewness parameter δ v − .
25. Two clumps with no visible 24 µ m sources have signsof infall, suggesting that they are in a dynamical state at the clumpscale prior to the formation of an intermediate / high mass core.The infall velocities are ≃ .
16 km s − on average, simi-lar to what is observed in other high-mass star forming regionswith hints of protostellar activity. Similarly the mass accretion rate,0 . ˙M . × − M ⊙ / yr, is comparable with other mas-sive star forming regions. With these accretion rates, a clump suchas 22.53-0.192 has the potential to form massive stars comparablewith the most massive protostellar cores observed in the Galaxy todate within one free-fall time, t f f ≃ . − . × yr. Assuminga lifetime of 10 yr, clumps with infall signatures and no 24 µ msources may embed faint, low-mass protostars not detected in theMIPSGAL survey.Finally we combine the dust properties with the gas dynamicsto discuss the evolution of these clumps and to search for di ff er-ences between clumps with and without 24 µ m counterparts. Wedivided the clumps in three groups: clumps with no 24 µ m counter-part and low values of the HCO + skewness parameter, ( N L , 2clumps), clumps with no 24 µ m counterparts but significant valueof the skewness parameter ( N S , 6 clumps), and finally objectswith at least one 24 µ m counterpart ( Y
24, 8 clumps). We found nosignificant di ff erences, within the dispersion, between these threegroups from indicators as L / M ratio, dust temperature, surface den-sity, N H + velocity dispersion and gas abundance. The only evi-dence is that the accretion rate increases from 24 µ m dark to 24 µ mbright clumps. This increase of the accretion rate may be the firstsign of evolution in massive clumps, as all the other indicators donot show any significant di ff erence between clumps with and with-out 24 µ m counterparts.We conclude that massive starless clumps are extremely rare.The lack of 70 µ m (and possibly 24 µ m) emission is a necessary, butnot su ffi cient condition to identify massive starless clumps. Mas-sive condensations may quickly form deeply embedded protostars,and the majority, if not all of these massive clumps may already har-bor low-mass fragments. High resolution observations are neededto reveal the embedded content of these high density regions. REFERENCES
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24 µm F l u x ( J y / b e a m ) D e c ( J )
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500 µm F l u x ( J y / b e a m ) Figure A14. (cid:13) , 000–000 A. Traficante et al. D e c ( J )
24 µm F l u x ( J y / b e a m ) D e c ( J )
70 µm F l u x ( J y / b e a m ) D e c ( J )
160 µm F l u x ( J y / b e a m ) D e c ( J )
250 µm F l u x ( J y / b e a m ) D e c ( J )
350 µm F l u x ( J y / b e a m ) D e c ( J )
500 µm F l u x ( J y / b e a m ) Figure A15. (cid:13)000
500 µm F l u x ( J y / b e a m ) Figure A15. (cid:13)000 , 000–000 assive 70 µ m quiet clumps I D e c ( J )
24 µm F l u x ( J y / b e a m ) D e c ( J )
70 µm F l u x ( J y / b e a m ) D e c ( J )
160 µm F l u x ( J y / b e a m ) D e c ( J )
250 µm F l u x ( J y / b e a m ) D e c ( J )
350 µm F l u x ( J y / b e a m ) D e c ( J )
500 µm F l u x ( J y / b e a m ) Figure A16. + (cid:13) , 000–000 A. Traficante et al. D e c ( J )
24 µm F l u x ( J y / b e a m ) D e c ( J )
70 µm F l u x ( J y / b e a m ) D e c ( J )
160 µm F l u x ( J y / b e a m ) D e c ( J )
250 µm F l u x ( J y / b e a m ) D e c ( J )
350 µm F l u x ( J y / b e a m ) D e c ( J )
500 µm F l u x ( J y / b e a m ) Figure A17. (cid:13)000
500 µm F l u x ( J y / b e a m ) Figure A17. (cid:13)000 , 000–000 assive 70 µ m quiet clumps I D e c ( J )
24 µm F l u x ( J y / b e a m ) D e c ( J )
70 µm F l u x ( J y / b e a m ) D e c ( J )
160 µm F l u x ( J y / b e a m ) D e c ( J )
250 µm F l u x ( J y / b e a m ) D e c ( J )
350 µm F l u x ( J y / b e a m ) D e c ( J )
500 µm F l u x ( J y / b e a m ) Figure A18. + (cid:13) , 000–000 A. Traficante et al.
APPENDIX B: N H + (1 − −
0) AND HCO + (1 − c (cid:13) , 000–000 assive 70 µ m quiet clumps I Figure B1. N H + (1 −
0) spectrac (cid:13) , 000–000 A. Traficante et al.
Figure B2. N H + (1 −
0) spectra continues c (cid:13)000
0) spectra continues c (cid:13)000 , 000–000 assive 70 µ m quiet clumps I Figure B3.
HNC (1 −
0) spectrac (cid:13) , 000–000 A. Traficante et al.
Figure B4.
HNC (1 −
0) spectra continues c (cid:13)000
0) spectra continues c (cid:13)000 , 000–000 assive 70 µ m quiet clumps I Figure B5.
HCO + (1 −
0) spectrac (cid:13) , 000–000 A. Traficante et al.
Figure B6.
HCO + (1 −
0) spectra continues c (cid:13)000