Chemical evolution of N2H+ in six massive star-forming regions
aa r X i v : . [ a s t r o - ph . GA ] M a y Chemical evolution of N H + in six massive star-forming regions Nai-Ping Yu ,Jin-Long Xu ,Jun-Jie Wang ,Xiao-Lan Liu ABSTRACT
To investigate how the abundance of N H + varies as massive clumps evolve, herewe present a multi-wavelength study toward six molecular clouds. All of these cloudscontain several massive clumps in di ff erent evolutionary stages of star formation. Us-ing archival data of Herschel InfraRed Galactic Plane Survey (Hi-GAL), we made H column density and dust temperature maps of these regions by the spectral energy dis-tribution (SED) method. We found all of the six clouds show distinct dust temperaturegradients, ranging from ∼
20 K to ∼
30 K. This makes them good candidates to studychemical evolution of molecules (such as N H + ) in di ff erent evolutionary stages of starformation. Our molecular line data come from Millimeter Astronomy Legacy TeamSurvey at 90 GHz (MALT90). We made column density and then abundance mapsof N H + . We found that when the dust temperature is above 27 K, the abundanceof N H + begins to decrease or reaches a plateau. We regard this is because that inthe photodissociation regions (PDRs) around classical HII regions, N H + is destroyedby free electrons heavily. However, when the dust temperature is below 27 K, theabundance of N H + increases with dust temperature. This seems to be inconsistentwith previous chemical models made in low-mass star-forming regions. In order tocheck out whether this inconsistency is caused by a di ff erent chemistry in high-massstar-forming clumps, higher angular resolution observations are necessary. Subject headings: astrochemistry - stars: massive - stars: formation - ISM: molecules- ISM: abundances - ISM: HII regions - ISM: clouds
1. Introduction
Massive stars ( ≥ ⊙ ) play an important role in the evolution of galaxies and molecularclouds. They release large amounts of energy into their surrounding interstellar medium (ISM) andhave an immense impact on the subsequent star formation therein. They enrich heavy elements ofthe cosmic space in the form of supernova explosions. Massive stars are rare and used to form in National Astronomical Observatories, Chinese Academy of Sciences, Beijing 100012, China ff erent due to the physical changes in di ff erent star formation processes(e.g. Sakai et al. 2008; Sanhueza et al. 2012; Vasyunina et al. 2011). Although many studies havebeen done to understand massive star formation, less is studied about their chemistry. What is thechemistry of massive star formations? Are they di ff erent from their low-mass counterparts? Canthe chemical evolution in massive star-forming regions be used as a chemical clock? Recently,several studies have focused on this subject (e.g. Vasyunina et al. 2011; Hoq et al. 2013; Miettinen2014; Yu & Wang 2015; Yu & Xu 2016 ).Compared with CO species, N H + is more resistant to freeze-out onto grains, thus it is re-garded as a good tracer of dense gas in the early stages of star formation (Bergin et al. 2001).CO would easily be depleted onto the dust grains when dust temperature is below 20 K. Lee etal. (2004) combined a sequence of Bonnor-Ebert spheres and the inside-out collapse model todescribe dynamics from the pre-protostellar stage to later stages. They found N H + is primarilyformed through the gas-phase reaction H + + N → N H + + H , and destroyed by CO moleculesin the gas phase. So it should be expected that in the early stages of star formation, the abundanceof N H + is relatively high as the depletion of CO in gas phase. As the central protostar evolves, thegas gets warm, and CO molecules begin to evaporate from the dust grains when dust temperatureis above 20 K (Tobin et al. 2013). CO could destroy N H + through the reaction N H + + CO → HCO + + N ( e.g. Jørgensen et al. 2004; Lee et al. 2004). According to the work of Vigren etal. (2012), N H + can also be destroyed by free electrons in HII regions: N H + + e − → N + Hor NH + N. Some researches have been done to check the chemical evolution of N H + in massivestar-forming regions. Using the molecular line maps from Year 1 of the MALT90 Survey, Hoq etal. (2013) found the abundance of N H + increase as a function of evolutionary stage. Sanhuezaet al. (2012) also found both the column density and abundance of N H + increase as the clumpsevolve from “quiescent” (starless candidates) to “red” stages (HII region candidates). These resultsare inconsistent with predictions of chemical models introduced above. As these authors analyze,several reasons may account this phenomena: (i) Chemical models of low-mass star forming re-gions (e.g. Bergin & Langer 1997; Lee et al. 2004) focused on single low-mass cores, whichcannot be compared with a clump (hosting tens or hundreds of cores). The Mopra beam is 38 ′′ , 3 –this makes it impossible to resolve a single massive star-forming core. The Mopra telescope notonly probes the warm core gas but also the surrounding cold di ff use material; (ii)The chemistry ofN H + in massive star-forming regions may be really di ff erent from their low-mass counterparts;(iii)The formation and destruction rate of N H + might not be as accurate as is currently believed.We should also mention here that the initial conditions may be very di ff erent in di ff erent molecularclouds. For example, in the Galaxy, the C / C ratio ranges from ∼
20 to ∼
70, depending onthe distance to the Galactic center (e.g. Savage et al. 2002). This might make their results notstatistically significant. Here we present multi-wavelength study toward six massive star-formingregions containing several clumps in di ff erent evolutionary stages of star formation, with the aimto make a more clear picture of chemical evolution of N H + in massive star-forming regions. Withonly six molecular clouds, our work cannot be statistically significant, but given that the clumpsin each cloud likely have similar initial conditions, we wish to use a di ff erent approach to foundout whether our result will be consistent with previous work. We introduce our sources and datain Section 2, analysis is given in Section 3, results and discussions are in Section 4, and finally wesummarize in Section 5.
2. Source and Data
In the followings we present our source introduction in Section 2.1, molecular data of MALT90in Section 2.2, and far-infrared data from ATLASGAL and Hi-GAL in Section 2.3.
The source sample of this paper involves two filamentary clouds identified by Li et al. (2016),two bubbles from Churchwell et al. (2006; 2007), and two dense clouds from Rathborne et al.(2016). The basic information of them is shown in table 1. We can see that each source involvesat least four dense clumps from the catalogue of Contreras et al. (2013). Guzm´an et al. (2015)classified the clumps into “Quiescent”, “Protostellar”, “HII region” and “PDRs” stages accordingto the schematic timeline of a massive star formation. All of the six sources involves at leasttwo stages. The e ff ective radius of most clumps are larger than 38 ′′ , which means they can beresolved by the Mopra telescope. For G351.776-0.527, CN148 and S36, previous studies havealready indicated they are candidates of massive star-forming regions (e.g. Klaassen et al. 2015;Dewangan et al. 2015; Torii et al. 2017). For the other three sources, all of them involves at leastone clumps with masses > M ⊙ . Given a typical star formation e ffi ciency of 10%-30% (Lada etal. 2010) and a cluster having a Salpeter-type initial stellar mass function (IMF), we could expecta 10 M ⊙ clump to form a star cluster with massive stars >
20 M ⊙ . Therefore, our sources are 4 –candidates of massive star forming regions. In the following analysis, it can be noted that all of thesix molecular clouds have distinct dust temperature gradient, ranging from ∼
20 K to ∼
30 K. Thismakes them good candidates to study chemical evolution of N H + in di ff erent evolutionary stagesof massive star formation. From figure 1 to figure 6, we show the composed Spitzer images of oursources overlaid with the ATLASGAL 870 µ m contours. We use archival molecular data from the MALT90 Survey. MALT90 is an international projectwith the aim to characterize physical and chemical properties of massive star formation in ourGalaxy (e.g., Foster et al. 2011; Foster et al. 2013; Jackson et al. 2013). This project was carriedout with the Mopra Spectrometer (MOPS) arrayed on the Mopra 22 m telescope, which is locatednear Coonabarabran in New South Wales, Australia. The full 8GHz bandwidth of MOPS was splitinto 16 zoom bands of 138 MHz, providing a velocity resolution of 0.11 km s − at frequenciesnear 90 GHz. The beamsize of Mopra is 38 ′′ at 86GHz, with a beam e ffi ciency between 0.49 at86 GHz and 0.42 at 115 GHz (Ladd et al. 2005). The target of this survey are selected from theATLASGAL clumps found by Contreras et al. (2013). The size of the data cube is 4.6 ′ × ′ ,with a step of 9 ′′ . We downloaded the data files from the MALT90 Home Page . For each source,we combined all the N H + data cube into a new data cube using the software package of CLASS(Continuum and Line Analysis Single-Disk Software). The analysis of other molecules (C H,HC N, H CO + and so on) will come in another paper. The combined images of N H + integratedemissions are also shown in figure 1-6. An example of N H + (1-0) averaged spectrum is shown infigure 7. ATLASGAL is the first systematic survey of the inner Galactic plane in the submillimetre(Siringo et al. 2009; Contreras et al. 2013). It provides high angular resolution ( ∼ column densities exceeding 10 cm − . The Hi-GAL data set is comprised of 5continuum images of the Milky Way Galaxy using the PACS (70 and 160 µ m) and SPIRE (250,350 and 500 µ m) instruments. The nominal angular resolutions ranges from 5.2 ′′ to 35.2 ′′ for 70 µ m and 500 µ m. The high-frequency components provide high angular resolution and is una ff ected http: // atoa.atnf.csiro.au / MALT90
3. Analysis3.1. Dust temperature and H column density We made H column density and dust temperature maps of each region by the spectral energydistribution (SED) method described by Wang et al. (2015). Given Hi-GAL is sensitive to low-density gas of about 10 cm − , background and / or foreground contaminations make a seriousproblem when analysing the Hi-GAL data. Following the steps described by Wang et al. (2015),we first remove the background and foreground emissions. After removing the background andforeground emissions, we re-gridded the pixels onto the same scale of 13 ′′ , and convolved all theimages to a spatial resolution of 45 ′′ which is the measured beamsize of Hi-GAL observations at500 µ m (Traficante et al. 2011). For each pixel, we use equation I ν = B ν (1 − e − τ ν ) (1)to model intensities at various wavelengths. The optical depth τ ν could be estimated through τ ν = µ H m H κ ν N H / R gd (2)We adopt a mean molecular weight per H molecule of µ H = m H is the mass of a hydrogen atom. N H is the column density.R gd is the gas-to-dust mass ratio which is set to be 100. According to Ossenkopf & Henning (1994),dust opacity per unit dust mass ( κ ν ) could be expressed as κ ν = . ν GHz ) β cm g − (3)where the value of the dust emissivity index β is fixed to 1.75 in our fitting. The two free parameters( N H and T d ) for each pixel could be fitted finally. The final resulting dust temperature and columndensity maps, which have a spatial resolution of 45 ′′ with a pixel size of 13 ′′ , are shown in figure8. H + To calculate the abundance of N H + in each pixel, we also smoothed the molecular data into anew beamsize of 45 ′′ with a new step of 13 ′′ . Assuming LTE conditions and a beam filling factor 6 –of 1, the column density of N H + in every pixel can thus be calculated through: N ( N H + ) = πν c R Q rot g u A ul exp ( E l / kT ex )1 − exp ( − h ν/ kT ex ) Z τ d ν (4)where c is the velocity of light in the vacuum, ν is the frequency of the transitions, g u is thestatistical weight of the upper level, A ul is the Einstein coe ffi cient, E l is the energy of the lowerlevel, Q rot is the partition function. For the excitation temperature of T ex , we here assume thatT ex is equal to the dust temperature derived above in each pixel. The value of R is 5 /
9, takinginto account the satellite lines corrected by their relative opacities (Sanhueza et al. 2012). We useapproximation Z τ d ν = τ − exp ( − τ ) R T mb dvJ ( T ex ) − J ( T bg ) (5)to take τ N H + into account. N H + (1-0) has 7 hyperfine components (e.g., Pagani et al. 2009; Keto& Rybicki 2010). As an example is shown in figure 7, the 7 hyperfine structures of N H + (1-0)blended into 3 groups because of turbulent line widths. We estimate the optical depth of N H + (1-0) by the method described by Purcell et al.(2009). The integrated intensities of Group 1 / Group 2should be in the ratio of 1:5, assuming that the line widths of the individual hyperfine componentsare all equal. The optical depth of N H + ( τ N H + ) can then be derived using the following equation: R T MB , Group dv R T MB , Group dv = − exp ( − . τ )1 − exp ( − τ ) (6)By solving Equation (4) and (6) in each pixel where the Group 2 emission of N H + is greater than6 σ , we got the column density maps of N(N H + ). For the uncertainties of column density, wehere only consider the errors from optical depth and integrated intensities. The mean uncertaintyis about 20%. The abundance value of N H + ( χ (N H + )) for each pixel can be calculated through χ (N H + ) = N(N H + ) / N(H ). The N H + abundance maps are shown in figure 9.
4. Results and Discussions4.1. G351.776-0.527
G351.776-0.527 is a filamentary cloud identified by Li et al. (2016). It is also known asinfrared dark cloud (IRDC) G351.77-0.51 (Simon et al. 2006). In the top panel of figure 1,dark absorption features against the Galactic mid-infrared background radiation field are obvi-ous. According to Leurini et al. (2011), the kinematic distance of this cloud is about 1 kpc. IRAS17233-3606 lies in the center of the filament, where the dust temperature is more than 30 K. Pre-vious studies indicate active massive star formation in this IRAS source (e.g. Caswell et al. 1980; 7 –Menten 1991; Leurini et al. 2009). Using high resolution observations of VLA, Klaassen et al.(2015) found a large scale outflow from IRAS 17233-3606. Some dense clumps have also foundon the northeast and southwest part of G351.776-0.527. The dust temperature there is relativelylow, indicating earlier evolutionary stages. From the top left panel of figure 9, we can see that theabundance of N H + is highest in the center and southeast side of IRAS 17233-3606. The top leftpanel of figure 10 indicates the abundance of N H + increase as the increment of dust temperaturein G351.776-0.527. This result is inconsistent with chemical models of low-mass star formation. G340.301-0.387 is also a filamentary cloud identified by Li et al. (2016). From figure 2, wecan see that most of the dense gas is in the northwest part of this cloud, where the gas is also moreevolved. The clumps in the southeast show no distinct Spitzer 8 µ m emissions, indicating the gashere is less evolved. The dust temperature map shows a distinct temperature gradient, decreasingfrom ∼
27 K in the northwest to ∼
20 K in the southeast. The N H + abundance map shows a similartrend, also decreasing from northwest to southeast. The T d - χ (N H + ) relation map in figure 10suggests a positive correlation. This trend is also inconsistent with chemical models of low-massstar formation, which suggests the abundance of N H + should decrease as star formation evolves. CN148 and S36 are two infrared bubbles found by Churchwell et al. (2006; 2007). Thearc-shaped distributions of N H + (1-0) emission, the 870 µ m dust emission, and the polycyclicaromatic hydrocarbon (PAH) features trace photodissociation regions (PDRs) around the two bub-bles. A multi-wavelength study of CN148 carried out by Dewangan et al. (2015) suggests triggeredstar formation when the bubble expands into the surrounding interstellar medium (ISM). A distinctdust temperature gradient can be noted around the two bubbles. The dust is quite hot (more than30 K) on the PDRs, and decreases to ∼
20 K outside of PDRs. According to Guzm´an et al. (2015),the clumps found by Contreras et al. (2013) in the two regions range from the “Quiescent” to PDRstages. Our study shows the abundance of N H + is relatively low on the PDRs. This is consistentwith chemical models, as N H + is regarded to be destroyed by free electrons (e.g. Dislaire et al.2012; Vigren et al. 2012). In a previous paper (Yu & Xu 2016), we also found the abundance ofN H + seems to decrease as a function of Lyman continuum fluxes (N L ) in compact HII regions, in-dicating that this molecule could be destroyed by UV photons when H II regions have formed. TheT d - χ (N H + ) relation maps in figure 10 suggest the abundance of N H + increases when the dusttemperature increases from ∼
18 K to ∼
27 K, and drops when dust temperature is more than 27 K. 8 –Again, the evolution trend of N H + is inconsistent with chemical models when dust temperature isbelow 27 K. + + These two sources are two dense clouds we selected from Rathborne et al. (2016). Eventhough they are not listed as bubbles by Churchwell et al. (2006; 2007), we can see distinct radioemissions from the Sydney University Molonglo Sky Survey (SUMSS) (843 MHz; Mauch et al.2003), indicating they are also two classical HII regions. The situation of these two sources arequite similar to CN148 and S36. The dust temperature in the two regions also decreases from PDRto the gas outside. The T d - χ (N H + ) relation maps of the two sources also show a turning pointnear 27 K. In the chemical models of low-mass star formation (e.g. Bergin & Langer 1997; Lee et al.2004), a relative enhancement of N H + abundance is expected in the cold prestellar phase, as COis thought to be depleted in starless cores. As the central star evolves, the gas gets warm andCO should evaporate from the dust grains if the dust temperature exceeds about 20 K (Tobin etal. 2013). Thus we could expect the N H + abundance decrease as a function of dust tempera-ture. In order to investigate the chemical evolution of N H + as clumps evolve, here we present amultiwavelength study toward six molecular clouds containing several clumps in di ff erent evolu-tionary stages of star formation. The initial conditions could be supposed to be the same in thesame molecular cloud. Our study indicates when dust temperature is below 27 K, the abundanceof N H + increases with dust temperature. Previous studies (e.g. Hoq et al. 2013, Sanhueza et al.2012, Miettinen 2014) also show this similar trend. The result of our study is consistent with thoseprevious studies, although we used a di ff erent approach. As Hoq et al. (2013) suggest, chemicalprocesses in massive star-forming regions may really di ff er from low-mass star formation. Themass infall rate, UV flux, and density in massive star formation regions are indeed di ff erent fromtheir low-mass counterparts. The large beam of Mopra may also be the reason for this inconsis-tency. Chemical models of low-mass star forming regions (e.g. Bergin & Langer 1997; Lee et al.2004) focused on single low-mass cores, which cannot be compared with a clump. Higher angularresolution observations and chemical models should be carried out to study the chemical evolutionof N H + in the early stages of massive star formation. In the PDRs where dust temperature is morethan 27 K, our study indicates that the abundance of N H + begins to decrease (CN148, S36 andG326.432 + + H + is prone to be destroyed by free electrons (e.g. Vigren et al. 2012;Yu & Xu 2016). Figure 14 in Sanhueza et al. (2012) and Figure 5 in Hoq et al. (2013) indicatethe increase of N H + abundance was up to the protostellar phase, and then for the “red clumps”or “PDR clumps” there was no significant increase or decrease. This phenomena is very simi-lar to our sources of G351.776-0.527 and G326.641 + H + abundance reachesa plateau and becomes more or less constant when T d >
27 K. Hoq et al. (2013) do not see thedecrease from their “Protostellar” to “HII / PDR” clumps, because most sources in their “HII / PDR”catalogue are compact HII regions, which means the size of ionized gas is no more than 0.1 pc.Given a typical distance of 3 kpc for a massive star forming region, the Mopra telescope not onlyprobes the ionized gas but also the surrounding cold di ff use material. Large scale infalls have alsobeen found in many compact HII regions (e.g. Keto & Wood χ (N H + ) when T d is above 27 K in CN148, S36 and G326.432 +
5. Summary
We present a multi-wavelength study toward six massive star-forming regions, to investigatethe chemical evolution of N H + as clumps evolve. Using archival data of Hi-GAL, we madeH column density and dust temperature maps of these regions through the SED method. Wefound all of the six sources show distinct dust temperature gradient, ranging from ∼
20 K to ∼
30 K. Previous infrared studies and the dust temperature images indicate physical and chemicalproperties are quite di ff erent in di ff erent parts of these sources. This makes them good candidatesfor us to study the chemical evolution of N H + in di ff erent evolutionary stages of massive starformation. Using molecular line data of MALT90, we made the abundance maps of N H + . Wefound that when the dust temperature is above 27 K, the abundance of N H + begins to decrease orreaches a plateau. We regard this is because that in the PDRs around classical HII regions, N H + is destroyed by electrons heavily. We also found that when the dust temperature is below 27 K, theabundance of N H + increases with dust temperature. This is inconsistent with chemical models oflow-mass star formation. In order to check out whether this inconsistency is caused by a di ff erentchemistry in high-mass star-forming clumps, higher angular resolution observations are necessary. 10 –
6. ACKNOWLEDGEMENTS
We are very grateful to the anonymous referee for his / her helpful comments and suggestions.This paper has made use of information from the ATLASGAL Database Server . The Red MSXSource survey was constructed with support from the Science and Technology Facilities Councilof the UK. The ATLASGAL project is a collaboration between the Max-Planck-Gesellschaft, theEuropean Southern Observatory (ESO) and the Universidad de Chile. This research made use ofdata products from the Millimetre Astronomy Legacy Team 90 GHz (MALT90) survey. The Mopratelescope is part of the Australia Telescope and is funded by the Commonwealth of Australia foroperation as National Facility managed by CSIRO. This paper is supported by National NaturalScience Foundation of China under grants of 11503037. REFERENCES
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This preprint was prepared with the AAS L A TEX macros v5.2.
13 –Table 1. Basic Information of our sources.Source Dist. a Associated l b R e f f b Mass c Type b Name (kpc) Clumps (deg) (deg) (pc) (M ⊙ )G351.776-0.527 1.0 AGAL351.774-00.537 351.774 -0.537 0.37 377.6 ProtostellarAGAL351.784-00.514 351.784 -0.514 0.34 461.4 ProtostellarAGAL351.744-00.577 351.744 -0.577 0.44 431.0 HII regionAGAL351.783-00.604 351.783 -0.604 0.15 40.8 QuiescentAGAL351.804-00.449 351.804 -0.449 0.38 392.4 ProtostellarG340.301-0.387 3.9 AGAL340.248-00.374 340.248 -0.374 1.77 6121.9 HII regionAGAL340.232-00.397 340.232 -0.397 0.55 220.7 QuiescentAGAL340.304-00.376 340.304 -0.376 0.89 1077.0 QuiescentAGAL340.301-00.402 340.301 -0.402 0.81 1130.1 QuiescentCN148 2.2 AGAL010.299-00.147 10.299 -0.147 0.70 923.3 HII regionAGAL010.286-00.164 10.286 -0.164 0.49 285.4 QuiescentAGAL010.323-00.161 10.323 -0.161 0.74 837.3 HII regionAGAL010.284-00.114 10.284 -0.114 0.37 615.7 QuiescentAGAL010.342-00.142 10.342 -0.142 0.44 393.0 ProtostellarAGAL010.344-00.172 10.344 -0.172 0.38 184.7 PDRAGAL010.356-00.149 10.356 -0.149 0.53 480.8 HII regionAGAL010.404-00.201 10.404 -0.201 0.44 397.2 HII regionS36 3.2 AGAL337.922-00.456 337.922 -0.456 1.27 3001.0 PDRAGAL337.916-00.477 337.916 -0.477 0.93 1927.5 ProtostellarAGAL337.934-00.507 337.934 -0.507 0.93 1153.7 PDRAGAL337.939-00.532 337.939 -0.532 0.53 400.3 HII regionAGAL337.974-00.519 337.974 -0.519 0.60 195.4 PDRG326.432 + + + + + + a Associated l b R e f f b Mass c Type b Name (kpc) Clumps (deg) (deg) (pc) (M ⊙ )AGAL326.641 + + + + + a The distance value of G351.776-0.527 comes from Leurini et al. (2011). The distance valueof CN148 comes from Dewangan et al. (2015). The distances of other sources come fromWhitaker et al. (2017). b Guzm´an et al. (2015). c Contreras et al. (2017). 15 – G a l a c t i c l a t i t ude G Fig. 1.— Top: Three colour mid-infrared image of G351.776-0.527 created using the SpitzerIRAC band filters (8.0 µ m in red, 4.5 µ m in green and 3.6 µ m in blue). The green boxes indicatethe five observed regions by MALT90. The blue circle marks IRAS 17233-3606. Bottom: Thenew combined image of N H + from MALT90 data set. The emission has been integrated from -6to 0 km s − . The ATLASGAL 870 µ m emissions (in white) are superimposed with levels 0.15,0.30, 0.60, 1.20, 2.40, 4.80 and 9.60 Jy / beam in each panel. The pluses mark the dense clumpslisted in table 1. The black circle shown in the lower right corner of this image indicates the beamsize of Mopra. The unit of the color bar on the right is in K km / s. 16 – G a l a c t i c l a t i t ude G340.301-0.387
Fig. 2.— Top: Three colour mid-infrared image of
G340.301-0.387 created using the Spitzer IRACband filters (8.0 µ m in red, 4.5 µ m in green and 3.6 µ m in blue). The green boxes indicate the fourobserved regions by MALT90. Bottom: The new combined image of N H + from MALT90 dataset. The emission has been integrated from -55 to -48 km s − . The ATLASGAL 870 µ m emissions(in white) are superimposed with levels 0.24, 0.48, 0.96, 1.92 and 3.84 Jy / beam in each panel. Thepluses mark the dense clumps listed in table 1. The white circle shown in the lower right corner ofthis image indicates the beam size of Mopra. The unit of the color bar on the right is in K km / s. 17 – G a l a c t i c l a t i t ude CN148
Fig. 3.— Top: Three colour mid-infrared image of CN148 created using the Spitzer IRAC bandfilters (8.0 µ m in red, 4.5 µ m in green and 3.6 µ m in blue). The green boxes indicate the sevenobserved regions by MALT90. Bottom: The new combined image of N H + from MALT90 dataset. The emission has been integrated from 10 to 15 km s − . The ATLASGAL 870 µ m emissions(in white) are superimposed with levels 0.27, 0.54, 1.08, 2.16 and 4.32 Jy / beam in each panel. Thepluses mark the dense clumps listed in table 1. The black circle shown in the lower right corner ofthis image indicates the beam size of Mopra. The unit of the color bar on the right is in K km / s. 18 – G a l a c t i c l a t i t ude S Fig. 4.— Top: Three colour mid-infrared image of S36 created using the Spitzer IRAC band filters(8.0 µ m in red, 4.5 µ m in green and 3.6 µ m in blue). The green boxes indicate the six observedregions by MALT90. Bottom: The new combined image of N H + from MALT90 data set. Theemission has been integrated from -42 to -36 km s − . The ATLASGAL 870 µ m emissions (inwhite) are superimposed with levels 0.27, 0.54, 1.08, 2.16 and 4.32 Jy / beam in each panel. Thepluses mark the dense clumps listed in table 1. The black circle shown in the lower right corner ofthis image indicates the beam size of Mopra. The unit of the color bar on the right is in K km / s. 19 – G a l a c t i c l a t i t ude G326.432+0.916
Fig. 5.— Top: Three colour mid-infrared image of G326.432 + µ m in red, 4.5 µ m in green and 3.6 µ m in blue). The green boxes indicatethe four observed regions by MALT90. The 843 MHz SUMSS radio continuum emissions (yellow)are superimposed with levels 0.2, 0.4,0.8, 1.6 and 3.2 Jy / beam. Bottom: The new combined imageof N H + from MALT90 data set. The emission has been integrated from -44 to -37 km s − . TheATLASGAL 870 µ m emissions (in white) are superimposed with levels 0.21, 0.42, 0.84 and 1.68Jy / beam in each panel. The pluses mark the dense clumps listed in table 1. The black circle shownin the lower right corner of this image indicates the beam size of Mopra. The unit of the color baron the right is in K km / s. 20 – G a l a c t i c l a t i t ude G326.641+0.612
Fig. 6.— Top: Three colour mid-infrared image of G326.641 + µ m in red, 4.5 µ m in green and 3.6 µ m in blue). The green boxes indicatethe four observed regions by MALT90. The 843 MHz SUMSS radio continuum emissions (yellow)are superimposed with levels 0.2, 0.4,0.8, 1.6 and 3.2 Jy / beam. Bottom: The new combined imageof N H + from MALT90 data set. The emission has been integrated from -44 to -37 km s − . TheATLASGAL 870 µ m emissions (in white) are superimposed with levels 0.42, 0.84, 1.68, 3.36 and6.72 Jy / beam in each panel. The pluses mark the dense clumps listed in table 1. The black circleshown in the lower right corner of this image indicates the beam size of Mopra. The unit of thecolor bar on the right is in K km / s. 21 –Fig. 7.— Averaged spectra of N H + (1-0) over the cloud of G351.776-00.527. The vertical dashedlines indicate the seven hyperfine structures. 22 –Fig. 8.— Left panels: H column density of the six sources built on the SED fitting pixel by pixel.The pluses mark the dense clumps listed in table 1. The unit of each color bar is in cm − . Rightpanels: Dust temperature maps in color scale derived from SED fitting pixel by pixel. The unit ofeach color bar is in K. 23 –Fig. 9.— Calculated N H + abundance maps of the six sources. The pluses mark the dense clumpslisted in table 1. The black circle shown in each image indicates the beam size of 45 ′′ . 24 –
18 20 22 24 26 28 30 32 34 36024681012
G351.776-0.527 X ( N H + ) x - T dust (K)
20 21 22 23 24 25 26 27 280246810121416182022
G340.301-0.387 X ( N H + ) x - T dust (K)
20 22 24 26 28 30 3202468101214 X ( N H + ) x - T dust (K) CN148
20 22 24 26 28 30 32 34 3602468101214 X ( N H + ) x - T dust (K) S36
18 20 22 24 26 28 30 32 34 36 3802468101214 X ( N H + ) x - T dust (K) G326.432+0.916
18 20 22 24 26 28 30 32 34 3602468101214 X ( N H + ) x - T dust (K) G326.641+0.612 Fig. 10.— Abundance of N H + plotted as a function of dust temperature in each pixel of the sixsources. The red dashed lines mark T d ==