# Characterizing the Transition from Diffuse Atomic to Dense Molecular Clouds in the Magellanic Clouds with [C II], [C I], and CO

Jorge L. Pineda, William D. Langer, Paul F. Goldsmith, Shinji Horiuchi, Thomas B. H. Kuiper, Erik Muller, Annie Hughes, Jürgen Ott, Miguel A. Requena-Torres, Thangasamy Velusamy, et al.

<figure><img src="image_1.png" /></figure>

## ► To cite this version:

Jorge L. Pineda, William D. Langer, Paul F. Goldsmith, Shinji Horiuchi, Thomas B. H. Kuiper, et al.: Characterizing the Transition from Diffuse Atomic to Dense Molecular Clouds in the Magellanic Clouds with [C II], [C I], and CO. The Astrophysical Journal, 2017, 839, 10.3847/1538-4357/aa683a. insu-03677018

HAL Id: insu-03677018

https://insu.hal.science/insu-03677018v1

Submitted on 2 Sep 2025

HAL is a multi-disciplinary open access
archive for the deposit and dissemination of sci-
entific research documents, whether they are pub-
lished or not. The documents may come from
teaching and research institutions in France or
abroad, or from public or private research centers.

L'archive ouverte pluridisciplinaire HAL, est
destinée au dépôt et à la diffusion de documents
scientifiques de niveau recherche, publiés ou non,
émanant des établissements d'enseignement et de
recherche français ou étrangers, des laboratoires
publics ou privés.

<figure><img src="image_2.png" /></figure>

Distributed under a Creative Commons Attribution 4.0 International License
<figure><img src="image_3.png" /></figure>

# Characterizing the Transition from Diffuse Atomic to Dense Molecular Clouds in the Magellanic Clouds with [C II], [C I], and CO

Jorge L. Pineda<fnref n="1" />, William D. Langer<fnref n="1" />, Paul F. Goldsmith<fnref n="1" />, Shinji Horiuchi<fnref n="2" />, Thomas B. H. Kuiper<fnref n="1" />, Erik Muller<fnref n="3" />, Annie Hughes<fnref n="4" />,<fnref n="5" />,
Jürgen Ott<fnref n="6" />, Miguel A. Requena-Torres<fnref n="7" />, Thangasamy Velusamy<fnref n="1" />, and Tony Wong<fnref n="8" />

Received 2016 April 19; revised 2017 March 7; accepted 2017 March 14; published 2017 April 21

## Abstract

We present and analyze deep Herschel/HIFI observations of the [C II] 158 µm, [C I] 609 µm, and [C I] 370 µm lines toward 54 lines of sight in the Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC). These observations are used to determine the physical conditions of the line-emitting gas, which we use to study the transition from atomic to molecular gas and from C⁺ to C⁰ to CO in their low-metallicity environments. We trace gas with molecular fractions in the range 0.1 < *f* (H₂) < 1, between those in the diffuse H₂ gas detected by UV absorption (*f* (H₂) < 0.2) and well-shielded regions in which hydrogen is essentially completely molecular. The C⁰ and CO column densities are only measurable in regions with molecular fractions *f* (H₂) > 0.45 in both the LMC and the SMC. Ionized carbon is the dominant gas-phase form of this element that is associated with molecular gas, with C⁰ and CO representing a small fraction, implying that most (89% in the LMC and 77% in the SMC) of the molecular gas in our sample is CO-dark H₂. The mean XCO conversion factors in our LMC and SMC sample are larger than the value typically found in the Milky Way. When applying a correction based on the filling factor of the CO emission, we find that the values of XCO in the LMC and SMC are closer to that in the Milky Way. The observed [C II] intensity in our sample represents about 1% of the total far-infrared intensity from the lines of sight observed in both Magellanic clouds.

**Key words:** evolution – ISM: clouds – ISM: general – ISM: structure – submillimeter: galaxies – submillimeter: ISM

## 1. Introduction

Understanding the life cycle of the interstellar medium (ISM) in galaxies is a primary goal in the study of galaxy evolution and star formation. The formation of molecular clouds, their subsequent evolution to initiate star formation, and the radiative and mechanical feedback of stars into their progenitor molecular gas that terminates or reignites star formation are fundamental aspects determining how galaxies evolve over cosmic time. Most of the star formation at early cosmological times took place in environments with reduced gas metallicity and dust-to-gas ratios (DGRs). With the implementation of Atacama Large Millimeter Array (ALMA), it is now possible to study a large number of high-redshift systems (e.g., Carilli & Walter 2013). To interpret these observations of poorly resolved high-redshift galaxies, it is necessary to understand local templates of the low-metallicity ISM that can be spatially and spectrally resolved in a level of detail that is not achievable for distant galaxies.

The best local templates for studying the life cycle of the ISM and star formation in low-metallicity environments are the Large and Small Magellanic Clouds (LMC, Z = 0.5 Z⊙, and SMC, Z = 0.2 Z⊙; Westerlund 1997). Due to their proximity (LMC: D = 50 kpc; SMC: D = 60 kpc; Hilditch et al. 2005; Schaefer 2008), the Magellanic clouds provide a unique opportunity to resolve individual clouds, allowing us to conduct detailed studies of the different phases of the ISM in

low-metallicity environments using various gas and dust tracers.

The ISM plays the critical role of fueling star formation in galaxies. Obtaining a complete census of the mass of the different phases of the ISM is thus key for understanding star formation and its effect on the evolution of galaxies. Traditionally, the distribution of ISM phases in galactic disks has been studied in Hα or radio continuum emission to trace ionized gas (e.g., Haffner et al. 2009), in the H I 21 cm line to trace atomic gas (e.g., Kalberla & Kerp 2009), and in the CO line to trace shielded molecular gas (e.g., Heyer & Dame 2015). The total dust mass of the ISM has been traced by modeling the spectral energy distribution (SED) of the dust continuum, but the conversion from dust mass to hydrogen mass is often complicated by finding appropriate values for the DGR, dust emissivity, and dust temperature, parameters that can vary with environmental conditions (Roman-Duval et al. 2014). We lack, however, observations of the very important ISM phases constituting the transition between diffuse atomic clouds and dense molecular clouds. This missing link between diffuse atomic and dense molecular gas can be isolated and characterized with observations of the [C II] 158 µm line that trace the diffuse ionized medium, warm and cold atomic clouds, clouds in transition from atomic to molecular form, and dense and warm photon-dominated regions (PDRs). In particular, the [C II] line is a tracer of the CO-dark H₂ gas (Madden et al. 1997; Grenier et al. 2005; Langer et al. 2010,

<footnote n="1">¹Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109-8099, USA; Jorge.Pineda@jpl.nasa.gov</footnote>

<footnote n="2">²CSIRO Astronomy & Space Science/NASA Canberra Deep Space Communication Complex, P.O. Box 1035, Tuggeranong ACT 2901, Australia</footnote>

<footnote n="3">³National Astronomical Observatory of Japan, Chile Observatory, Tokyo, Mitaka, 181-8588, Japan</footnote>

<footnote n="4">⁴CNRS, IRAP, 9 Av. Colonel Roche, BP 44346, 31028, Toulouse Cedex 4, France</footnote>

<footnote n="5">⁵Université de Toulouse, UPS-OMP, IRAP, 31028, Toulouse Cedex 4, France</footnote>

<footnote n="6">⁶National Radio Astronomy Observatory, P.O. Box O, 1003 Lopezville Road, Socorro, NM 87801, USA</footnote>

<footnote n="7">⁷Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA</footnote>

<footnote n="8">⁸Department of Astronomy, University of Illinois, Urbana, IL 61801, USA</footnote>
2014; Wolfire et al. 2010), which is gas in which hydrogen is molecular but carbon is ionized and thus not traced by CO but by [C II]. The CO-dark H₂ gas represents ~30% of the molecular mass of the Milky Way, and this fraction increases with galactocentric distance, which is an effect of the metallicity gradient of the Galaxy (Pineda et al. 2013).

The reduced DGR and lower abundance of species responsible for gas cooling (e.g., C and O) in metal-poor environments impact the relative distributions of different ISM phases and their thermal balance. In low-DGR environments, the dust column density required to form enough H₂ to self-shield efficiently against photodestruction, enabling the transition from atomic to molecular gas, is achieved at larger gas column densities compared with larger DGR environments (e.g., Gnedin et al. 2009; Krumholz et al. 2009; Sternberg et al. 2014). Additionally, due to the reduced attenuation of far-UV (FUV) photons by dust in low-metallicity gas, the CO molecule is more readily photodissociated, pushing the C⁺/C⁰/CO transition to higher molecular hydrogen column densities compared to higher metallicity systems (van Dishoeck & Black 1988). This effect results in an enhanced CO-to-H₂ conversion factor in the Magellanic clouds (Cohen et al. 1988; Rubio et al. 1991; Israel 1997, 2000; Bolatto et al. 2013). Observations of the distribution of [C II], [C I], and CO line emission in the Magellanic clouds are thus important tools for studying how metallicity affects the properties of the different phases of the ISM.

The [C II] line was imaged in the entire LMC with the Balloon-borne Infrared Carbon (BICE) balloon by Mochizuki et al. (1994), with 15′ angular resolution and 175 km s⁻¹ velocity resolution. These observations showed an enhanced [C II]/CO ratio compared with the Milky Way and that [C II] constitutes about 1.32% of the far-infrared (FIR) luminosity of the LMC (Rubin et al. 2009). A handful of star-forming regions in the LMC and SMC have been studied with the Kuiper Airborne Observatory both with low (Poglitsch et al. 1995; Israel et al. 1996; Israel & Maloney 2011) and with high (Boreiko & Betz 1991) velocity resolution observations. The 30 Doradus region in the LMC has been studied in detail with the PACS instrument on Herschel (Chevance et al. 2016). High velocity resolution [C II] images of H II regions in the LMC and SMC are starting to become available using Stratospheric Observatory for Far-Infrared Astronomy (SOFIA) (Okada et al. 2015; Requena-Torres et al. 2016).

The [C II] 158 µm, [C I] 609 and 370 µm fine structure, and ¹²CO and ¹³CO rotational transitions are diagnostics of the physical conditions of PDRs (e.g., Kim 2006; Minamidani et al. 2008, 2011; Pineda et al. 2008, 2012; Okada et al. 2015; Chevance et al. 2016; Lee et al. 2016). When compared with predictions from PDR models, the [C II] line constrains the strength of the FUV field and volume density. The excitation of the [C I] 609 and 370 µm lines can give independent estimates of the kinetic temperature and H₂ volume density (Stutzki et al. 1997). The ionized gas component can be characterized with the [N II] 122 and 205 µm fine structure line emission (Goldsmith et al. 2015; Langer et al. 2016), as well as with hydrogen recombination lines and free-free continuum emission in the cm wavelength regime. Characterizing the physical conditions of the gas is key for our understanding of the origin of the [C II] emission in the Magellanic clouds and will have implications for the interpretation of observations of distant galaxies.

In this paper, we present deep observations of the [C II] 158 µm, [C I] 609 µm, [C I] 370 µm, and ¹²CO J = 7 → 6 lines

toward 54 LMC and SMC lines of sight (LOSs). These data were obtained using the HIFI (de Graauw et al. 2010) instrument onboard the Herschel Space Observatory<fnref n="9" /> (Pilbratt et al. 2010). We complement this data set with observations of the *J* = 1 → 0 and *J* = 3 → 2 transitions of ¹²CO and ¹³CO from the Mopra and Atacama Pathfinder Experiment (APEX) telescopes, respectively. We base our target selection on maps of H I, 160 µm dust continuum emission, and CO emission, as well as results from the Far Ultraviolet Spectroscopic Explorer (FUSE) survey of H₂ absorption toward the Magellanic clouds (Cartledge et al. 2005). The targets are distributed throughout the LMC and SMC in order to study the spatial variations of the properties of their ISMs. By studying clouds with different physical conditions, we aim to determine the key factors that characterize the evolution of the interstellar matter in the Magellanic clouds.

This paper is organized as follows. In Section 2, we describe our sample selection and observations. We discuss the determination of the contributions of the different phases of the ISM to the observed [C II] emission in Section 3, and we estimate in Section 4 the physical parameters of different ISM components. In Section 5, we study the H₀-to-H₂ and the C⁺/C⁰/CO transitions in the LMC and SMC. We also discuss the CO-to-H₂ conversion factor and the relationship between [C II] and FIR dust emission in our sample, and we compare our single-dish observations with pencil-beam FUV and optical studies to characterize the substructure of the ISM in the Magellanic clouds. We present our conclusions in Section 6.

## 2. Sample Selection and Observations

### 2.1. Sample Selection

To characterize the transition from diffuse atomic to dense molecular clouds in the Magellanic clouds, we need to study many different clouds with different physical conditions in their low-metallicity environments. Therefore, our work results in statistical properties of the sample of clouds rather than a full understanding of an individual region. This strategy also allows us to search for locations that could be followed up with current or future observatories (e.g., SOFIA, GUSTO).

Our sample consists of 36 LOSs in the LMC and 18 in the SMC. We show the locations used in our analysis in Figure 1. The locations of our LOSs were selected to be as uniformly distributed as possible over the LMC and SMC, and they do not necessarily represent the brightest emission peaks in these galaxies. We present sample spectra representing diffuse and dense PDRs in Figure 2. The full spectral line data set used in this paper, including images of the spectra in each LOS, is available as a Herschel User Provided Data Product.<fnref n="10" /> In Tables 1 and 2, we show the observational parameters in our sample, including the H₂ column densities and molecular fractions derived in our analysis and discussed in Section 5.1. By fitting Gaussians to the observed [C II], [C I], and CO spectra, we identified 49 velocity components in the LMC and 28 velocity components in the SMC. We show the integrated intensities of the spectral lines detected in our survey for the LMC in Table 3 and for the SMC in Table 4.

Our selection of positions to sample was initially based on H I (LMC: Kim et al. 2003; SMC: Stanimirovic et al. 1999)

<footnote n="9">⁹ Herschel is an ESA space observatory with science instruments provided by European-led principal investigator consortia with important participation from NASA.</footnote>

<footnote n="10">¹⁰ http://www.cosmos.esa.int/web/herschel/user-provided-data-products</footnote>
<figure><img src="image_4.png" /><figcaption>**Figure 1.** Images showing the H<sup>0</sup> column density map (blue; Stanimirovic et al. 1999; Kim et al. 2003), Herschel 160 µm continuum emission (red; Meixner et al. 2013), and MAGMA CO line emission (green; Wong et al. 2011; E. Muller et al. 2017, in preparation) in the LMC (top) and SMC (bottom). The white circles denote the positions studied in this paper. The size of the circles does not denote the beam size employed in any of the observations presented here.</figcaption></figure>

and CO (LMC: Wong et al. 2011; SMC: Muller et al. 2010; Rubio et al. 1993) maps of the LMC and SMC. In the case of the LMC, we also used Spitzer 160 µm continuum emission (Meixner et al. 2006).<sup>11</sup> We included seven H I peaks in both the LMC and SMC that have faint or no associated CO line emission or, in the case of the LMC, 160 µm continuum. These peaks represent atomic hydrogen-dominated LOSs. We also included 12 H I peaks in the LMC that are associated with

160 µm continuum emission but are undetected in the CO maps. This sample likely traces clouds that are in transition from diffuse molecular to dense molecular clouds. We also included 10 CO peaks in the LMC and seven in the SMC representing regions that have enough column density to show CO emission. They might still have a large fraction of H<sub>2</sub> gas traced by [C II] if they are clumpy and the volume filling factor of CO cores is low. Finally, our sample also includes eight LOSs in the LMC and three in the SMC studied in UV absorption with FUSE by Cartledge et al. (2005). With these sources, we have a priori knowledge of the H<sub>2</sub> column density; thus, they can be used to compare with our determination of the H<sub>2</sub> column densities from spectral line data.

We studied whether the physical conditions derived in our sample are representative of the average conditions in the LMC and SMC. For that purpose, we compared the ratio of the total infrared (TIR) intensity (Section 5.4) to the H<sup>0</sup> column density (Section 3.2) as a function of the visual extinction in our sample with that derived in maps over the entire LMC and SMC. In Figure 3, we show the TIR/N(H<sup>0</sup>) ratio as a function of A<sub>V</sub> for the entire LMC and SMC and for the locations of our sample. The grayscale in Figure 3 represents the number of pixels at a given TIR/N(H<sup>0</sup>) and A<sub>V</sub> bin, with contour lines representing 2, 10, 100, and 1000 pixels. The TIR/N(H<sup>0</sup>) ratio is an approximate measure of the FUV radiation field per unit hydrogen atom; therefore, it is sensitive to how closely an LOS is associated with star formation, while the A<sub>V</sub> is a measure of the total column density. Thus, a diffuse LOS away from star formation would be in the lower left corner of the plot, while a warm and dense PDR close to newly formed stars would be in the upper right corner of the plot. The TIR and A<sub>V</sub> maps used here were smoothed to 60'' to match the resolution of the H I data. The bulk of the LOSs have TIR/N(H<sup>0</sup>) ≃ 10<sup>-25</sup>–10<sup>-23</sup> erg s<sup>-1</sup> sr<sup>-1</sup> and A<sub>V</sub> ≲ 2 mag in the LMC and ≃10<sup>-25.5</sup>–10<sup>-24.5</sup> erg s<sup>-1</sup> sr<sup>-1</sup> and ≲0.5 mag in the SMC. In the LMC, a large fraction (90%) of our sample has TIR/N(H<sup>0</sup>) and A<sub>V</sub> values that are similar to those of the bulk of the pixels in this galaxy. This correspondence suggests that the properties derived in our sample are representative of the average properties in the LMC. There are four LOSs that have larger TIR/N(H<sup>0</sup>) and A<sub>V</sub> values corresponding to CO peaks associated with dense PDRs, which represent a small volume fraction in the LMC. In the SMC, half the LOSs have conditions that are similar to the bulk of the pixels in this galaxy, while the other half tend to have larger values of TIR/N(H<sup>0</sup>) and A<sub>V</sub>. However, most of the [C II] emission is detected in the latter half of the sample. Thus, the conditions derived in our sample in the SMC might represent those of more active regions compared with the bulk of the SMC.

## 2.2. [C II] and [C I] Observations

We surveyed the Magellanic clouds in the [C I] <sup>3</sup>P<sub>1</sub> – <sup>3</sup>P<sub>0</sub>, [C I] <sup>3</sup>P<sub>2</sub>–<sup>3</sup>P<sub>1</sub>, and [C II] <sup>2</sup>P<sub>3/2</sub>–<sup>2</sup>P<sub>1/2</sub> fine structure lines at 492.1607, 809.3420, and 1900.5469 GHz (rest frequency), respectively, with the HIFI (de Graauw et al. 2010) instrument aboard the Herschel Space Observatory (Pilbratt et al. 2010). These observations are part of the Herschel Open Time 1 Project OT1_jpineda_1. The <sup>12</sup>CO J = 7 → 6 (806.6518 GHz) rotational line was observed simultaneously with the [C I] <sup>3</sup>P<sub>2</sub>–<sup>3</sup>P<sub>1</sub> line. There are nine LOSs where only the <sup>12</sup>CO J = 7 → 6 line was observed due to an error in the frequency

<footnote n="11"><sup>11</sup> The 160 µm continuum map, used as a proxy for higher column density warmer gas, was only available for the LMC when the sample was originally selected.</footnote>
<figure><img src="image_5.png" /><figcaption>**Figure 2.** Sample spectra of ionic, atomic, and molecular species in the LMC and SMC. We show examples of diffuse regions and dense PDRs. The shaded region in the H I spectra represents the velocity range where we assume that the H I emission is in the form of CNM (see Section 3.2).</figcaption></figure>

configuration. These LOSs correspond to diffuse LOSs where the [C I] $^3P_1–^3P_0$ line was not detected; therefore, we did not expect a detection of the usually weaker [C I] $^3P_2–^3P_1$ line.

The [C II] 1.9 THz observations were carried out with the HIFI Band 7b receiver, which employs hot electron bolometer (HEB) mixers, in the LoadChop with reference observing mode. The HEB bands in HIFI show prominent electrical standing waves that are produced between the HEB mixing element and the first low-noise amplifier. The standing wave shape is not a standard sinusoid and is difficult to remove from the resulting spectrum using standard fitting methods (Higgins & Kooi 2009). To remove these standing waves, we used a procedure available in HIPE (Ott et al. 2006) version 12, which uses a library of standing wave shapes to find the best fit to correct the observed spectrum (see Higgins 2011 for a detailed description of this method). In HIPE, we also removed residual standing waves by fitting a single sinusoidal function using the *FitHIFIFringe()* procedure. After all standing waves were removed, we exported our data to the *CLASS90*<fnref n="12" /> data analysis software, which we used to combine the two polarizations, fit the data with polynomial baselines (typically of order 3), and smooth in velocity.

The angular resolution of the [C II] observations is 12". We divided the data by a factor of 0.61 to transform it from an antenna temperature to a main-beam temperature scale.<fnref n="13" /> The data were produced by the wide band spectrometer (WBS), which has a channel width of 1 MHz (0.16 km s⁻¹ at 1.9 THz). We later smoothed the data to a resolution of 0.8 km s⁻¹. For this resolution, the average rms noise of our data is<fnref n="14" /> 0.1 K.

We observed the [C I] $^3P_1–^3P_0$ line using HIFI Band 1, while the [C I] $^3P_2–^3P_1$ and $^{12}$CO $J = 7 \rightarrow 6$ lines were observed simultaneously using HIFI Band 3. The angular resolutions are 44" and 26."5 for bands 1 and 3, respectively. We applied main-beam efficiencies of 0.651 and 0.645, respectively, to convert these data from an antenna temperature to a main-beam temperature scale. The typical rms noise for the [C I] $^3P_1–^3P_0$ line was 0.035 K in a 0.91 km s⁻¹ channel width. For the [C I] $^3P_2–^3P_1$ and $^{12}$CO $J = 7 \rightarrow 6$ lines, the typical rms noise was 0.02 K in a 1.1 km s⁻¹ channel width.

<footnote n="12">¹² http://www.iram.fr/IRAMFR/GILDAS</footnote>

<footnote n="13">¹³ The beam efficiencies of all HIFI bands are presented in http://herschel.esac.esa.int/twiki/pub/Public/HifiCalibrationWeb/HifiBeamReleaseNote_Sep2014.pdf.</footnote>

<footnote n="14">¹⁴ For the typical [C II] FWHM line width of about 3 km s⁻¹, this sensitivity limit corresponds to 1.1 × 10⁻⁶ erg s⁻¹ cm⁻² sr⁻¹. The integrated intensity in units of K km s⁻¹ can be converted to that in units of erg s⁻¹ cm⁻² sr⁻¹ using *I* [K km s⁻¹] = 1.43 × 10⁵ *I* [erg s⁻¹ cm⁻² sr⁻¹] (Goldsmith et al. 2012).</footnote>
<table><caption>**Table 1**

Derived Parameters for the LMC and SMC Sample</caption><thead><tr><th>LOS</th><th>R.A.</th><th>Decl.</th><th>N(H<sup>0</sup>)<sup>a</sup><br>log<br>(cm<sup>-2</sup>)</th><th>N(H<sup>+</sup>)<sup>b</sup><br>log<br>(cm<sup>-2</sup>)</th><th>N(H<sub>2</sub>)<sup>c</sup><br>log<br>(cm<sup>-2</sup>)</th><th>f(H<sub>2</sub>)<sup>d</sup></th><th>A<sub>v</sub><sup>e</sup><br>(mag)</th><th>Total FIR<sup>f</sup><br>log<br>(erg s<sup>-1</sup> cm<sup>-2</sup> sr<sup>-1</sup>)</th></tr><tr><td colspan="9">LMC</td></tr></thead><tbody><tr><td>Diff1_NW<sup>h</sup></td><td>5:31:59.2</td><td>-66:22:52.3</td><td>21.23 ± 0.02</td><td>20.51 ± 0.01</td><td>...</td><td>...</td><td>0.91 ± 0.36</td><td>-2.61 ± 0.04</td></tr><tr><td>Diff2_SE<sup>h</sup></td><td>4:59:35.5</td><td>-70:11:04.6</td><td>21.39 ± 0.01</td><td>20.48 ± 0.01</td><td>20.9<sup>+0.2</sup><sub>-0.1</sub></td><td>0.4<sup>+0.1</sup><sub>-0.1</sub></td><td>0.70 ± 0.28</td><td>-2.70 ± 0.04</td></tr><tr><td>Diff3_RIDGE<sup>h</sup></td><td>5:31:50.6</td><td>-71:12:41.6</td><td>21.53 ± 0.01</td><td>19.61 ± 0.10</td><td>...</td><td>...</td><td>2.00 ± 0.80</td><td>-2.82 ± 0.06</td></tr><tr><td>Diff4_NE<sup>h</sup></td><td>5:01:47.7</td><td>-65:59:05.2</td><td>21.39 ± 0.01</td><td>19.85 ± 0.06</td><td>...</td><td>...</td><td>0.38 ± 0.15</td><td>-2.91 ± 0.07</td></tr><tr><td>Diff5_SE<sup>h</sup></td><td>4:58:54.0</td><td>-69:08:29.9</td><td>21.47 ± 0.01</td><td>19.67 ± 0.09</td><td>...</td><td>...</td><td>0.42 ± 0.17</td><td>-2.78 ± 0.05</td></tr><tr><td>Diff6_NW<sup>h</sup></td><td>5:43:34.9</td><td>-67:56:08.2</td><td>21.35 ± 0.01</td><td>20.06 ± 0.03</td><td>...</td><td>...</td><td>0.46 ± 0.18</td><td>-2.95 ± 0.08</td></tr><tr><td>Diff7_NW<sup>h</sup></td><td>5:25:17.3</td><td>-67:08:03.6</td><td>21.48 ± 0.01</td><td>19.91 ± 0.05</td><td>20.9 ± 0.1</td><td>0.4<sup>+0.0g</sup><sub>-0.0</sub></td><td>0.50 ± 0.20</td><td>-2.68 ± 0.04</td></tr><tr><td>LMC10_NE<sup>i</sup></td><td>4:51:51.1</td><td>-67:05:45.0</td><td>21.58 ± 0.01</td><td>20.84 ± 0.16</td><td>...</td><td>...</td><td>2.23 ± 0.89</td><td>-2.63 ± 0.04</td></tr><tr><td>LMC11_Ridge<sup>i</sup></td><td>5:25:33.8</td><td>-69:50:16.6</td><td>21.36 ± 0.01</td><td>19.92 ± 0.05</td><td>21.3<sup>+0.3</sup><sub>-0.4</sub></td><td>0.6<sup>+0.2</sup><sub>-0.2</sub></td><td>1.21 ± 0.48</td><td>-1.77 ± 0.01</td></tr><tr><td>LMC12_SE<sup>i</sup></td><td>5:02:13.7</td><td>-69:02:16.4</td><td>21.40 ± 0.01</td><td>20.25 ± 0.02</td><td>21.5<sup>+0.2</sup><sub>-0.5</sub></td><td>0.7<sup>+0.1</sup><sub>-0.2</sub></td><td>2.94 ± 1.18</td><td>-1.91 ± 0.01</td></tr><tr><td>LMC1_NW<sup>i</sup></td><td>5:28:1.9</td><td>-67:25:14.0</td><td>21.43 ± 0.01</td><td>21.14 ± 0.00<sup>g</sup></td><td>21.6<sup>+0.0g</sup><sub>-0.1</sub></td><td>0.8<sup>+0.0g</sup><sub>-0.0</sub></td><td>5.05 ± 2.02</td><td>-1.72 ± 0.00<sup>g</sup></td></tr><tr><td>LMC2_NW<sup>i</sup></td><td>5:25:16.3</td><td>-66:24:40.8</td><td>21.58 ± 0.01</td><td>20.32 ± 0.02</td><td>21.0 ± 0.1</td><td>0.3<sup>+0.1</sup><sub>-0.1</sub></td><td>0.91 ± 0.36</td><td>-2.44 ± 0.02</td></tr><tr><td>LMC3_NW<sup>i</sup></td><td>5:20:44.8</td><td>-66:06:58.2</td><td>21.25 ± 0.02</td><td>19.74 ± 0.07</td><td>20.5<sup>+0.3</sup><sub>-0.2</sub></td><td>0.3<sup>+0.1</sup><sub>-0.1</sub></td><td>0.63 ± 0.25</td><td>-2.85 ± 0.06</td></tr><tr><td>LMC4_RIDGE<sup>i</sup></td><td>5:28:22.5</td><td>-69:28:22.5</td><td>21.49 ± 0.01</td><td>20.53 ± 0.01</td><td>21.2 ± 0.1</td><td>0.5<sup>+0.0g</sup><sub>-0.0</sub></td><td>1.79 ± 0.72</td><td>-2.18 ± 0.01</td></tr><tr><td>LMC5_SE<sup>i</sup></td><td>5:06:23.1</td><td>-70:28:08.7</td><td>21.26 ± 0.02</td><td>20.47 ± 0.01</td><td>21.2 ± 0.1</td><td>0.6<sup>+0.0g</sup><sub>-0.1</sub></td><td>1.35 ± 0.54</td><td>-2.08 ± 0.01</td></tr><tr><td>LMC7_RIDGE<sup>i</sup></td><td>5:45:6.9</td><td>-69:50:42.6</td><td>21.74 ± 0.01</td><td>20.49 ± 0.01</td><td>21.2 ± 0.1</td><td>0.4<sup>+0.0g</sup><sub>-0.0</sub></td><td>1.00 ± 0.40</td><td>-2.20 ± 0.01</td></tr><tr><td>LMC8_RIDGE<sup>i</sup></td><td>5:47:11.8</td><td>-69:28:35.1</td><td>21.81 ± 0.00<sup>g</sup></td><td>20.52 ± 0.01</td><td>...</td><td>...</td><td>1.75 ± 0.70</td><td>-2.20 ± 0.01</td></tr><tr><td>LMC9_NE<sup>i</sup></td><td>5:03:20.8</td><td>-67:11:44.2</td><td>21.61 ± 0.01</td><td>20.17 ± 0.03</td><td>20.5<sup>+0.3</sup><sub>-0.2</sub></td><td>0.1<sup>+0.1</sup><sub>-0.2</sub></td><td>1.56 ± 0.62</td><td>-2.63 ± 0.04</td></tr><tr><td>NT127<sup>j</sup></td><td>5:24:19.8</td><td>-70:27:48.7</td><td>21.05 ± 0.03</td><td>19.49 ± 0.13</td><td>21.1<sup>+0.4</sup><sub>-0.7</sub></td><td>0.7<sup>+0.2</sup><sub>-0.2</sub></td><td>1.23 ± 0.49</td><td>-2.74 ± 0.05</td></tr><tr><td>NT2_NE<sup>j</sup></td><td>4:47:36.8</td><td>-67:12:13.7</td><td>21.35 ± 0.01</td><td>...</td><td>...</td><td>...</td><td>1.32 ± 0.53</td><td>-3.01 ± 0.09</td></tr><tr><td>NT74<sup>j</sup></td><td>5:14:33.4</td><td>-70:10:51.9</td><td>21.33 ± 0.01</td><td>20.87 ± 0.15</td><td>21.4<sup>+0.6</sup><sub>-0.3</sub></td><td>0.7<sup>+0.3</sup><sub>-0.1</sub></td><td>1.11 ± 0.45</td><td>-2.49 ± 0.03</td></tr><tr><td>NT77<sup>j</sup></td><td>5:15:9.4</td><td>-70:35:42.0</td><td>21.00 ± 0.03</td><td>20.66 ± 0.24</td><td>21.3<sup>+0.2</sup><sub>-0.3</sub></td><td>0.8<sup>+0.1</sup><sub>-0.1</sub></td><td>1.82 ± 0.73</td><td>-1.56 ± 0.00<sup>g</sup></td></tr><tr><td>NT97<sup>j</sup></td><td>5:19:27.8</td><td>-71:13:52.4</td><td>21.03 ± 0.03</td><td>19.52 ± 0.12</td><td>...</td><td>...</td><td>0.44 ± 0.18</td><td>-3.04 ± 0.10</td></tr><tr><td>NT99<sup>j</sup></td><td>5:19:57.6</td><td>-70:42:21.7</td><td>21.17 ± 0.02</td><td>20.87 ± 0.15</td><td>...</td><td>...</td><td>0.68 ± 0.27</td><td>-2.83 ± 0.06</td></tr><tr><td>PDR1_NW<sup>j</sup></td><td>5:25:46.9</td><td>-66:13:41.6</td><td>21.71 ± 0.01</td><td>20.40 ± 0.01</td><td>21.7<sup>+0.3</sup><sub>-0.6</sub></td><td>0.6<sup>+0.2</sup><sub>-0.2</sub></td><td>3.27 ± 1.31</td><td>-1.93 ± 0.01</td></tr><tr><td>PDR2_NW<sup>j</sup></td><td>5:35:22.4</td><td>-67:35:00.5</td><td>21.37 ± 0.01</td><td>21.78 ± 0.00<sup>g</sup></td><td>22.2 ± 0.1</td><td>0.9<sup>+0.0g</sup><sub>-0.0</sub></td><td>3.38 ± 1.35</td><td>-0.51 ± 0.00<sup>g</sup></td></tr><tr><td>PDR3_NE<sup>j</sup></td><td>4:52:8.3</td><td>-66:55:13.7</td><td>21.50 ± 0.01</td><td>21.47 ± 0.00<sup>g</sup></td><td>22.5 ± 0.1</td><td>1.0<sup>+0.0g</sup><sub>-0.0</sub></td><td>8.84 ± 3.54</td><td>-0.33 ± 0.00<sup>g</sup></td></tr><tr><td>PDR4_RIDGE<sup>j</sup></td><td>5:39:48.7</td><td>-71:09:27.4</td><td>21.65 ± 0.01</td><td>20.43 ± 0.01</td><td>21.5<sup>+0.2</sup><sub>-0.4</sub></td><td>0.6<sup>+0.1</sup><sub>-0.2</sub></td><td>11.70 ± 4.68</td><td>-1.42 ± 0.00<sup>g</sup></td></tr><tr><td>SK-66D35<sup>k</sup></td><td>4:57:4.5</td><td>-66:34:38.0</td><td>21.49 ± 0.01</td><td>20.64 ± 0.00<sup>g</sup></td><td>21.1<sup>+0.1</sup><sub>-0.2</sub></td><td>0.6<sup>+0.3</sup><sub>-0.1</sub></td><td>1.40 ± 0.56</td><td>-2.29 ± 0.02</td></tr><tr><td>SK-67D2<sup>k</sup></td><td>4:47:4.4</td><td>-67:06:53.0</td><td>21.31 ± 0.01</td><td>20.13 ± 0.03</td><td>...</td><td>...</td><td>0.31 ± 0.12</td><td>-2.61 ± 0.04</td></tr><tr><td>SK-68D129<sup>k</sup></td><td>5:36:26.8</td><td>-68:57:32.0</td><td>21.55 ± 0.01</td><td>20.66 ± 0.00<sup>g</sup></td><td>...</td><td>...</td><td>0.98 ± 0.39</td><td>-2.34 ± 0.02</td></tr><tr><td>SK-68D140<sup>k</sup></td><td>5:38:57.3</td><td>-68:56:53.0</td><td>21.69 ± 0.01</td><td>21.20 ± 0.00<sup>g</sup></td><td>21.0 ± 0.1</td><td>0.3<sup>+0.1</sup><sub>-0.1</sub></td><td>0.45 ± 0.18</td><td>-1.63 ± 0.00<sup>g</sup></td></tr><tr><td>SK-68D155<sup>k</sup></td><td>5:42:54.9</td><td>-68:56:54.0</td><td>21.76 ± 0.01</td><td>20.81 ± 0.00<sup>g</sup></td><td>21.4 ± 0.0</td><td>0.5<sup>+0.0g</sup><sub>-0.0</sub></td><td>0.92 ± 0.37</td><td>-1.77 ± 0.01</td></tr><tr><td>SK-68D26<sup>k</sup></td><td>5:01:32.2</td><td>-68:10:43.0</td><td>21.24 ± 0.02</td><td>19.91 ± 0.05</td><td>...</td><td>...</td><td>0.61 ± 0.24</td><td>-2.86 ± 0.06</td></tr><tr><td>SK-69D228<sup>k</sup></td><td>5:37:9.2</td><td>-69:20:20.0</td><td>21.52 ± 0.01</td><td>20.85 ± 0.00<sup>g</sup></td><td>20.9 ± 0.1</td><td>0.3<sup>+0.1</sup><sub>-0.1</sub></td><td>0.48 ± 0.19</td><td>-2.17 ± 0.01</td></tr><tr><td>SK-69D279<sup>k</sup></td><td>5:41:44.7</td><td>-69:35:15.0</td><td>21.62 ± 0.01</td><td>20.41 ± 0.01</td><td>...</td><td>...</td><td>0.64 ± 0.26</td><td>-2.61 ± 0.04</td></tr></tbody></table>

**Notes.**

<sup>a</sup> H<sup>0</sup> column density derived from the integrated intensity of the H I 21 cm line over the full velocity range (see Section 3.2).

<sup>b</sup> H<sup>+</sup> column density derived from Hα observations as described in Section 3.1.

<sup>c</sup> H<sub>2</sub> column density derived from [C II], [C I], and CO observations as described in Section 5.1.

<sup>d</sup> Molecular fraction, $f(H_2) = 2N(H_2)/(N(H^0) + 2N(H_2))$ (Section 5.1).

<sup>e</sup> Visual extinction derived from dust continuum maps as described in Section 4.1.

<sup>f</sup> Total FIR intensity derived from Spitzer 24 µm and Herschel 100 µm maps as described in Section 5.4.

<sup>g</sup> The uncertainty is below 0.005 or 0.05.

<sup>h</sup> H I peak (see Section 2).

<sup>i</sup> 160 µm peak (see Section 2).

<sup>j</sup> CO peak (see Section 2).

<sup>k</sup> FUSE LOSs (see Section 2).

## 2.3. CO Observations

To complement the Herschel data, we observed the $J = 1 \rightarrow 0$ transitions of $^{12}$CO, $^{13}$CO, and C$^{18}$O with the Australia Telescope National Facility (ATNF) Mopra$^{15}$ Telescope

(project M580). The Mopra 22 m telescope has an angular resolution of 33" at 115 GHz. Typical system temperatures were 600, 300, and 250 K for $^{12}$CO, $^{13}$CO, and C$^{18}$O, respectively. To convert from an antenna to a main-beam temperature scale, we used a main-beam efficiency of 0.42 (Ladd et al. 2005). All lines were observed simultaneously with the Mopra spectrometer in zoom mode. The spectra were smoothed in velocity to 0.87 km s$^{-1}$ for $^{12}$CO and 0.91 km s$^{-1}$

<footnote>$^{15}$ The Mopra radio telescope is part of the Australia Telescope, which is funded by the Commonwealth of Australia for operation as a national facility managed by CSIRO.</footnote>
<table><caption>**Table 2**

Derived Parameters for the LMC and SMC Sample</caption><thead><tr><th>LOS</th><th>R.A.</th><th>Decl.</th><th>N(H<sup>0</sup>)<sup>a</sup><br>log<br>(cm<sup>-2</sup>)</th><th>N(H<sup>+</sup>)<sup>b</sup><br>log<br>(cm<sup>-2</sup>)</th><th>N(H<sub>2</sub>)<sup>c</sup><br>log<br>(cm<sup>-2</sup>)</th><th>f(H<sub>2</sub>)<sup>d</sup></th><th>A<sub>V</sub><sup>e</sup><br>(mag)</th><th>Total FIR<sup>f</sup><br>log<br>(erg s<sup>-1</sup> cm<sup>-2</sup>sr<sup>-1</sup>)</th></tr><tr><td></td><td></td><td></td><td colspan="6" style="text-align: center;">SMC</td></tr><tr><th></th><th>J2000</th><th>J2000</th><th></th><th></th><th></th><th></th><th></th><th></th></tr></thead><tbody><tr><td>AzV18<sup>k</sup></td><td>0:47:13.1</td><td>-73:06:25.0</td><td>22.10 ± 0.00<sup>g</sup></td><td>20.54 ± 0.00<sup>g</sup></td><td>21.9 ± 0.1</td><td>0.5<sup>+0.1g</sup><sub>-0.0</sub></td><td>0.55 ± 0.22</td><td>-1.88 ± 0.01</td></tr><tr><td>AzV456<sup>k</sup></td><td>1:10:55.8</td><td>-72:42:55.0</td><td>21.64 ± 0.00<sup>g</sup></td><td>20.52 ± 0.01</td><td>...</td><td>...</td><td>0.09 ± 0.04</td><td>-3.15 ± 0.12</td></tr><tr><td>AzV462<sup>k</sup></td><td>1:11:25.9</td><td>-72:32:21.0</td><td>21.62 ± 0.00<sup>g</sup></td><td>19.77 ± 0.07</td><td>...</td><td>...</td><td>0.00 ± 0.00<sup>g</sup></td><td>-3.54 ± 0.30</td></tr><tr><td>SMC_HL_1<sup>h</sup></td><td>0:58:40.5</td><td>-72:34:52.4</td><td>21.99 ± 0.00<sup>g</sup></td><td>20.10 ± 0.03</td><td>...</td><td>...</td><td>0.28 ± 0.11</td><td>-3.04 ± 0.10</td></tr><tr><td>SMC_HL_2<sup>h</sup></td><td>0:57:35.7</td><td>-72:48:56.3</td><td>21.99 ± 0.00<sup>g</sup></td><td>20.11 ± 0.03</td><td>...</td><td>...</td><td>0.40 ± 0.16</td><td>-2.98 ± 0.08</td></tr><tr><td>SMC_HL_3<sup>h</sup></td><td>0:53:2.0</td><td>-73:15:27.9</td><td>22.02 ± 0.00<sup>g</sup></td><td>20.27 ± 0.02</td><td>...</td><td>...</td><td>0.58 ± 0.23</td><td>-2.82 ± 0.06</td></tr><tr><td>SMC_HL_4<sup>h</sup></td><td>0:48:41.5</td><td>-73:06:08.3</td><td>22.13 ± 0.00<sup>g</sup></td><td>20.36 ± 0.01</td><td>21.8 ± 0.1</td><td>0.5<sup>+0.0g</sup><sub>-0.0</sub></td><td>0.85 ± 0.34</td><td>-2.11 ± 0.01</td></tr><tr><td>SMC_HL_5<sup>h</sup></td><td>0:47:20.4</td><td>-73:18:28.4</td><td>22.05 ± 0.00<sup>g</sup></td><td>20.63 ± 0.01</td><td>...</td><td>...</td><td>0.73 ± 0.29</td><td>-2.73 ± 0.05</td></tr><tr><td>SMC_HL_6<sup>h</sup></td><td>0:49:53.2</td><td>-72:56:15.6</td><td>22.05 ± 0.00<sup>g</sup></td><td>20.48 ± 0.01</td><td>...</td><td>...</td><td>0.58 ± 0.23</td><td>-2.76 ± 0.05</td></tr><tr><td>SMC_HL_7<sup>h</sup></td><td>0:49:38.1</td><td>-73:01:14.2</td><td>22.06 ± 0.00<sup>g</sup></td><td>20.16 ± 0.03</td><td>...</td><td>...</td><td>0.35 ± 0.14</td><td>-2.80 ± 0.06</td></tr><tr><td>SMC_B2_6<sup>i</sup></td><td>0:47:57.2</td><td>-73:17:16.4</td><td>22.00 ± 0.00<sup>g</sup></td><td>20.98 ± 0.00<sup>g</sup></td><td>22.2 ± 0.1</td><td>0.7<sup>+0.0g</sup><sub>-0.0</sub></td><td>0.48 ± 0.19</td><td>-1.79 ± 0.01</td></tr><tr><td>SMC_LIRS36<sup>j</sup></td><td>0:46:40.3</td><td>-73:06:10.5</td><td>22.04 ± 0.00<sup>g</sup></td><td>20.96 ± 0.00<sup>g</sup></td><td>22.2 ± 0.2</td><td>0.7<sup>+0.1</sup><sub>-0.1</sub></td><td>0.60 ± 0.24</td><td>-1.53 ± 0.00<sup>g</sup></td></tr><tr><td>SMC_LIRS49<sup>j</sup></td><td>0:48:21.1</td><td>-73:05:29.0</td><td>22.12 ± 0.00<sup>g</sup></td><td>20.53 ± 0.00<sup>g</sup></td><td>22.5<sup>+0.1</sup><sub>-0.2</sub></td><td>0.8<sup>+0.0g</sup><sub>-0.1</sub></td><td>0.90 ± 0.36</td><td>-1.69 ± 0.00<sup>g</sup></td></tr><tr><td>SMC_NE_1a<sup>i</sup></td><td>0:59:43.8</td><td>-71:44:47.0</td><td>21.60 ± 0.00<sup>g</sup></td><td>20.23 ± 0.02</td><td>21.5<sup>+0.1</sup><sub>-0.2</sub></td><td>0.6<sup>+0.1</sup><sub>-0.1</sub></td><td>0.41 ± 0.16</td><td>-2.38 ± 0.02</td></tr><tr><td>SMC_NE_3c<sup>j</sup></td><td>1:03:30.0</td><td>-71:57:00.0</td><td>21.78 ± 0.00<sup>g</sup></td><td>20.48 ± 0.01</td><td>21.5<sup>+0.3</sup><sub>-0.8</sub></td><td>0.5<sup>+0.2</sup><sub>-0.4</sub></td><td>0.46 ± 0.19</td><td>-2.70 ± 0.04</td></tr><tr><td>SMC_NE_3g<sup>j</sup></td><td>1:03:9.9</td><td>-72:03:46.9</td><td>21.79 ± 0.00<sup>g</sup></td><td>20.84 ± 0.00<sup>g</sup></td><td>21.4<sup>+0.1</sup><sub>-0.2</sub></td><td>0.5<sup>+0.1</sup><sub>-0.1</sub></td><td>0.82 ± 0.33</td><td>-2.29 ± 0.02</td></tr><tr><td>SMC_NE_4a Hi<sup>j</sup></td><td>0:57:0.0</td><td>-72:22:40.0</td><td>21.86 ± 0.00<sup>g</sup></td><td>20.21 ± 0.02</td><td>22.1<sup>+0.6</sup><sub>-1.9</sub></td><td>0.8<sup>+0.3</sup><sub>-0.2</sub></td><td>1.13 ± 0.45</td><td>-2.77 ± 0.05</td></tr><tr><td>SMC_NE_4c low<sup>j</sup></td><td>0:58:40.0</td><td>-72:27:40.0</td><td>21.93 ± 0.00<sup>g</sup></td><td>20.39 ± 0.02</td><td>22.2 ± 0.7</td><td>0.8<sup>+0.4</sup><sub>-0.2</sub></td><td>0.67 ± 0.27</td><td>-2.63 ± 0.04</td></tr></tbody></table>

**Notes.**

<sup>a</sup> H<sup>0</sup> column density derived from the integrated intensity of the H I 21 cm line over the full velocity range (see Section 3.2).

<sup>b</sup> H<sup>+</sup> column density derived from Hα observations as described in Section 3.1.

<sup>c</sup> H<sub>2</sub> column density derived from [C II], [I], and CO observations as described in Section 5.1.

<sup>d</sup> Molecular fraction, $f(H_2) = 2N(H_2)/(N(H^0) + 2N(H_2))$ (Section 5.1).

<sup>e</sup> Visual extinction derived from dust continuum maps as described in Section 4.1.

<sup>f</sup> Total FIR intensity derived from *Spitzer* 24 µm and *Herschel* 100 µm maps as described in Section 5.4.

<sup>g</sup> The uncertainty is below 0.005 or 0.05.

<sup>h</sup> H I peak (see Section 2).

<sup>i</sup> 160 µm peak (see Section 2).

<sup>j</sup> CO peak (see Section 2).

<sup>k</sup> FUSE LOSs (see Section 2).

for <sup>13</sup>CO. The typical rms noise was 0.06 K for <sup>12</sup>CO and 0.05 K for <sup>13</sup>CO. The C<sup>18</sup>O line was detected in only a small number of LOSs, and we did not use these observations in the present analysis. We checked pointing accuracy every 60 minutes using a nearby SiO maser.

We used the APEX<sup>16</sup> 12 m telescope (Güsten et al. 2006) to observe the <sup>12</sup>CO J = 3 → 2 and <sup>13</sup>CO J = 3 → 2 lines in 25 LOSs in the LMC and SMC that have been detected in our low-J observations (project M0036_93). The angular resolution of the APEX telescope is 17"/5 at 345 GHz. We converted from an antenna to a main-beam temperature scale using an efficiency of 0.69 for 345 GHz. The typical rms noise in the <sup>12</sup>CO J = 3 → 2 data is 0.08 K for a 0.33 km s<sup>-1</sup> channel width, and the value in the <sup>13</sup>CO J = 3 → 2 data is 0.05 K for a 0.33 km s<sup>-1</sup> channel width.

## 2.4. H I, Hα, H66α, and Dust Continuum Data

We used the H I 21 cm maps of the entire LMC and SMC presented by Kim et al. (2003) and Stanimirovic et al. (1999)

(see also Staveley-Smith et al. 2003), respectively. These maps of the Magellanic clouds were made by combining interferometric (Australia Telescope Compact Array; ATCA) and single-dish (Parkes) observations. The map of the LMC has an angular resolution of 60", corresponding to a spatial resolution of 15 pc, while the SMC map has an angular resolution of 98", corresponding to a spatial scale of 29 pc. The rms noise in the LMC map is 2.4 K in a 1.65 km s<sup>-1</sup> channel width, and the value in the SMC map is 1.3 K in a 1.65 km s<sup>-1</sup> channel width.

To estimate the contribution to the [C II] emission from ionized gas, we used the Hα map from the Southern Alpha Sky Survey Atlas (SHASSA; Gaustad et al. 2001). The SHASSA maps cover the entire LMC and SMC with an angular resolution of 48" and a sensitivity level of 0.5 Rayleigh ($R = 10^6/4\pi$ photons cm<sup>-2</sup> s<sup>-1</sup> sr<sup>-1</sup>). We also used the NASA 70 m Deep Space Network telescope (DSS-43) to observe the H66α hydrogen radio recombination line at 22.364 GHz in one of our sources with the aim of testing the effect of dust extinction on the Hα observations. The angular resolution of the DSS-43 at 22.364 GHz is 48". We converted the data from an antenna temperature to a main-beam temperature scale using a main-beam efficiency of 0.50. The H66α spectrum has rms noise of 5.4 mK in a 3.5 km s<sup>-1</sup> channel width.

<footnote n="16"><sup>16</sup> This publication is based in part on data acquired with the Atacama Pathfinder Experiment (APEX). APEX is a collaboration between the Max-Planck-Institut für Radioastronomie, the European Southern Observatory, and the Onsala Space Observatory.</footnote>
<table><caption>**Table 3**
Integrated Intensities of Spectral Lines for the LMC Sample</caption><thead><tr><th>LOS</th><th>V<sub>LSR</sub><br>(km s<sup>-1</sup>)</th><th>I<sub>[H I]</sub>CNM<sup>a</sup></th><th>I<sub>[C II]</sub></th><th>I<sub>[C II]</sub>(1-0)</th><th>I<sub>[C II]</sub>(2-1)</th><th>I<sub>CO</sub>(1-0)<br>(K km s<sup>-1</sup>)</th><th>I<sub>CO</sub>(3-2)</th><th>I<sub>CO</sub>(7-6)</th><th>I<sub>13CO</sub>(1-0)</th><th>I<sub>13CO</sub>(3-2)</th></tr></thead><tbody><tr><td colspan="12" style="text-align: center;">LMC</td></tr><tr><td>Diff1_NW_1</td><td>288.8</td><td>510 ± 8</td><td>2.4 ± 0.3</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>Diff2_SE_1</td><td>233.3</td><td>169 ± 6</td><td>2.8 ± 0.6</td><td>0.9 ± 0.3</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>Diff3_RIDGE_1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>0.7 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>Diff5_SE_1</td><td>245.0</td><td>405 ± 7</td><td>0.9 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>Diff6_NW_1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>1.0 ± 0.1</td><td>0.3 ± 0.1</td><td>...</td><td>...</td><td>...</td></tr><tr><td>Diff7_NW_1</td><td>288.1</td><td>253 ± 6</td><td>2.7 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC10_NE_1</td><td>275.4</td><td>496 ± 7</td><td>1.8 ± 0.2</td><td>0.3 ± 0.1</td><td>...</td><td>1.7 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC11_Ridge_1</td><td>250.3</td><td>341 ± 7</td><td>7.9 ± 0.1</td><td>0.6 ± 0.1</td><td>0.4 ± 0.1</td><td>4.1 ± 0.2</td><td>3.4 ± 0.1</td><td>...</td><td>...</td><td>0.3 ± 0.1</td></tr><tr><td>LMC12_SE_1</td><td>265.7</td><td>224 ± 6</td><td>6.4 ± 0.2</td><td>0.1 ± 0.1</td><td>...</td><td>1.9 ± 0.2</td><td>2.8 ± 0.1</td><td>...</td><td>...</td><td>0.2 ± 0.1</td></tr><tr><td>LMC12_SE_2</td><td>270.6</td><td>288 ± 6</td><td>5.3 ± 0.2</td><td>1.1 ± 0.1</td><td>0.9 ± 0.1</td><td>4.7 ± 0.2</td><td>7.4 ± 0.1</td><td>...</td><td>0.4 ± 0.1</td><td>0.5 ± 0.1</td></tr><tr><td>LMC1_NW_1</td><td>293.1</td><td>288 ± 7</td><td>6.6 ± 0.0<sup>b</sup></td><td>0.6 ± 0.1</td><td>0.4 ± 0.0<sup>b</sup></td><td>1.8 ± 0.3</td><td>1.3 ± 0.2</td><td>1.5 ± 0.1</td><td>...</td><td>0.3 ± 0.0<sup>b</sup></td></tr><tr><td>LMC1_NW_2</td><td>287.6</td><td>341 ± 7</td><td>6.1 ± 0.0<sup>b</sup></td><td>0.9 ± 0.1</td><td>0.6 ± 0.1</td><td>2.3 ± 0.3</td><td>2.5 ± 0.3</td><td>0.4 ± 0.1</td><td>...</td><td>0.3 ± 0.1</td></tr><tr><td>LMC1_NW_3</td><td>288.0</td><td>890 ± 12</td><td>10.1 ± 0.0<sup>b</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC2_NW_1</td><td>290.5</td><td>716 ± 8</td><td>4.7 ± 0.4</td><td>0.5 ± 0.1</td><td>0.4 ± 0.1</td><td>2.3 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC3_NW_1</td><td>289.7</td><td>257 ± 6</td><td>1.6 ± 0.3</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC4_RIDGE_1</td><td>246.6</td><td>618 ± 9</td><td>5.7 ± 0.3</td><td>...</td><td>...</td><td>0.6 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC5_SE_1</td><td>219.0</td><td>75 ± 4</td><td>9.7 ± 0.1</td><td>0.7 ± 0.1</td><td>0.8 ± 0.2</td><td>2.5 ± 0.2</td><td>3.2 ± 0.1</td><td>...</td><td>...</td><td>0.4 ± 0.1</td></tr><tr><td>LMC7_RIDGE_1</td><td>232.7</td><td>745 ± 9</td><td>6.2 ± 0.5</td><td>...</td><td>...</td><td>1.9 ± 0.2</td><td>0.5 ± 0.2</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC8_RIDGE_1</td><td>238.5</td><td>688 ± 9</td><td>3.7 ± 0.4</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC9_NE_1</td><td>274.2</td><td>310 ± 7</td><td>2.0 ± 0.3</td><td>...</td><td>...</td><td>3.3 ± 0.3</td><td>1.3 ± 0.1</td><td>...</td><td>...</td><td>...</td></tr><tr><td>NT127_1</td><td>235.0</td><td>73 ± 4</td><td>0.6 ± 0.1</td><td>1.1 ± 0.1</td><td>0.2 ± 0.0<sup>b</sup></td><td>3.7 ± 0.2</td><td>4.2 ± 0.1</td><td>...</td><td>...</td><td>0.3 ± 0.1</td></tr><tr><td>NT2_NE_1</td><td>...</td><td>...</td><td>...</td><td>1.2 ± 0.2</td><td>...</td><td>8.2 ± 0.2</td><td>...</td><td>...</td><td>0.9 ± 0.1</td><td>...</td></tr><tr><td>NT74_1</td><td>236.1</td><td>310 ± 8</td><td>1.6 ± 0.4</td><td>2.6 ± 0.1</td><td>0.5 ± 0.1</td><td>17.1 ± 0.3</td><td>11.0 ± 0.1</td><td>...</td><td>2.5 ± 0.2</td><td>0.6 ± 0.1</td></tr><tr><td>NT77_1</td><td>217.4</td><td>59 ± 4</td><td>8.3 ± 0.1</td><td>1.0 ± 0.2</td><td>1.0 ± 0.1</td><td>8.1 ± 0.2</td><td>15.5 ± 0.1</td><td>1.1 ± 0.1</td><td>1.5 ± 0.1</td><td>2.7 ± 0.1</td></tr><tr><td>NT97_1</td><td>...</td><td>...</td><td>...</td><td>0.4 ± 0.1</td><td>...</td><td>3.8 ± 0.1</td><td>...</td><td>...</td><td>0.7 ± 0.1</td><td>...</td></tr><tr><td>NT99_1</td><td>227.0</td><td>270 ± 7</td><td>1.3 ± 0.3</td><td>...</td><td>...</td><td>5.0 ± 0.1</td><td>...</td><td>...</td><td>0.6 ± 0.1</td><td>...</td></tr><tr><td>PDR1_NW_1</td><td>285.8</td><td>775 ± 9</td><td>13.7 ± 0.3</td><td>3.1 ± 0.2</td><td>1.7 ± 0.1</td><td>13.6 ± 0.2</td><td>23.0 ± 0.1</td><td>...</td><td>1.6 ± 0.1</td><td>1.6 ± 0.1</td></tr><tr><td>PDR1_NW_2</td><td>293.3</td><td>287 ± 6</td><td>1.0 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>PDR2_NW_1</td><td>291.4</td><td>184 ± 7</td><td>3.1 ± 0.2</td><td>...</td><td>0.4 ± 0.0<sup>b</sup></td><td>2.1 ± 0.5</td><td>3.9 ± 0.1</td><td>1.5 ± 0.1</td><td>...</td><td>...</td></tr><tr><td>PDR2_NW_2</td><td>284.3</td><td>144 ± 7</td><td>47.0 ± 0.3</td><td>0.8 ± 0.1</td><td>1.8 ± 0.1</td><td>7.1 ± 0.6</td><td>11.8 ± 0.1</td><td>5.9 ± 0.2</td><td>...</td><td>1.0 ± 0.2</td></tr><tr><td>PDR2_NW_3</td><td>279.6</td><td>203 ± 8</td><td>21.0 ± 0.2</td><td>1.2 ± 0.1</td><td>0.5 ± 0.0<sup>b</sup></td><td>4.2 ± 0.3</td><td>2.8 ± 0.1</td><td>...</td><td>...</td><td>...</td></tr><tr><td>PDR2_NW_4</td><td>301.4</td><td>173 ± 11</td><td>13.6 ± 0.5</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>PDR3_NE_1</td><td>275.8</td><td>737 ± 9</td><td>42.7 ± 1.4</td><td>0.9 ± 0.4</td><td>0.4 ± 0.1</td><td>3.7 ± 0.5</td><td>14.7 ± 0.3</td><td>5.0 ± 0.2</td><td>0.3 ± 0.3</td><td>2.6 ± 0.1</td></tr><tr><td>PDR3_NE_2</td><td>271.4</td><td>391 ± 7</td><td>34.1 ± 0.2</td><td>2.5 ± 0.4</td><td>4.0 ± 0.1</td><td>15.1 ± 0.5</td><td>32.0 ± 0.3</td><td>7.6 ± 0.2</td><td>2.1 ± 0.3</td><td>5.3 ± 0.0<sup>b</sup></td></tr><tr><td>PDR3_NE_3</td><td>276.6</td><td>1370 ± 13</td><td>36.6 ± 1.4</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>PDR4_RIDGE_1</td><td>228.0</td><td>602 ± 9</td><td>9.3 ± 0.1</td><td>4.2 ± 0.2</td><td>2.2 ± 0.1</td><td>33.9 ± 0.3</td><td>34.2 ± 0.1</td><td>2.4 ± 0.1</td><td>5.4 ± 0.2</td><td>4.0 ± 0.1</td></tr><tr><td>SK-66D35_1</td><td>...</td><td>...</td><td>...</td><td>1.0 ± 0.3</td><td>0.3 ± 0.1</td><td>0.7 ± 0.1</td><td>2.0 ± 0.1</td><td>0.5 ± 0.1</td><td>...</td><td>0.3 ± 0.0<sup>b</sup></td></tr><tr><td>SK-66D35_2</td><td>279.5</td><td>584 ± 9</td><td>5.9 ± 0.4</td><td>...</td><td>...</td><td>0.7 ± 0.1</td><td>0.3 ± 0.0<sup>b</sup></td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-67D2_1</td><td>259.0</td><td>708 ± 10</td><td>2.2 ± 0.5</td><td>...</td><td>...</td><td>0.6 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-68D129_1</td><td>250.0</td><td>635 ± 9</td><td>7.0 ± 0.6</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-68D140_1</td><td>273.9</td><td>809 ± 14</td><td>6.2 ± 1.5</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-68D140_2</td><td>262.0</td><td>590 ± 9</td><td>5.4 ± 1.2</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-68D155_1</td><td>246.1</td><td>383 ± 8</td><td>8.0 ± 0.5</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-68D155_2</td><td>256.6</td><td>374 ± 7</td><td>1.8 ± 0.4</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-68D26_1</td><td>251.3</td><td>476 ± 9</td><td>2.2 ± 0.5</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-69D228_1</td><td>247.0</td><td>314 ± 9</td><td>6.1 ± 0.3</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr></tbody></table>

**Notes.**

We also compared our *Herschel* spectral line observations with dust continuum maps of the LMC and SMC taken using the PACS and SPIRE instruments on *Herschel*. These maps are part of the HERITAGE survey (Meixner et al. 2013), which provides images at 100 and 160 µm (PACS) and at 250, 350, and 500 µm (SPIRE). The angular resolutions of these maps are 7"/7 (100 µm), 12" (160 µm), 18" (250 µm), 25" (350 µm), and 40" (500 µm). The foreground dust emission originating from

the Milky Way was subtracted using a linear baseline (Meixner et al. 2013). We also used the *Spitzer* 24 µm maps of the LMC and SMC (Meixner et al. 2006). These maps have an angular resolution of 6".

## 2.5. Angular Resolution

The analysis presented in this paper uses a multiwavelength data set from different telescopes. Therefore, the angular

<footnote><sup>a</sup> H I intensity integrated over a velocity range defined by the full width at half maximum (FWHM) of the observed [C II] line (Section 3.2).</footnote>

<footnote><sup>b</sup> The error is below 0.05.</footnote>
<table><caption>**Table 4**
Integrated Intensities of Spectral Lines for the SMC Sample</caption><thead><tr><th rowspan="2">LOS</th><th rowspan="2">V<sub>LSR</sub><br>(km s<sup>-1</sup>)</th><th colspan="8">SMC</th></tr><tr><th>I<sub>[H I];CNM</sub><sup>a</sup></th><th>I<sub>[C II]</sub></th><th>I<sub>[C I](1-0)</sub></th><th>I<sub>[C I](2-1)</sub></th><th>I<sub>CO(1-0)</sub><br>(K km s<sup>-1</sup>)</th><th>I<sub>CO(3-2)</sub></th><th>I<sub>CO(7-6)</sub></th><th>I<sup>13</sup>CO(1-0)</th><th>I<sup>13</sup>CO(3-2)</th></tr></thead><tbody><tr><td>AzV18_1</td><td>122.0</td><td>1709 ± 11</td><td>5.9 ± 0.6</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>AzV18_2</td><td>136.7</td><td>764 ± 8</td><td>3.4 ± 0.5</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_B2_6_1</td><td>120.9</td><td>731 ± 4</td><td>15.6 ± 0.4</td><td>0.9 ± 0.1</td><td>0.8 ± 0.1</td><td>5.5 ± 0.2</td><td>5.6 ± 0.1</td><td>0.8 ± 0.1</td><td>0.8 ± 0.1</td><td>0.7 ± 0.1</td></tr><tr><td>SMC_B2_6_2</td><td>125.3</td><td>1273 ± 5</td><td>5.3 ± 0.4</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_HI_2_1</td><td>...</td><td>...</td><td>...</td><td>0.7 ± 0.2</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_HI_3_1</td><td>145.4</td><td>2347 ± 8</td><td>4.4 ± 0.8</td><td>0.6 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>0.4 ± 0.1</td><td>...</td><td>...</td></tr><tr><td>SMC_HI_4_1</td><td>123.3</td><td>1103 ± 4</td><td>5.5 ± 0.4</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_HI_4_2</td><td>114.8</td><td>508 ± 3</td><td>2.3 ± 0.3</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_HI_6_1</td><td>149.7</td><td>431 ± 4</td><td>1.3 ± 0.3</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_LIRS36_1</td><td>126.4</td><td>383 ± 3</td><td>16.4 ± 0.1</td><td>2.0 ± 0.1</td><td>1.8 ± 0.1</td><td>7.6 ± 0.2</td><td>18.9 ± 0.1</td><td>1.5 ± 0.1</td><td>0.7 ± 0.1</td><td>2.8 ± 0.2</td></tr><tr><td>SMC_LIRS36_2</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>0.5 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_LIRS49_1</td><td>110.4</td><td>448 ± 3</td><td>5.6 ± 0.0<sup>b</sup></td><td>...</td><td>...</td><td>...</td><td>0.8 ± 0.1</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_LIRS49_2</td><td>115.0</td><td>502 ± 3</td><td>14.3 ± 0.0<sup>b</sup></td><td>2.5 ± 0.2</td><td>1.3 ± 0.1</td><td>8.6 ± 0.2</td><td>18.1 ± 0.1</td><td>0.6 ± 0.1</td><td>0.6 ± 0.1</td><td>2.0 ± 0.2</td></tr><tr><td>SMC_LIRS49_3</td><td>125.9</td><td>1173 ± 5</td><td>5.2 ± 0.0<sup>b</sup></td><td>...</td><td>...</td><td>1.3 ± 0.2</td><td>2.4 ± 0.1</td><td>...</td><td>...</td><td>0.3 ± 0.1</td></tr><tr><td>SMC_LIRS49_4</td><td>137.2</td><td>191 ± 2</td><td>1.4 ± 0.0<sup>b</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_NE_1a_1</td><td>148.9</td><td>496 ± 4</td><td>4.5 ± 0.1</td><td>0.5 ± 0.1</td><td>0.5 ± 0.1</td><td>4.0 ± 0.1</td><td>4.3 ± 0.1</td><td>0.6 ± 0.1</td><td>0.4 ± 0.0<sup>b</sup></td><td>...</td></tr><tr><td>SMC_NE_3c_1</td><td>175.1</td><td>170 ± 2</td><td>1.1 ± 0.1</td><td>1.2 ± 0.1</td><td>0.7 ± 0.1</td><td>4.2 ± 0.2</td><td>4.1 ± 0.1</td><td>...</td><td>0.7 ± 0.0<sup>b</sup></td><td>0.4 ± 0.1</td></tr><tr><td>SMC_NE_3g_1</td><td>169.1</td><td>690 ± 4</td><td>4.0 ± 0.1</td><td>0.6 ± 0.2</td><td>0.6 ± 0.1</td><td>4.3 ± 0.1</td><td>3.2 ± 0.1</td><td>0.5 ± 0.1</td><td>0.4 ± 0.0<sup>b</sup></td><td>0.5 ± 0.1</td></tr><tr><td>SMC_NE_4a_hi_1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>0.7 ± 0.1</td><td>...</td><td>...</td><td>0.2 ± 0.1</td></tr><tr><td>SMC_NE_4a_hi_2</td><td>153.9</td><td>131 ± 2</td><td>1.6 ± 0.2</td><td>0.5 ± 0.1</td><td>...</td><td>3.1 ± 0.1</td><td>2.1 ± 0.1</td><td>...</td><td>0.2 ± 0.0<sup>b</sup></td><td>...</td></tr><tr><td>SMC_NE_4a_hi_3</td><td>122.2</td><td>551 ± 4</td><td>1.3 ± 0.3</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_NE_4c_low_1</td><td>121.6</td><td>780 ± 4</td><td>2.1 ± 0.4</td><td>0.7 ± 0.1</td><td>...</td><td>5.8 ± 0.1</td><td>1.7 ± 0.1</td><td>0.2 ± 0.1</td><td>0.7 ± 0.0<sup>b</sup></td><td>...</td></tr></tbody></table>

**Notes.**

resolution of the observations is not uniform. While maps can be convolved to a uniform resolution, this is not the case for pointed [C II], [C I], and CO observations, which have angular resolutions ranging from 12" for [C II] to 60" for H I. Therefore, the absolute intensities and line ratios used in our analysis can be affected by beam dilution effects.

We studied the effect of beam dilution in our observations by smoothing the 12" angular resolution 160 µm HERITAGE dust continuum map, which is assumed to be a proxy for the distribution of gas in the LMC and SMC. We then compared the resulting intensities at different angular resolutions. In the case of ionized gas traced by Hα observations, we used the *Spitzer* 24 µm images of the Magellanic clouds, which are a proxy for hot dust emission associated with ionized gas. Details of this study are presented in the Appendix. The effects of beam dilution on our results will be discussed as they are presented throughout the paper. In general, we found that beam dilution has a minor effect on our results for most of our sample, with the exception of a handful of point-like, unresolved PDR sources. Note that the data used in our analysis are not corrected for beam dilution effects, due to the uncertainty of whether the 24 µm or 160 µm are good tracers of the ionized and neutral gas components, respectively.

### 3. The Origin of the [C II] Emission in the Magellanic Clouds

The observed [C II] emission arises from gas associated with hydrogen in the form of H$^0$, H$^2$, and H$^+$. In the latter gas component, the collisions are mainly with electrons due to their

higher speeds, so we refer to it as e⁻ gas. Thus, the observed [C II] emission is given by

$$I_{[C II]} = I_{[C II]}^{H^0} + I_{[C II]}^{H_2} + I_{[C II]}^{e^{-}} \quad (1)$$

For optically thin emission, the [C II] intensity (in units of K km s⁻¹) is related to the C⁺ column density, $N_{C^+}$ (cm⁻²), and the volume density of the collisional partner (assumed to be uniform along the LOS), n (e⁻, H$^0$, or H$^2$; cm⁻³), as (see, e.g., Goldsmith et al. 2012)

$$I_{[C II]} = N_{C^+} \left[ 3.05 \times 10^{15} \left( 1 + 0.5 \left( 1 + \frac{A_{ul}}{R_{ul}n} \right) e^{\frac{91.21}{T_{kin}}} \right) \right]^{-1}, \quad (2)$$

where $A_{ul} = 2.3 \times 10^{-6}$ s⁻¹ is the Einstein spontaneous decay rate and $R_{ul}$ is the collisional de-excitation rate coefficient at a kinetic temperature $T_{kin}$, with u and l denoting the upper and lower energy levels. Values of $R_{ul}$ for collisions with H$^0$, H$^2$, and e⁻ as a function of the kinetic temperature are available from Barinovs et al. (2005), Wiesenfeld & Goldsmith (2014), and Wilson & Bell (2002), respectively. For example, at $T_{kin} = 100$ K, $R_{ul} = 7.58 \times 10^{-10}$ and $5.12 \times 10^{-10}$ cm³ s⁻¹ for collisions with atomic and molecular hydrogen,<fnref n="17" /> respectively, while at $T_{kin} = 8000$ K, $R_{ul} = 5.2 \times 10^{-8}$ cm³ s⁻¹ for collisions with electrons.

Dufour et al. (1982) studied the carbon abundance in H II regions in the Magellanic clouds, deriving $12 + \log(C/H) = 7.9$

<footnote>a H I intensity integrated over a velocity range defined by the FWHM of the observed [C II] line (Section 3.2).</footnote>

<footnote>b The error is below 0.05.</footnote>

<footnote n="17">¹⁷ Assuming an ortho-to-para H₂ ratio, OPR = 1. See discussion in Wiesenfeld & Goldsmith (2014).</footnote>
<figure><img src="image_6.png" /><figcaption>**Figure 3.** Ratio of the TIR intensity (Section 5.4) to the H$^0$ column density (Section 3.2) as a function of the visual extinction (Section 4.1) for the entire LMC and SMC. We also include data points that correspond to those in our sample in the LMC and SMC. The grayscale represents the number of pixels at a given TIR/N(H$^0$) and A$_V$ bin, with contour lines representing 2, 10, 100, and 1000 pixels.</figcaption></figure>

<figure><img src="image_7.png" /></figure>

for the LMC and 12+log(C/H) = 7.16 for the SMC. Relative to the gas-phase carbon fractional abundance of [C/H] = 1.4 × 10⁻⁴ determined in the Milky Way (Sofia et al. 1997), the carbon abundances in the LMC and SMC are 1.8 and 10 times lower, respectively. We used the fractional carbon abundances derived by Dufour et al. (1982) in the analysis presented here.

In the following, we studied the relative contribution to the observed [C II] emission from ionized, atomic, and molecular gas. In Figure 4, we show an example of the decomposition of the [C II] emission originating from different ISM components. The derived [C II] emission fractions as a function of the observed [C II] intensity for the entire sample are summarized in Figure 5.

### 3.1. Ionized Gas

We estimated the [C II] intensity originating from ionized gas (e⁻ gas) using the SHASSA (Gaustad et al. 2001) Hα emission maps of the LMC and SMC. A commonly defined quantity used to relate the properties of the ionized gas and the observed Hα emission is the emission measure (EM), defined as the integral of the electron volume density squared along the LOS,

$$EM = \int n_e^2 dl. \qquad (3)$$

Assuming that the electron density is constant along the LOS, this equation can be simplified to

$$EM = \langle n_e \rangle N_e \approx \langle n_e \rangle N_{H^+}. \qquad (4)$$

The EM is related to the intensity of the Hα line in units of Rayleigh (R) as (e.g., Reynolds 1991)

$$\left( \frac{\text{EM}}{\text{pc cm}^{-6}} \right) = 2.75 \left( \frac{T_{\text{kin}}}{10^4 \text{ K}} \right)^{0.9} \left( \frac{I_{\text{H}\alpha}}{\text{R}} \right). \qquad (5)$$

The critical density, $n_{cr} \equiv A_{ul}/R_{ul}$, is the density at which the collisional de-excitation rate is equal to the effective spontaneous decay rate. For collisions of C⁺ with electrons at

<figure><img src="image_8.png" /><figcaption>**Figure 4.** Sample [C II] and [C I] $^3P_2-^3P_1$ spectrum illustrating the decomposition between [C II] emission associated with PDRs, CO-dark H$_2$ gas, cold H I, and ionized gas. See Section 3 for details on the decomposition of the [C II] emission.</figcaption></figure>

$T_{\text{kin}} = 8000$ K, this is equal to $44 \text{ cm}^{-3}$. For electron densities much smaller than the critical density $n_e \ll n_{cr}$, Equation (2) can be written as

$$I_{[C \text{ II}]}^{-} = n_e N_{C^{+}} [1.52 \times 10^{15} n_{cr} e^{\frac{91.21}{T_{\text{kin}}}}]^{-1}. \qquad (6)$$

The term $n_e N_{C^+}$ can be written in terms of the fractional abundance of ionized carbon, $X_{C^+} = N_{C^+}/N_{H^+}$, as $X_{C^+} n_e N_{H^+} = X_{C^+} \text{EM}$. Paradis et al. (2011) estimated the typical electron
<figure><img src="image_9.png" /><figcaption>Figure 5. Fraction of the [C II] emission that we estimated arises from ionized gas, cold atomic gas, CO-dark H₂, and PDRs as a function of the observed [C II] emission in the LMC and SMC.</figcaption></figure>

densities for regions emitting different regimes of Hα emission. For the range of Hα intensities in our sample, electron densities vary from 0.05 to 3.98 cm⁻³, thus validating the assumption that n_e ≪ n_cr. Equation (6) can thus be rewritten in terms of the EM as

$$I_{[C \text{II}]}^{-} = X_{[C^{+}]} \text{EM}[1.52 \times 10^{15} n_{\text{cr}} e^{\frac{91.21}{T_{\text{kin}}}}]^{-1} \quad (7)$$

and can be used to estimate the contribution of ionized gas to the observed [C II] emission from Hα observations.

In Figure 5, we show the fraction of the [C II] emission that we have estimated arises from ionized gas as a function of the observed [C II] emission in the LMC and SMC. We assumed T_kin = 8000 K for the kinetic temperature of the ionized gas. In LOSs with multiple velocity components, we assumed that each component contributes equally to the derived [C II] emission from ionized gas, as suggested by the H66α spectrum in Figure 6 that shows two [C II] velocity components with varying peak intensities but uniform recombination line emission. We found that ionized gas tends to contribute a small fraction of the [C II] emission, with typical fractions around 19% in the LMC and 5% in the SMC. These contributions from ionized gas to the observed [C II] emission are in agreement with those estimated using the unobscured [N II] fine structure lines by Chevance et al. (2016) and Okada et al. (2015) in the 30 Doradus and N159 regions in the LMC, respectively, and by Requena-Torres et al. (2016) in several star-forming regions in the SMC. The derived contributions from ionized gas to the observed [C II] emission in the LMC and SMC are also consistent with those estimated in the Galactic plane (Pineda et al. 2013).

We estimated the H⁺ column density in our LOSs using EM = n_e N_H⁺ and the volume densities of the ionized gas suggested by Paradis et al. (2011) for different ranges of Hα emission. Typical H⁺ column densities in our sample are 10²⁰.⁸±⁰.³⁵ cm⁻² for the LMC and 10²⁰.⁵±⁰.¹⁴ cm⁻² for the SMC.

Note that the Hα emission used here might be affected by extinction from dust grains, making our estimate of the EM

<figure><img src="image_10.png" /><figcaption>Figure 6. Comparison between [C II] and the H66α radio recombination line in the PDR2_NW position.</figcaption></figure>

from Hα observations a lower limit. We tested the effect of dust extinction in the Hα observations by assuming that half of the visual extinction (A_Hα = 0.81A_V; e.g., Viallefond & Goss 1986; Parker et al. 1992) derived for each LOS (Section 4.1) is in front of the Hα sources. We found that the Hα intensities would be underestimated by an average factor of 4.6 in the LMC and 1.4 in the SMC. In the LMC, the large factor is dominated by three warm and dense PDRs—PDR4_RIDGE, PDR3_NW, and LMC_1_NW—that are bright in [C II] and have the largest visual extinctions. Without these PDR regions, we found that the Hα emission would be affected by a factor of 1.65 in the LMC. If a correction to the Hα
intensity for dust extinction is applied to the calculation of the
contribution of ionized gas to the observed [C II] emission,
the fraction would increase by a similar factor. Note that for the
warm and dense PDRs mentioned above, applying a correction
to the contribution from ionized gas to the observed [C II]
emission would result in larger [C II] intensities than are ob-
served, suggesting that less than $A_V/2$ of extinction is in the
foreground of these sources.

We tested the reliability of Hα as a tracer of ionized gas by
observing the unobscured H66α radio recombination line using
the DSS-43 NASA Deep Space Network telescope in one of
our LOSs (Figure 6). Following Alves et al. (2015), we
converted the integrated intensity of the H66α line to the EM,
assuming T<sub>kin</sub> = 8000 K, as 1.9 × 10<sup>4</sup> cm<sup>-6</sup>pc. For the same
position, we obtained EM = 7.8 × 10<sup>3</sup> cm<sup>-6</sup>pc from Hα, a
factor of ~2.4 lower than that obtained from H66α. The
position we observed in H66α corresponds to a warm and
dense PDR that is likely associated with a large column of dust,
so we expected a smaller effect in more diffuse regions in our
sample.

3.2. Atomic Gas

We estimated the contribution from atomic gas to the
observed [C II] emission in our sample in the Magellanic clouds
using the H I 21 cm line observations. We calculated N (H�0) for
each LOS, assuming optically thin emission, using
N(H⁰) = 1.82 × 10¹⁸ I(H I) cm⁻², with I(H I) in units of
K km s⁻¹. We converted from N(H⁰) to N(C⁺) using the
carbon fractional abundances of the LMC and SMC, assuming
that all gas-phase carbon associated with atomic gas is in the
form of C⁺.

We only considered H I emission that is associated in
velocity space with [C II] emission. To calculate I (H I), we
integrated the H I spectrum over a velocity range defined by the
full width at half maximum (FWHM) of the [C II] lines.
Because the [C II] emission is volume density-sensitive for
nH < 3000 cm⁻³, the critical density for collisions with H
(Goldsmith et al. 2012), we expected that the [C II] emission
was associated with cold, dense H I (CNM; Wolfire et al. 2003)
rather than diffuse, warm H I (WNM; see discussion in Pineda
et al. 2013). We therefore considered the H I emission that is
associated in velocity with [C II] to be CNM, while the
remaining H I emission is WNM. In the top panels of Figure 2,
we illustrate this decomposition by showing the velocity range
as a shaded region in which we assign the H I to be CNM. Note
that, due to the typical complexity of the H I line profiles, it is
difficult to reliably separate components in velocity for our
entire data set.

Our method for calculating the H I column could be
overestimating the column density of cold atomic hydrogen
associated with [C II] if a broad WNM component overlaps in
velocity with the [C II] emission. To quantify this effect, we
attempted to fit multiple Gaussian components to the H I
spectra of SMC_HI_4 and SMC_LIRS36 shown in Figure 2.
For the Gaussian fitting, we fixed the Local Standard of Rest
(LSR) velocity of the [C II] components (see Table 4) and left
the peak main-beam temperature and line width as free
parameters. We then subtracted all Gaussian components that
were not associated with [C II] from the observed spectra in
order to isolate the Gaussian component associated with [C II].
In this component, we recalculated the integrated intensity by

integrating over the velocity range defined by the FWHM of
the [C II] line, as above. We found that we would typically
overestimate the H I intensity associated with CNM by a factor
of about 2 in SMC_HI_4 and SMC_LIRS36. Note, however,
that complex H I spectra, such as those of SMC_HI_4 and
SMC_LIRS36, are typically rare in our sample and mostly
correspond to SMC LOSs. In the SMC, there are five LOSs
with such complexity out of 10 detected in [C II]. In the LMC,
the H I spectra are mostly similar to those of Diff1_NW and
PDR1_NW, as shown in Figure 2.

To calculate the [C II] intensity from atomic gas using
Equation (2), we needed to assume a gas volume density and a
kinetic temperature. As discussed in Section 3.4, we adopted a
kinetic temperature of T<sub>kin</sub> = 70 K in the case of velocity
components in which H I and [C II] are detected but <sup>13</sup>CO
J = 1 → 0 and/or J = 3 → 2 are not. In these LOSs, we were
unable to determine the physical conditions of the C<sup>0</sup> or CO
layer. In LOSs with kinetic temperatures of the C<sup>0</sup> and/or CO
layers derived in the excitation analysis (Section 4.4.3) and
with [C II] and H I detected, we still assumed 70 K if the kinetic
temperature of the C<sup>0</sup> layer was lower than 70 K. Otherwise,
we assumed the kinetic temperature of the C<sup>0</sup> layer. The
volume density was then derived by combining the kinetic
temperatures with the thermal pressures of the diffuse ISM
derived in Section 4.3.

In Figure 5, we show the fraction of the observed [C II]
emission from atomic gas as a function of the observed [C II].
The contribution from atomic gas to the observed [C II]
emission shows a large scatter with average values of about
43% for the LMC and 28% for the SMC. These average
contributions are consistent with the 20% contribution derived
in the plane of the Milky Way by Pineda et al. (2013).

3.3. CO-Dark H₂ Gas and Dense PDRs

When we subtracted the contribution from the ionized and
atomic gas components from the observed total [C II] emission,
we typically found that there was residual [C II] emission. This
residual emission is produced in regions where carbon is
ionized but hydrogen is molecular. The molecular hydrogen
component in which carbon is mainly ionized so that it is
traced by [C II] and not by CO is what we refer to as CO-dark
H₂ gas. The CO-dark H₂ gas can originate from (1) a diffuse
cloud that has enough dust shielding to maintain only a trace
amount of C⁰ and CO, thus with most of the carbon in the form
of C⁺, or (2) the envelopes of a warm and dense molecular
cloud in which the inner parts have significant column densities
of C⁰ and CO. In order to associate the observed [C II] velocity
components with either of these gas conditions, we used the
detection limit of the ¹³CO lines observed in our survey to
identify clouds that have enough shielding to maintain
significant amounts of CO and therefore are likely associated
warm and dense molecular clouds with ongoing star formation.
We assigned the [C II] emission to warm and dense PDRs in
velocity components with enough CO column density for the
¹³CO J = 1 → 0 and/or ¹³CO J = 3 → 2 lines to be detected.
In LOSs where the ¹³CO lines are detected, we were able to
perform an excitation analysis (Section 4.4) that confirms our
assumption of warm and dense gas in these velocity
components (see Table 5). In LOSs that are detected in
[C II] and H I but not in ¹³CO J = 1 → 0 and/or ¹³CO
J = 3 → 2, we assumed that the residual [C II] emission arises
<table><caption>**Table 5**

Derived Parameters for the LMC and SMC Sample</caption><thead><tr><th>LOS</th><th>N (H<sup>0</sup>)<sub>CNM</sub><sup>a</sup><br>log<br>(cm<sup>-2</sup>)</th><th>T<sub>kin</sub><sup>C+,b</sup><br>(K)</th><th>n<sub>H2</sub><sup>C+,c</sup><br>log<br>(cm<sup>-3</sup>)</th><th>N (C<sup>+</sup>)<sub>H2</sub><sup>d</sup><br>log<br>(cm<sup>-2</sup>)</th><th>T<sub>kin</sub><sup>C0,b</sup><br>(K)</th><th>n<sub>H2</sub><sup>C0,c</sup><br>log<br>(cm<sup>-3</sup>)</th><th>N (C<sup>0</sup>)<sup>d</sup><br>log<br>(cm<sup>-2</sup>)</th><th>T<sub>kin</sub><sup>CO,b</sup><br>(K)</th><th>n<sub>H2</sub><sup>CO,c</sup><br>log<br>(cm<sup>-3</sup>)</th><th>N (CO)<sup>d</sup><br>log<br>(cm<sup>-2</sup>)</th><th>Filling<sup>e</sup><br>Factor</th></tr></thead><tbody><tr><th colspan="13" style="text-align: center;">LMC</th></tr><tr><td>Diff2_SE_1</td><td>20.5<sup>f</sup></td><td>49</td><td>2.8</td><td>17.1<sup>+0.2</sup><sub>-0.1</sub></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>Diff7_NW_1</td><td>20.7<sup>f</sup></td><td>49</td><td>2.8</td><td>17.1 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC11_Ridge_1</td><td>20.8<sup>f</sup></td><td>40 ± 10</td><td>3.4<sup>+0.0f</sup><sub>-0.3</sub></td><td>17.5 ± 0.3</td><td>40 ± 10</td><td>3.9<sup>+0.2</sup><sub>-0.6</sub></td><td>15.9<sup>+0.3</sup><sub>-0.7</sub></td><td>20 ± 10</td><td>4.5 ± 0.1</td><td>15.9<sup>+0.3</sup><sub>-0.7</sub></td><td>0.10<sup>+0.04</sup><sub>-0.32</sub></td><td></td></tr><tr><td>LMC12_SE_1</td><td>20.6<sup>f</sup></td><td>30 ± 10</td><td>3.7 ± 0.0<sup>f</sup></td><td>17.5<sup>+0.3</sup><sub>-0.6</sub></td><td>30 ± 10</td><td>4.4 ± 0.1</td><td>15.3<sup>+0.2</sup><sub>-0.3</sub></td><td>40 ± 10</td><td>4.4 ± 0.1</td><td>16.0<sup>+0.2</sup><sub>-0.3</sub></td><td>0.10<sup>+0.02</sup><sub>-0.05</sub></td><td></td></tr><tr><td>LMC12_SE_2</td><td>20.7<sup>f</sup></td><td>60 ± 10</td><td>3.1 ± 0.0<sup>f</sup></td><td>17.1<sup>+0.1</sup><sub>-0.2</sub></td><td>60 ± 10</td><td>3.5 ± 0.1</td><td>16.2<sup>+0.2</sup><sub>-0.4</sub></td><td>30 ± 10</td><td>4.6<sup>+0.3</sup><sub>-0.1</sub></td><td>16.3<sup>+0.2</sup><sub>-0.4</sub></td><td>0.34<sup>+0.09</sup><sub>-0.32</sub></td><td></td></tr><tr><td>LMC1_NW_1</td><td>20.7<sup>f</sup></td><td>120 ± 20</td><td>2.6 ± 0.0<sup>f</sup></td><td>17.2 ± 0.1</td><td>120 ± 20</td><td>2.7 ± 0.1</td><td>16.3<sup>+0.2</sup><sub>-0.4</sub></td><td>50<sup>+10</sup><sub>-20</sub></td><td>6.0<sup>+1.9</sup><sub>-0.1</sub></td><td>16.2<sup>+0.4</sup><sub>-0.5</sub></td><td>0.01<sup>+0.00g</sup><sub>-0.02</sub></td><td></td></tr><tr><td>LMC1_NW_2</td><td>20.8<sup>f</sup></td><td>160 ± 10</td><td>2.6 ± 0.0<sup>f</sup></td><td>16.9 ± 0.1</td><td>160 ± 10</td><td>2.9 ± 0.1</td><td>16.1<sup>+0.2</sup><sub>-0.3</sub></td><td>60 ± 10</td><td>4.0 ± 0.1</td><td>15.8<sup>+0.2</sup><sub>-0.3</sub></td><td>0.02<sup>+0.00g</sup><sub>-0.01</sub></td><td></td></tr><tr><td>LMC1_NW_3</td><td>21.2<sup>f</sup></td><td>49</td><td>2.8</td><td>17.6 ± 0.0<sup>f</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC2_NW_1</td><td>21.1<sup>f</sup></td><td>49</td><td>2.8</td><td>17.2 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC3_NW_1</td><td>20.7<sup>f</sup></td><td>49</td><td>2.8</td><td>16.7<sup>+0.3</sup><sub>-0.2</sub></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC4_RIDGE_1</td><td>21.1<sup>f</sup></td><td>49</td><td>2.8</td><td>17.4 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC5_SE_1</td><td>20.1<sup>f</sup></td><td>170<sup>+10</sup><sub>-20</sub></td><td>2.8</td><td>17.3 ± 0.0<sup>f</sup></td><td>170<sup>+10</sup><sub>-20</sub></td><td>3.3 ± 0.1</td><td>16.0<sup>+0.2</sup><sub>-0.4</sub></td><td>30 ± 10</td><td>4.6<sup>+0.2</sup><sub>-0.1</sub></td><td>15.9<sup>+0.2</sup><sub>-0.4</sub></td><td>0.09<sup>+0.02</sup><sub>-0.07</sub></td><td></td></tr><tr><td>LMC7_RIDGE_1</td><td>21.1<sup>f</sup></td><td>49</td><td>2.8</td><td>17.4 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>LMC9_NE_1</td><td>20.8<sup>f</sup></td><td>49</td><td>2.8</td><td>16.7<sup>+0.3</sup><sub>-0.2</sub></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>NT127_1</td><td>20.1<sup>f</sup></td><td>30 ± 10</td><td>2.6 ± 0.0<sup>f</sup></td><td>17.2<sup>+0.5</sup><sub>-0.7</sub></td><td>30 ± 10</td><td>2.1 ± 0.1</td><td>16.7<sup>+0.3</sup><sub>-0.4</sub></td><td>30 ± 10</td><td>4.3 ± 0.1</td><td>16.0<sup>+0.3</sup><sub>-0.4</sub></td><td>0.16<sup>+0.05</sup><sub>-0.15</sub></td><td></td></tr><tr><td>NT77_1</td><td>20.0<sup>f</sup></td><td>190<sup>+20</sup><sub>-10</sub></td><td>2.7 ± 0.0<sup>f</sup></td><td>17.3<sup>+0.0f</sup><sub>-0.1</sub></td><td>190<sup>+20</sup><sub>-10</sub></td><td>3.2 ± 0.1</td><td>16.1<sup>+0.6</sup><sub>-0.2</sub></td><td>50 ± 20</td><td>4.3<sup>+0.6</sup><sub>-0.4</sub></td><td>16.9<sup>+1.1</sup><sub>-0.6</sub></td><td>0.74<sup>+0.49</sup><sub>-0.26</sub></td><td></td></tr><tr><td>PDR1_NW_1</td><td>21.1<sup>f</sup></td><td>30 ± 10</td><td>3.7 ± 0.0<sup>f</sup></td><td>17.8<sup>+0.3</sup><sub>-0.7</sub></td><td>30 ± 10</td><td>4.4<sup>+0.2</sup><sub>-0.1</sub></td><td>16.6<sup>+0.3</sup><sub>-0.2</sub></td><td>30 ± 10</td><td>4.4<sup>+0.2</sup><sub>-0.1</sub></td><td>16.8<sup>+0.7</sup><sub>-0.2</sub></td><td>0.79<sup>+0.24</sup><sub>-0.21</sub></td><td></td></tr><tr><td>PDR2_NW_2</td><td>20.4<sup>f</sup></td><td>40 ± 10</td><td>4.4 ± -0.0</td><td>17.9<sup>+0.2</sup><sub>-0.3</sub></td><td>40 ± 10</td><td>5.9 ± 0.1</td><td>16.1 ± 0.2</td><td>40 ± 10</td><td>5.9 ± 0.1</td><td>16.5 ± 0.2</td><td>0.11<sup>+0.02</sup><sub>-0.03</sub></td><td></td></tr><tr><td>PDR2_NW_3</td><td>20.6<sup>f</sup></td><td>49</td><td>2.8</td><td>18.0 ± 0.0<sup>f</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>PDR2_NW_4</td><td>20.5<sup>f</sup></td><td>49</td><td>2.8</td><td>17.6 ± 0.0<sup>f</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>PDR3_NE_1</td><td>21.1<sup>f</sup></td><td>150 ± 10</td><td>2.3 ± 0.0<sup>f</sup></td><td>18.4 ± 0.1</td><td>150 ± 10</td><td>2.3 ± 0.1</td><td>16.2<sup>+0.2</sup><sub>-0.3</sub></td><td>60<sup>+10</sup><sub>-20</sub></td><td>4.2<sup>+0.2</sup><sub>-0.3</sub></td><td>16.8<sup>+0.2</sup><sub>-0.3</sub></td><td>0.09<sup>+0.02</sup><sub>-0.04</sub></td><td></td></tr><tr><td>PDR3_NE_2</td><td>20.9<sup>f</sup></td><td>170 ± 10</td><td>3.4 ± 0.0<sup>f</sup></td><td>17.5 ± 0.0<sup>f</sup></td><td>170 ± 10</td><td>4.5 ± 0.1</td><td>16.6<sup>+0.3</sup><sub>-0.3</sub></td><td>40 ± 10</td><td>4.5 ± 0.4</td><td>17.3<sup>+0.2</sup><sub>-0.3</sub></td><td>0.38<sup>+0.08</sup><sub>-0.28</sub></td><td></td></tr><tr><td>PDR3_NE_3</td><td>21.4<sup>f</sup></td><td>49</td><td>2.8</td><td>18.3 ± 0.0<sup>f</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>PDR4_RIDGE_1</td><td>21.0<sup>f</sup></td><td>80 ± 20</td><td>2.7 ± 0.0<sup>f</sup></td><td>17.5<sup>+0.1</sup><sub>-0.2</sub></td><td>80 ± 20</td><td>2.8 ± 0.1</td><td>16.7<sup>+0.3</sup><sub>-0.5</sub></td><td>30 ± 10</td><td>4.6<sup>+0.7</sup><sub>-0.4</sub></td><td>17.2<sup>+0.4</sup><sub>-0.6</sub></td><td>0.40<sup>+0.16</sup><sub>-0.60</sub></td><td></td></tr><tr><td>SK-66D35_2</td><td>21.0<sup>f</sup></td><td>49</td><td>2.8</td><td>17.1 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-68D140_1</td><td>21.1<sup>f</sup></td><td>49</td><td>2.8</td><td>17.2 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-68D155_1</td><td>20.8<sup>f</sup></td><td>49</td><td>2.8</td><td>17.6 ± 0.0<sup>f</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SK-69D228_1</td><td>20.8<sup>f</sup></td><td>49</td><td>2.8</td><td>17.1 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><th colspan="13" style="text-align: center;">SMC</th></tr><tr><td>AzV18_1</td><td>21.5<sup>f</sup></td><td>49</td><td>3.3</td><td>17.1 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>AzV18_2</td><td>21.1<sup>f</sup></td><td>49</td><td>3.3</td><td>16.9 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_B2_6_1</td><td>21.1<sup>f</sup></td><td>160 ± 10</td><td>2.9 ± 0.0<sup>f</sup></td><td>17.4 ± 0.1</td><td>160 ± 10</td><td>3.1 ± 0.1</td><td>16.1<sup>+0.3</sup><sub>-0.5</sub></td><td>...</td><td>4.1<sup>+0.3</sup><sub>-0.2</sub></td><td>16.4<sup>+0.4</sup><sub>-0.6</sub></td><td>0.06<sup>+0.03</sup><sub>-0.07</sub></td><td></td></tr><tr><td>SMC_B2_6_2</td><td>21.4<sup>f</sup></td><td>49</td><td>3.3</td><td>17.1 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_HI_4_1</td><td>21.3<sup>f</sup></td><td>49</td><td>3.3</td><td>17.1 ± 0.0<sup>f</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_HI_4_2</td><td>21.0<sup>f</sup></td><td>49</td><td>3.3</td><td>16.7 ± 0.1</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_LIRS36_1</td><td>20.8<sup>f</sup></td><td>60 ± 10</td><td>3.5 ± 0.0<sup>f</sup></td><td>17.5 ± 0.1</td><td>60 ± 10</td><td>3.8 ± 0.1</td><td>16.4<sup>+0.4</sup><sub>-0.1</sub></td><td>...</td><td>4.3<sup>+0.4</sup><sub>-0.2</sub></td><td>16.9<sup>+0.8</sup><sub>-0.3</sub></td><td>0.89<sup>+0.42</sup><sub>-0.11</sub></td><td></td></tr><tr><td>SMC_LIRS49_1</td><td>20.9<sup>f</sup></td><td>49</td><td>3.3</td><td>17.2 ± 0.0<sup>f</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_LIRS49_2</td><td>21.0<sup>f</sup></td><td>50 ± 10</td><td>3.1 ± 0.0<sup>f</sup></td><td>17.8<sup>+0.1</sup><sub>-0.2</sub></td><td>50 ± 10</td><td>3.0 ± 0.1</td><td>16.5 ± 0.4</td><td>...</td><td>4.4 ± 0.3</td><td>16.8<sup>+0.8</sup><sub>-0.6</sub></td><td>0.51<sup>+0.26</sup><sub>-0.49</sub></td><td></td></tr><tr><td>SMC_LIRS49_3</td><td>21.3<sup>f</sup></td><td>50 ± 10</td><td>3.8<sup>+0.1f</sup><sub>-0.0</sub></td><td>16.8<sup>+0.1</sup><sub>-0.2</sub></td><td>...</td><td>...</td><td>...</td><td>...</td><td>4.3<sup>+0.3</sup><sub>-0.1</sub></td><td>16.3<sup>+0.3</sup><sub>-0.5</sub></td><td>0.15<sup>+0.03</sup><sub>-0.24</sub></td><td></td></tr><tr><td>SMC_LIRS49_4</td><td>20.5<sup>f</sup></td><td>49</td><td>3.3</td><td>16.6 ± 0.0<sup>f</sup></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td></tr><tr><td>SMC_NE_1a_1</td><td>21.0<sup>f</sup></td><td>110 ± 10</td><td>3.3</td><td>16.7 ± 0.1</td><td>110 ± 10</td><td>3.6 ± 0.1</td><td>15.8<sup>+0.1</sup><sub>-0.2</sub></td><td>...</td><td>3.9<sup>+0.3</sup><sub>-0.2</sub></td><td>16.6<sup>+0.3</sup><sub>-0.2</sub></td><td>0.12<sup>+0.00g</sup><sub>-0.05</sub></td><td></td></tr></tbody></table>
<table><caption>**Table 5**
(Continued)</caption><thead><tr><th>LOS</th><th>N (H<sup>0</sup>)<sub>CNM</sub><sup>a</sup><br>log<br>(cm<sup>-2</sup>)</th><th>T<sub>kin</sub><sup>C+,b</sup><br>(K)</th><th>n<sub>H<sub>2</sub></sub><sup>C+,c</sup><br>log<br>(cm<sup>-3</sup>)</th><th>N (C<sup>+</sup>)<sub>H<sub>2</sub></sub><sup>d</sup><br>log<br>(cm<sup>-2</sup>)</th><th>T<sub>kin</sub><sup>C0,b</sup><br>(K)</th><th>n<sub>H<sub>2</sub></sub><sup>C0,c</sup><br>log<br>(cm<sup>-3</sup>)</th><th>N (C<sup>0</sup>)<sup>d</sup><br>log<br>(cm<sup>-2</sup>)</th><th>T<sub>kin</sub><sup>CO,b</sup><br>(K)</th><th>n<sub>H<sub>2</sub></sub><sup>CO,c</sup><br>log<br>(cm<sup>-3</sup>)</th><th>N (CO)<sup>d</sup><br>log<br>(cm<sup>-2</sup>)</th><th>Filling<sup>e</sup><br>Factor</th></tr></thead><tbody><tr><td>SMC_NE_3c_1</td><td>20.5<sup>f</sup></td><td>40 ± 10</td><td>3.3 ± -0.0</td><td>16.6<sup>+0.2</sup><sub>-0.4</sub></td><td>40 ± 10</td><td>3.3 ± 0.1</td><td>16.3<sup>+0.3</sup><sub>-1.0</sub></td><td>...</td><td>3.5<sup>+0.7</sup><sub>-0.1</sub></td><td>16.5<sup>+0.4</sup><sub>-1.0</sub></td><td>0.06<sup>+0.03</sup><sub>-0.42</sub></td><td></td></tr><tr><td>SMC_NE_3g_1</td><td>21.1<sup>f</sup></td><td>190 ± 10</td><td>3.0 ± 0.0<sup>f</sup></td><td>16.6 ± 0.1</td><td>190 ± 10</td><td>3.3 ± 0.1</td><td>15.9 ± 0.3</td><td>...</td><td>3.4 ± 0.1</td><td>16.4 ± 0.3</td><td>0.03 ± 0.01</td><td></td></tr><tr><td>SMC_NE_4a_hi_2</td><td>20.4<sup>f</sup></td><td>20 ± 10</td><td>3.8<sup>+0.0f</sup><sub>-0.1</sub></td><td>17.5<sup>+0.7</sup><sub>-2.0</sub></td><td>20 ± 10</td><td>4.0 ± 0.1</td><td>15.9<sup>+0.1</sup><sub>-0.2</sub></td><td>...</td><td>4.0 ± 0.1</td><td>16.2<sup>+0.1</sup><sub>-0.2</sub></td><td>0.10<sup>+0.00g</sup><sub>-0.02</sub></td><td></td></tr><tr><td>SMC_NE_4a_hi_3</td><td>21.0<sup>f</sup></td><td>49</td><td>3.3</td><td>16.3<sup>+0.3</sup><sub>-0.2</sub></td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td>...</td><td></td></tr><tr><td>SMC_NE_4c_low_1</td><td>21.2<sup>f</sup></td><td>30<sup>+10</sup><sub>-20</sub></td><td>2.8 ± 0.0<sup>f</sup></td><td>17.4<sup>+0.7</sup><sub>-0.8</sub></td><td>30<sup>+10</sup><sub>-20</sub></td><td>2.1<sup>+0.1</sup><sub>-0.2</sub></td><td>16.4<sup>+0.5</sup><sub>-0.4</sub></td><td>...</td><td>2.1<sup>+0.1</sup><sub>-0.2</sub></td><td>17.1<sup>+0.6</sup><sub>-0.5</sub></td><td>0.15<sup>+0.09</sup><sub>-0.17</sub></td><td></td></tr></tbody></table>

**Notes.**

a H<sup>0</sup> column density derived from the integrated intensity of the H I 21 cm line over the velocity range defined by the FWHM of the [C II] velocity components (see Section 3.2).

b Kinetic temperature of the H<sub>2</sub> gas associated with C<sup>+</sup> (Section 4.4.3), C<sup>0</sup> (Section 4.4.2), and CO (Section 4.4.1).

c Volume density of the H<sub>2</sub> gas associated with C<sup>+</sup> (Section 4.4.3), C<sup>0</sup> (Section 4.4.2), and CO (Section 4.4.1).

d C<sup>+</sup> (Section 4.4.3), C<sup>0</sup> (Section 4.4.2), and CO (Section 4.4.1) column densities associated with H<sub>2</sub> gas.

e Filling factor of the [C I]- and CO-emitting layers derived from <sup>13</sup>CO observations in Section 4.4.

f The uncertainty is below 0.05.

g The uncertainty is below 0.005.
from diffuse CO-dark H₂ gas clouds. In LOSs associated with PDRs, the [C II] emission from this ISM component tends to be brighter than the emission arising from other components, which is the result of the higher volume densities and temperatures of the gas. The typical contributions to the [C II] emission associated with PDRs are 62% and 66% of the observed emission in LOSs in the LMC and SMC, respectively. The fraction of the [C II] emission arising from CO-dark H₂ shows a large scatter, ranging from ~10% to ~80%. A similar range of CO-dark H₂ fractions is found in clouds in the Milky Way (Langer et al. 2014; Tang et al. 2016). This large scatter could be the result of clouds at different stages in the transition from diffuse to dense molecular gas, with clouds having varying fractions of CO-dark H₂ gas.

Our separation between [C II] from diffuse CO-dark H₂ gas and from warm and dense PDRs is sensitive to our ability to detect the ¹³CO J = 1 → 0 and/or ¹³CO J = 3 → 2 lines. For the thermal pressure of the diffuse ISM in the LMC and SMC derived in Section 4.3 and a kinetic temperature of the CO-dark H₂ of 49 K (Section 5.1), we obtain n (H₂) = 694 cm⁻³ in the LMC and 2040 cm⁻³ in the SMC. For these conditions, the ¹³CO column density required for a 3σ detection of ¹³CO J = 1 → 0 and ¹³CO J = 3 → 2 is 4 × 10¹⁴ and 3 × 10¹⁵ cm⁻² for the LMC and 4.5 × 10¹⁴ and 1 × 10¹⁵ cm⁻² for the SMC, respectively.

### 3.4. Uncertainties

The main uncertainty in our decomposition originates from the assumption of the temperature of atomic and molecular gas. The assumed kinetic temperature is used in the derivation of thermal pressures (Section 4.3) and in the calculation of the contribution of atomic gas to the observed [C II] emission. Observations of H I absorption against background continuum sources in the LMC and SMC suggest temperatures of the atomic gas as low as 10 K, with typical values of 30–40 K (Dickey et al. 2000; Marx-Zimmer et al. 2000). Note that H I absorption features are stronger with lower kinetic temperatures of the H I gas; thus, there might be a systematic tendency to detect colder atomic gas in absorption. At these low temperatures, most of the gas-phase carbon is likely to be in the form of CO rather than C⁺, as suggested by the detection of the ¹²CO J = 1 → 0 line in several H I absorption sources in the LMC (Marx-Zimmer et al. 1999). While it is possible that the H I temperatures are lower in the Magellanic clouds compared with our Galaxy, we considered it unlikely that [C II] will be detected for temperatures below 40 K, as the [C II] intensity is significantly reduced for kinetic temperatures below this value. We therefore assumed a temperature of 70 K for the atomic gas in our calculations, which is an intermediate value between the temperature derived from H I absorption in the Magellanic clouds and the 100 K typically assumed in the Milky Way (Pineda et al. 2013). For the temperature of the H₂ gas, we assumed that T<sub>kin</sub> = 0.7T<sub>kin</sub><sup>0</sup>, which is based on results in the Milky Way that suggest a temperature of 70 K for the CO-dark H₂ layer (Goldsmith 2013) and is in agreement with PDR model calculations (e.g., Goldsmith et al. 2016). The temperature of the CO-dark H₂ gas corresponds to 49 K for our assumed H I temperature of 70 K.

We estimated the uncertainty associated with our H I and H₂ kinetic temperature assumption by calculating the range in the [C II] intensities associated with the different ISM phases

resulting from assuming T<sub>kin</sub><sup>H0</sup> = 40 and T<sub>kin</sub><sup>H0</sup> = 100 K in the LMC and T<sub>kin</sub><sup>H0</sup> = 55 and T<sub>kin</sub><sup>H0</sup> = 100 K in the SMC. We used T<sub>kin</sub><sup>H0</sup> = 55 K as a lower kinetic temperature limit in the SMC, as this is the lowest temperature at which we could find solutions of the thermal pressure in the SMC (see Section 4.3). The derived uncertainties are shown as error bars in Figure 5.

Note that the assumed temperature of the CO-dark H₂ gas is lower than the H₂ rotational temperature in LOSs compiled by Welty et al. (2012) in the LMC and SMC. These H₂ temperatures range from 40 to 120 K, with average values of 85 K in both the LMC and SMC. This range corresponds, for our assumption of T<sub>kin</sub><sup>H2</sup> = 0.7T<sub>kin</sub><sup>H0</sup>, to a range in H₀ kinetic temperature between 57 and 170 K. However, the LOSs studied by Welty et al. (2012) have total hydrogen column densities that are much lower compared to the values derived in our survey (see Figure 13) and therefore are diffuse and likely warmer LOSs.

Another source of uncertainty originates from beam filling effects produced by the different angular resolutions of our observations. The angular resolution of the data used in our [C II] decomposition varies from 12'' for [C II] to 60'' for H I. As described in the Appendix, we derived correction factors using the 160 µm map for [C II] and H I and the 24 µm continuum maps for the SHASSA observations. We estimated that the contribution from the different ISM phases to the observed [C II] emission would typically change by 10%, 14%, 17%, and 16% for PDRs, CO-dark H₂ gas, H I gas, and ionized gas in the LMC, respectively, and by 6%, 4%, 15%, and 8% for PDRs, CO-dark H₂ gas, H I gas, and ionized gas in the SMC, respectively. We chose not to include the variation in the intensities due to beam filling factors in Figure 5 due to the uncertainty in whether the 24 µm and 160 µm emission represent the spatial distribution of the [C II] associated with ionized and neutral gas, respectively.

## 4. Physical Parameters

### 4.1. Determination of Visual Extinction

We used the Herschel HERITAGE (Meixner et al. 2013) dust continuum maps of the LMC and SMC to determine the dust column density in terms of the visual extinction A<sub>V</sub> in our sample. We assumed that the dust SED in the LMC and SMC can be described by the emission predicted from an optically thin modified blackbody at an equilibrium temperature T<sub>dust</sub>. The dust opacity at 160 µm is given by

$$ \tau_{160 \mu m} = \frac{I_{160 \mu m}}{B_{\nu}(T_{\text{dust}}, 160 \mu m)}, \qquad (8) $$

where I<sub>160 µm</sub> is the dust continuum intensity at 160 µm and B<sub>ν</sub>(T<sub>dust</sub>, 160 µm) is the intensity of a blackbody with a temperature T<sub>dust</sub> at 160 µm. We determined the opacity at 160 µm using the HERITAGE 160 µm map together with the dust temperatures fitted by Gordon et al. (2014) for a single-temperature blackbody modified by a broken power-law emissivity, their preferred dust emission model. The 160 µm and dust temperature maps have a common resolution of 40''.

The average (± standard deviation) dust temperature in our sample is 22.8 ± 4.1 K in the LMC and 23.1 ± 2.5 K in the SMC.
Lee et al. (2015) presented maps of the visual extinction in the LMC and SMC derived from dust continuum emission. They converted the 160 µm opacity to $A_V$ using $A_V = 2200\tau_{160\mu m}$, which is the average of different methods to relate these quantities based on data in the Milky Way. (We refer the reader to their paper for more details on how this conversion factor is determined.) The different methods described in Lee et al. (2015) show a scatter from the adopted conversion factor between $A_V$ and $\tau_{160 \mu m}$ of about 40%, and we adopted this value as the uncertainty in our determination of $A_V$.

We compared the values of $A_V$ derived here with those presented by Lee et al. (2015), who derived the 160 µm opacity using the HERITAGE 160 µm map and dust temperatures determined by fitting a single-temperature blackbody with an assumed $\beta = 1.5$ wavelength dependence of the dust opacity. We found that our $A_V$ values are in reasonable agreement with those derived by Lee et al. (2015), with differences typically within the assumed 40% uncertainty.

## 4.2. Dust-to-Gas Ratio

The DGR is a fundamental parameter relating the quantities of gas and dust in the ISMs of galaxies, and the variation of the DGR as a function of environment reflects the evolution of galaxies (Dwek 1998). In the following, we test our determination of gas column densities and visual extinctions in our sample by checking whether these two quantities are related by the DGRs that are typically found in the LMC and SMC. We initially test diffuse LOSs in our sample, but we will extend this comparison to our entire sample in Section 5.1 to test our determination of $H_2$ column densities.

The ratio of the total hydrogen column density, $N(H) = N(H^+) + N(H^0) + 2N(H_2)$, to color excess, $E(B-V)$, in the ISM of the Milky Way is $N(H)/E(B-V) = 5.8 \times 10^{21} \text{ cm}^{-2}\text{mag}^{-1}$ (Bohlin et al. 1978). Assuming a ratio of the total to selective extinction of $R_V = 3.1$ results in a relationship between the hydrogen column density and visual extinction of $N(H)/A_V = 1.88 \times 10^{21} \text{ cm}^{-2}\text{mag}^{-1}$ in the Milky Way. In the LMC, $N(H)/E(B-V)$ is observed to be $2 \times 10^{22} \text{ cm}^{-2}\text{mag}^{-1}$ (Koorneef 1982; Fitzpatrick 1985a), while in the SMC, $N(H)/E(B-V)$ ranges between $2.6 \times 10^{22}$ (Martin et al. 1989) and $8.7 \times 10^{22} \text{ cm}^{-2}\text{mag}^{-1}$ (Fitzpatrick 1985b). We adopted an intermediate value for the SMC of $6.6 \times 10^{22} \text{ cm}^{-2}\text{mag}^{-1}$. Assuming $R_V = 3.4$ in the LMC and $R_V = 2.7$ in the SMC (Gordon et al. 2003) results in $N(H)/A_V = 5.56 \times 10^{21} \text{ cm}^{-2}\text{mag}^{-1}$ for the LMC and $2.43 \times 10^{22} \text{ cm}^{-2}\text{mag}^{-1}$ for the SMC. Note that the value of $R_V$ assumed here corresponds to that in diffuse regions, and it might be larger in denser regions (Weingartner & Draine 2001; Whittet et al. 2001). Considering that a given LOS might intersect both dense and diffuse regions, Whittet et al. (2001) estimated in the Milky Way an effective $R_V$ along the LOS that increases up to ~4.0 for $A_V \simeq 10$ mag. Thus, the $N(H)/A_V$ ratio can be lower by up to ~30% in denser LOSs.

We tested our estimate of the visual extinction by comparing $A_V$ with $N(H^+) + N(H^0)$ in LOSs where [C II] was not detected. These LOSs (five in the LMC and seven in the SMC) likely have little or no $H_2$ and/or have densities and temperatures that are insufficient for producing [C II] emission that will be detectable in the sensitivity limits of our survey. They are therefore well suited for our comparison between visual extinctions and the $H^+ + H^0$ column densities, as they are likely to have a negligible or no contribution from dust

<figure><img src="image_11.png" /><figcaption>**Figure 7.** The $H^+ + H^0$ column density as a function of the visual extinction associated with atomic gas for LOSs in the LMC and SMC where H I and H α are the only spectral lines detected and therefore assumed to be diffuse gas. The lines in the figure represent the total hydrogen column density predicted for a given $A_V$ assuming $N(H)/A_V = 7.17 \times 10^{21} \text{ cm}^{-2}\text{mag}^{-1}$ for the LMC, $1.68 \times 10^{22} \text{ cm}^{-2}\text{mag}^{-1}$ for the SMC, and $1.88 \times 10^{21} \text{ cm}^{-2}\text{mag}^{-1}$ for the Milky Way (Section 4.2).</figcaption></figure>

associated with $H_2$ to the observed $A_V$. The $H^+$ column densities for these LOSs were determined in Section 3.1 and, on average, represent 5% and 1.4% of the $N(H^+) + N(H^0)$ column densities in the LMC and SMC, respectively. The $A_V$ and $N(H^+)$ data points used here were derived from dust continuum and H α maps that were smoothed to the 60" resolution of the H I data. In Figure 7, we show $N(H^+) + N(H^0)$ as a function of $A_V$ for LOSs in our sample where [C II] was not detected. We also include lines that represent the predicted $H^+ + H^0$ column density for a given value of $A_V$ using the DGRs for the LMC, SMC, and Milky Way discussed above. Given the uncertainties in the determination of $A_V$, we consider that our derived values of $A_V$ are consistent with the observed $H^0$ column densities and the independently measured gas-to-dust ratios in the LMC and SMC.

## 4.3. Thermal Pressure of the Diffuse ISM in the Magellanic Clouds

The thermal pressure of the diffuse ISM is an important parameter that, despite being a small fraction of the total ISM pressure (Boulares & Cox 1990), plays a fundamental role in the phase transition from warm and diffuse to cold and dense atomic gas (Pikel’Ner 1968; Field et al. 1969; Wolfire et al. 1995; Wolfire et al. 2003; Jenkins & Tripp 2011). Thus, studying the diffuse ISM thermal pressure is important for our understanding of molecular cloud formation and the regulation of star formation in galaxies (McKee 1989; Blitz & Rosolowsky 2006; Ostriker et al. 2010).

In the Magellanic clouds, the thermal pressure of the gas is expected to be larger than that in the solar neighborhood because of the larger FUV fields and lower dust abundances affecting the thermal balance (Sandstrom et al. 2010; Welty
et al. 2016; M. Wolfire et al. 2017, in preparation). Bolatto et al. (2011) compared the relationship between gas surface density and star formation rate in the SMC and found that a thermal pressure $p_{th}/k_B = 3 \times 10^4$ K cm$^{-3}$ is required to find agreement between the observations and the theoretical predictions from Ostriker et al. (2010).

An estimate of the thermal pressure of the diffuse ISM is also a requirement for the determination of the column density of the CO-dark H$_2$ gas using [C II] observations (Pineda et al. 2013). In the following, we derive the thermal pressures of the diffuse ISM of the LMC and SMC by using LOSs in which we detect H I and [C II] but no CO and/or [C I] is detected. These LOSs (14 in the LMC and six in the SMC) are likely diffuse, and the [C II] emission arises from both cold atomic and CO-dark H$_2$ gas.

As discussed in Section 3, the observed [C II] intensity is the sum of the emission associated with the ionized, atomic, and molecular gas components. For each LOS used here, we subtracted the [C II] emission associated with ionized gas derived in Section 3.1. The remaining [C II] intensity is thus described by six parameters (see Equation (2)): the kinetic temperature ($T_{kin}^{H0}$ and $T_{kin}^{H2}$), volume density ($n_{H0}$ and $n_{H2}$), and C$^+$ column densities ($N(C^+)_H$ and $N(C^+)_H$) of gas associated with the H$^0$ and H$_2$ layers.

The total column density of C$^+$ along the LOS is the sum of the column of C$^+$ associated with the different collisional partners,

$$N(C^+)_\text{total} = N(C^+)_H + N(C^+)_H + N(C^+)_\text{e}^- \quad (9)$$

We used the H I 21 cm observations to determine the column density of atomic hydrogen that is associated in velocity with the observed [C II] emission in Section 3.2 and the H$^+$ column density associated with ionized gas in Section 3.1. The H$^0$ and H$^+$ column densities are in turn converted to a C$^+$ column density by multiplying the carbon fractional abundances of the LMC and SMC discussed in Section 3.

To derive the column density of C$^+$ associated with H$_2$, we need an estimate of the total hydrogen column density along the LOS. The visual extinction along the LOS is the sum of that from dust associated with molecular hydrogen, with CNM and WNM atomic gas, and with ionized gas, so that

$$A_V^{\text{total}} = A_V^{H_2} + A_V^{H_0, \text{CNM}} + A_V^{H_0, \text{WNM}} + A_V^{H^+} \quad (10)$$

We estimated $A_V^{H1,CNM}$ and $A_V^{H1,WNM}$ by converting the integrated intensities of H I associated (CNM) and not associated (WNM) to the observed [C II] emission and applying the DGRs discussed in Section 4.2. The same DGR was applied to the H$^+$ column densities to estimate the visual extinction associated with ionized gas. The visual extinction associated with H$_2$ was then estimated by subtracting that associated with H$^0$ and H$^+$ from the observed value. We then converted $A_V^{H2}$ to an H$_2$ column density using a DGR and to $N(C^+)_H$ using a fractional abundance of carbon.

With an estimate of $N(C^+)_H$ and $N(C^+)_H$, the remaining parameters are the kinetic temperatures and volume densities of the molecular and cold atomic gas. We assumed that, for the diffuse LOS where only H I and [C II] are detected, there is thermal pressure equilibrium in the interface between the cold

<figure><img src="image_12.png" /><figcaption>Figure 8. Histograms of the thermal pressure of the diffuse ISM of the LMC and SMC derived in Section 4.3.</figcaption></figure>

H I and the CO-dark H$_2$ layers,

$$p_{\text{th}}/k_{\text{B}} = n_{\text{H}} T_{\text{kin}}^{\text{H0}} = n_{\text{H2}} T_{\text{kin}}^{\text{H2}} \quad (11)$$

Assuming a kinetic temperature for the H I layer of 70 K and for the CO-dark H$_2$ layer of 49 K (see discussion in Section 3.4), we searched for a thermal pressure that could reproduce the observed [C II] intensity. Figure 8 shows histograms of the thermal pressure derived in the LMC and SMC. The average (± standard deviation) thermal pressure is $p_{th}/k_B = 10^{4.5 \pm 0.3}$ K cm$^{-3}$ in the LMC and $p_{th}/k_B = 10^{5.0 \pm 0.2}$ K cm$^{-3}$ in the SMC. We note that components that are associated with the 30 Doradus nebula in the LMC tend to higher thermal pressure compared to those in other regions of the LMC. Without the LOSs associated with 30 Doradus, we found an average thermal pressure in the LMC of $p_{th}/k_B = 10^{4.2 \pm 0.1}$ K cm$^{-3}$, while LOSs associated with 30 Doradus have $p_{th}/k_B = 10^{4.9 \pm 0.2}$ K cm$^{-3}$. These LOSs
appear to be influenced by the feedback effects of the R136 super star cluster (Pellegrini et al. 2011). The average value for the SMC is in reasonable agreement with that suggested by Bolatto et al. (2011).

We tested the sensitivity of the derived thermal pressures to the assumed kinetic temperature of the gas. For low temperatures, the thermal pressure solutions become uncertain as higher volume densities are needed to reproduce the observed emission, and collisional de-excitation becomes important, reducing the dependence of the [C II] intensity on volume density (Goldsmith et al. 2012). In the LMC, the thermal pressure increases by a factor of 1.2 between 100 and 70 K and by a factor of 6 between 100 and 40 K. In the SMC, we could not find reliable solutions for temperatures lower than 55 K. Between 100 and 70 K, the thermal pressure increases by a factor of 1.5, and, between 100 and 55 K, it rises by a factor of 3.5.

We also studied the effects of beam dilution in our determination of thermal pressure in the LMC and SMC. For that purpose, we smoothed the dust continuum data used to calculate $A_V$ and the $H\alpha$ data used to estimate the contribution from ionized gas to the [C II] emission to the 60" resolution of the H I data used to calculate $N(H^0)$. We additionally corrected the observed [C II] emission with the beam dilution factor we estimated in the Appendix to correspond to emission at 60" resolution. We used this data to recalculate the thermal pressures in our subsample. We found that applying a beam dilution factor to the data used in our calculation results in thermal pressures that are reduced by factors of ~1.4 in both the LMC and SMC.

We note that the statistical significance of our determination of the thermal pressure in the LMC and SMC is limited by the small number of LOSs we were able to use. Using the technique presented here in future large-scale maps of the Magellanic clouds in [C II] will dramatically improve the significance of this result and allow us to compare the distribution of thermal pressures in the LMC and SMC with that in the Milky Way (Jenkins & Tripp 2011).

## 4.4. Kinetic Temperatures and H₂ Column and Volume Densities

We constrained the properties of the CO- and [C I]-emitting gas by comparing line intensity ratios and absolute intensities of [C I], $^{12}$CO, and $^{13}$CO with the results of a radiative transfer code. We used the RADEX non-Local Thermodynamic Equilibrium (LTE) radiative transfer code (van der Tak et al. 2007) using the uniform sphere approximation to calculate line intensities as a function of the kinetic temperature, H₂ volume density, and the column density per velocity interval $N/\Delta v$ of the species of interest. The collisional rate coefficients were taken from the Leiden Atomic and Molecular Database (LAMDA; Schöier et al. 2005).

In the PDRs located at the borders of molecular clouds, the innermost region is at a higher density deeper in the cloud, and the outermost region is at a lower density. Thus, the physical conditions of the C⁰ layer are not necessarily the same as those at the CO layer. We analyzed the [C I] and CO data separately, assuming their emission originates from spheres with uniform volume density and kinetic temperature. We nevertheless assumed that the [C I] and CO emission originates from the same regions with identical filling factors. This assumption is justified considering that the CO and [C I] emission arises from

the FUV-illuminated surfaces of clouds that are unresolved at the scale of our observations. For optically thin clumps (or for clumps with moderate opacity), the observed-to-model integrated intensity ratio gives the ratio of the beam-averaged column density to the clump intrinsic column density, assuming that the total integrated intensity is proportional to the number of clumps; i.e., no velocity crowding is present. We therefore determined the beam filling factor by comparing the observed absolute intensity of an optically thin line ($^{13}$CO $J=3 \to 2$ or $^{13}$CO $J=1 \to 0$) with a constrained model.

We compared the observed line intensity ratios with the predictions from a grid of RADEX models. The model grids predict line ratios for kinetic temperatures ranging from 10 to 200 K in steps of 10 K as a function of the CO or C⁰ column densities per velocity interval, $N_{CO}/\Delta v$ (or $N_{C0}/\Delta v$), and H₂ volume density, $n_{H2}$. The coverage in $N_{CO}/\Delta v$ and $N_{C0}/\Delta v$ is $10^{15} - 10^{19} \text{ cm}^{-2}(\text{km s}^{-1})^{-1}$, and, in $n_{H2}$, the coverage is $10^2-10^6 \text{ cm}^{-3}$, both in steps of 0.1 dex. In our calculations, we initially assumed $\Delta v = 1 \text{ km s}^{-1}$, but we later converted the constrained $N_{C0}/\Delta v$ and $N_{CO}/\Delta v$ to total C⁰ and CO column density by multiplying by the line FWHM resulting from the Gaussian fit in each velocity component. For velocity components where the lines in the numerator of the line ratios are not detected but those in the denominator are, we used the $3\sigma$ upper limits of the lines in the numerator to calculate line ratios that were used to discard solutions in the RADEX grids (Figure 9). In Table 5, we summarize the physical parameters derived in the LMC and SMC. We only list velocity components in which the [C II] line is detected.

We quantified how a RADEX model $M(T_{\text{kin}}, N_{\text{C}0}/\Delta v, n_{\text{H2}})$ reproduces a given line ratio $R$ with uncertainty $\sigma_R$ by calculating $\chi_R^2 = \sum_i (R - M_i)^2 / \sigma_R^2$ over the entire model grid. We later added the resulting $\chi_R^2$ for all line ratios available in a given velocity component to determine a total $\chi_{\text{total}}^2$. We then identified the best-matching model by searching for the minimum $\chi_{\text{total}}^2$ in the model grid, $\chi_{\text{total,min}}^2$. To determine the uncertainties in the derived parameters, we defined a region in the parameter space where models have a $\chi_{\text{total}}^2$ that is within three times the minimum $\chi_{\text{total,min}}^2$. We then determined the range of possible values of a given parameter within this region when the other two are fixed. For well-constrained models, $\chi_{\text{total,min}}^2$ can be very small, and the adjacent models in the parameter space often have $\chi_{\text{total}}^2$ that are higher than $3\chi_{\text{total,min}}^2$. In this case, the accuracy with which a model parameter can be determined is given by the step size of the model grid.

We studied the effect of beam dilution in the line ratios used in our analysis in the Appendix. Note that we did not correct the line ratios for beam dilution effects in our analysis, due to the uncertainty of whether dust continuum emission traces the distribution of gas. In Table 6, we present the effect of beam dilution in the line ratios that are calculated using observations at different angular resolutions. We found that these line ratios typically show a 10% variation. This small variation suggests that the observed structures are relatively extended at the resolution of our observations. In Table 7, we summarize the impact of the variation in the line ratios due to beam dilution in the derived physical properties of the C⁰ and CO layers. The kinetic temperature, column density, and filling factor of the C⁰ and CO layers are typically affected by factors between 1.1 and 1.6. The largest impact of the variation in the line ratios is in the derived H₂ volume density with variations between 1.5
<figure><img src="image_13.png" /><figcaption>**Figure 9.** Sample comparison between the observed CO and ¹³CO line ratios and the predictions from the RADEX model. The lines represent the observed line ratios ± the uncertainties in their determination. The comparisons are shown at the kinetic temperature at which the minimum value of χ²<sub>total,min</sub> is located. In the middle and lower panels, we show cases in which the 3σ upper limit of the ¹³CO *J* = 1 → 0/¹²CO *J* = 1 → 0 and/or the ¹²CO *J* = 7 → 6/*J* = 3 → 2 ratios are used to discard solutions.</figcaption></figure>

<table><caption>**Table 6**
Beam Dilution Effect on Line Ratios</caption><thead><tr><th>Ratio<sup>a</sup></th><th>Θ<sub>a</sub>/Θ<sub>b</sub><sup>b</sup></th><th>LMC<sup>c</sup></th><th>SMC<sup>c</sup></th></tr></thead><tbody><tr><td><sup>12</sup>CO J = 7 → 6/*J* = 3 → 2</td><td>26<sup>"</sup>5/17<sup>"</sup>5</td><td>1.07</td><td>1.08</td></tr><tr><td><sup>12</sup>CO J = 3 → 2/*J* = 1 → 0</td><td>17<sup>"</sup>5/33<sup>"</sup></td><td>1.11</td><td>1.11</td></tr><tr><td>[C I] <sup>3</sup>P<sub>2</sub> → <sup>3</sup>P<sub>1</sub>/<sup>3</sup>P<sub>1</sub> → <sup>3</sup>P<sub>0</sub></td><td>26<sup>"</sup>5/44<sup>"</sup></td><td>1.09</td><td>1.05</td></tr></tbody></table>

**Notes.**

<sup>a</sup> Line ratio involving spectral lines observed at different angular resolutions.

<sup>b</sup> Angular resolution of the spectral line in the numerator and denominator of the line ratio.

<sup>c</sup> Average variation in a line ratio due to beam dilution effects.

<table><caption>**Table 7**
Impact of Beam Dilution Effects on Derived Physical Parameters</caption><thead><tr><th>Physical Parameter</th><th>LMC<sup>a</sup></th><th>SMC<sup>a</sup></th></tr></thead><tbody><tr><td>N (C<sup>0</sup>)</td><td>1.2</td><td>1.5</td></tr><tr><td>N (CO)</td><td>1.4</td><td>1.1</td></tr><tr><td>n<sup>C0</sup><sub>H2</sub></td><td>2.2</td><td>3.2</td></tr><tr><td>n<sup>CO</sup><sub>H2</sub></td><td>4</td><td>1.5</td></tr><tr><td>T<sup>C0</sup><sub>kin</sub></td><td>1.2</td><td>1.3</td></tr><tr><td>T<sup>CO</sup><sub>kin</sub></td><td>1.5</td><td>1.6</td></tr><tr><td>Filling factor</td><td>1.4</td><td>1.2</td></tr></tbody></table>

**Note.**

<sup>a</sup> Average variation on the physical parameter resulting from beam dilution effects.

and 4. Note that the largest variation seen, in the resulting H₂ volume density of CO gas in the LMC, is dominated by three sources, PDR4_RIDGE_1, PDR3_NW, and NT77, which vary by factors of 25, 8, and 8, respectively. Without these three sources, we found typical variations of a factor of 1.6 for n<sup>CO</sup><sub>H₂</sub> in the LMC.

### 4.4.1. CO-Emitting Layer

We started by deriving the excitation conditions of the CO-emitting cloud layer. Depending on availability, we compared the observed ¹²CO *J* = 7 → 6/*J* = 3 → 2, ¹²CO *J* = 3 → 2/1 → 0, ¹³CO *J* = 3 → 2/¹²CO *J* = 3 → 2, and ¹³CO *J* = 1 → 0/¹²CO *J* = 1 → 0 intensity ratios with the predictions from a grid of RADEX models. In Figure 9, we show an example of the comparison between the model grid and the observed line ratios. We assumed a ¹²CO/¹³CO abundance ratio of 49 derived in the N113 region in the LMC by Wang et al. (2009).

Note that there are several velocity components in the LMC and SMC where ¹²CO *J* = 1 → 0 was detected but ¹³CO *J* = 1 → 0 and/or ¹³CO *J* = 3 → 2 were not. We were thus unable to derive the physical conditions of the CO-emitting layer in these LOSs. These LOSs are likely diffuse (see Section 3.3), as ¹²CO *J* = 1 → 0 is known to be detectable even in low column density regions, where ¹³CO *J* = 1 → 0 is not detected, and CO is not the dominant form of gas-phase carbon (Pineda et al. 2010).

We found that single-component models often provide a reasonable fit to the observed line ratios in our sample. There are, however, some indications by the ¹²CO/¹³CO ratios of a colder gas component. When a best-fitting model is identified, we compared the observed ¹³CO *J* = 3 → 2 intensities,
<figure><img src="image_14.png" /><figcaption>**Figure 10.** Sample comparison between the [C I] $^3P_2 \rightarrow {}^3P_1/{}^3P_1 \rightarrow {}^3P_0$ line intensity ratio and the peak intensity of the [C I] $^3P_1 \rightarrow {}^3P_0$ line with the grid of RADEX models. The comparison is shown at the kinetic temperature at which the minimum value of $\chi^2$ is found.</figcaption></figure>

assuming this line to be optically thin, with those predicted by the model to estimate the beam filling factor. If $^{13}$CO $J=3 \rightarrow 2$ was not detected, we instead used $^{13}$CO $J=1 \rightarrow 0$ when available. We derived the physical conditions of 23 velocity components in 20 LOSs where at least two pairs of line ratios could be calculated. The average value and standard deviation of the physical conditions in the sample are 43 ± 19 K in the LMC and 50 ± 20 K in the SMC for the kinetic temperature, $10^{16.7 \pm 0.5}$ cm$^{-2}$ in the LMC and $10^{16.6 \pm 0.3}$ cm$^{-2}$ in the SMC for the beam-averaged CO column density, $10^{5.1 \pm 0.9}$ cm$^{-3}$ in the LMC and $10^{4.0 \pm 0.3}$ cm$^{-3}$ in the SMC for the H$_2$ volume density, and (0.2 ± 0.2) in the LMC and (0.20 ± 0.3) in the SMC for the filling factor.

### 4.4.2. [C I]-Emitting Layer

We constrained the volume, column density, and kinetic temperature of the [C I]-emitting layer, comparing the observed [C I] $^3P_2 \rightarrow {}^3P_1/{}^3P_1 \rightarrow {}^3P_0$ line intensity ratio and the peak intensity of the [C I] $^3P_1 \rightarrow {}^3P_0$ line with the grid of RADEX models. We show a sample comparison between observations and models in Figure 10.

For kinetic temperatures above ~50 K and C$^0$ column densities per velocity interval below $\sim 10^{18}$ cm$^{-2}$ (K km s$^{-1}$)$^{-1}$, a typical range of solutions found in the [C I]-emitting layer, the [C I] $^3P_2 \rightarrow {}^3P_1/{}^3P_1 \rightarrow {}^3P_0$ ratio provides a good constraint on the volume density of the [C I]-emitting gas. The peak intensity of the [C I] $^3P_1 \rightarrow {}^3P_0$ line, corrected by the filling factor derived in the CO-emitting layer, depends mostly on the C$^0$ column density per velocity interval, thus constraining $N(C^0)/\Delta v$. The multiplication by a filling factor is appropriate because, in the range of kinetic temperatures and H$_2$ densities inferred in our analysis, the [C I] $^3P_1 - {}^3P_0$ line is likely optically thin.

While the [C I] lines typically provide good constraints on the column and volume density of the [C I]-emitting gas, the kinetic temperature is not well constrained. We therefore further constrained the solutions by requiring that the [C I]-emitting region have a lower or equal volume density and higher or equal kinetic temperature than the CO-emitting layer.

This assumption is justified, as we expect that [C I] and mid-J CO lines are emitted from adjacent layers in a PDR-like structure$^{18}$, where the neutral carbon is located in a somewhat warmer region with equal or lower volume density compared to the one in which CO is located.

There are some velocity components in which only the [C I] $^3P_1 \rightarrow {}^3P_0$ line is detected; therefore, we could only calculate a $3\sigma$ upper limit to the [C I] $^3P_2 \rightarrow {}^3P_1/{}^3P_1 \rightarrow {}^3P_0$ ratio. In these cases, we assumed that the kinetic temperature and volume density at the [C I]-emitting layer are the same as those at the CO-emitting layer, and we used the [C I] $^3P_1 \rightarrow {}^3P_0$, corrected with its corresponding filling factor, to constrain the C$^0$ column density per velocity interval.

We derived the physical conditions of 22 velocity components in 20 LOSs. The average value and standard deviation of the physical conditions in the [C I]-emitting layer are 91 ± 60 K in the LMC and 73 ± 62 K in the SMC for the kinetic temperature, $10^{16.4 \pm 0.4}$ cm$^{-2}$ in the LMC and $10^{16.2 \pm 0.3}$ cm$^{-2}$ in the SMC for the beam-averaged C$^0$ column density, and $10^{4.8 \pm 1.4}$ cm$^{-3}$ in the LMC and $10^{3.5 \pm 0.5}$ cm$^{-3}$ in the SMC for the H$_2$ volume density.

### 4.4.3. [C II]-Emitting Layer

As mentioned above, the observed [C II] emission is the result of the combined emission from molecular, atomic, and ionized gas. In Sections 3.1 and 3.2, we estimated the physical conditions of the ionized and atomic gas associated with C$^+$ to determine their contribution to the observed [C II] emission. In the following, we describe the derivation of the physical conditions of the [C II]-emitting gas associated with molecular gas.

We estimated the volume density of the [C II]-emitting layer associated with molecular hydrogen assuming a thermal pressure and a kinetic temperature. In Section 4.3, we estimated the typical thermal pressure of the diffuse gas of $3.4 \times 10^4$ and $1 \times 10^5$ K cm$^{-3}$ in the LMC and SMC, respectively. In LOSs where H I and [C II] are detected but no [C I], CO, or $^{13}$CO emission is detected, we assumed the derived diffuse ISM pressure and a kinetic temperature of 49 K (Section 3.4) to derive a typical volume density of 694 and 2040 cm$^{-3}$ for the LMC and SMC, respectively. We also applied this criteria to LOSs where weak CO emission was detected but no $^{13}$CO was detected.

For velocity components in which we were able to determine the physical conditions of the [C I]- or $^{12}$CO-emitting layer, we used a thermal pressure for the [C II]-emitting layer that is the geometrical mean between the diffuse ISM pressure derived in Section 4.3 and that of the [C I]- or $^{12}$CO-emitting layer, in case we were not able to determine the conditions at the [C I]-emitting layer. For the kinetic temperature, we assumed 49 K, in case the temperature of the denser molecular gas was lower than or equal to this temperature. We adopted the kinetic temperature of the denser molecular gas, in case it was larger than 49 K. With these assumptions, we derived typical volume densities (and standard deviations) of $10^{3.1 \pm 0.5}$ cm$^{-3}$ for the LMC and $10^{3.2 \pm 0.2}$ cm$^{-3}$ for the SMC.

With the estimated volume density and kinetic temperatures of the [C II]-emitting layer, we used Equation (2) to derive the

<footnote>$^{18}$ For an example of a typical temperature distribution and abundance structure of the main carbon species in PDRs, see, e.g., Figures 7 and 9 in Tielens & Hollenbach (1985).</footnote>
$C^+$ column density for each velocity component. Typical values of the beam-averaged column densities are $10^{17.4 \pm 0.3}$ and $10^{17.0 \pm 0.2}$ cm$^{-2}$ for the LMC and SMC, respectively.

## 5. Discussion

### 5.1. The H-to-H₂ Transition

The abundance of molecular hydrogen in interstellar clouds is set by the balance between H₂ formation on dust grains and photodestruction by the ambient FUV photons. The gas transitions from being H⁰-dominated to H₂-dominated when the dust column density is large enough to allow H₂ self-shielding to halt photodestruction, leading to a rapid increase in H₂ abundance (e.g., Gnedin et al. 2009; Krumholz et al. 2009; Sternberg et al. 2014; Bialy & Sternberg 2016). The transition between atomic and molecular hydrogen marks the onset of the formation of molecular clouds and therefore has important implications in star formation and galaxy evolution.

The atomic gas component in clouds is often traced by observations of the H I 21 cm line, while molecular gas is determined by observing the ¹²CO J = 1 → 0 line, which is converted to an H₂ column density using an $X_{CO}$ conversion factor. Observations of CO, however, trace mostly well-shielded regions, and the diffuse H₂ component (CO-dark H₂) is not well probed. Diffuse H₂ has been observed in absorption against stars in the FUV (e.g., Tumlinson et al. 2002; Cartledge et al. 2005; Sheffer et al. 2008; Welty et al. 2012) but is limited to individual LOSs and tends to trace low H₂ column densities, as LOSs with higher molecular content will absorb the starlight required for reliable FUSE measurements. Observations of [C II], which can potentially image large regions in the Magellanic clouds, can be used to trace the CO-dark H₂ component and study how it relates to the atomic gas.

We studied the transition from atomic to molecular gas in the Magellanic clouds by estimating the column density of H⁰, N(H⁰), and H₂, N(H₂), in our sample. As discussed in Section 3.2, we estimated the column density of atomic hydrogen using N(H⁰) = 1.82 × 10¹⁸ I(H I) cm⁻², with I(H I) in units of K km s⁻¹. The typical uncertainty in the H⁰ column density is 10¹⁹.⁸ cm⁻². To calculate the H₂ column density, we first calculated the total carbon column density, N(total(C) = N(C⁺)H₂ + N(C⁰) + N(CO), where N(C⁺)H₂ is the column density of C⁺ associated with H₂ (Section 4.4.3). We converted from N(total(C) to N(H₂) by applying the gas-phase carbon fractional abundances for the LMC and SMC discussed in Section 3. The estimated N(H₂) can include well-shielded gas associated with C⁰ and CO, as well as the diffuse CO-dark H₂ gas associated with C⁺. The uncertainties in N(H₂) are the result of the propagation of the uncertainties from N(C⁺), N(C⁰), and N(CO). For N(C⁺), the uncertainty is given by the propagation of the errors in the integrated intensity of the [C II] line and the errors on the thermal pressures due to the uncertainties in the physical conditions of the [C I] or CO layer. The uncertainties of N(C⁰) and N(CO) are estimated as part of the excitation analysis described in Section 4.4.

In order to test our determination of the H⁰ and H₂ column densities, in Figure 11, we show a comparison between the N(H⁺) + N(H⁰) + 2N(H₂) hydrogen column density derived

<figure><img src="image_15.png" /><figcaption>**Figure 11.** Comparison between the N(H⁺) + N(H⁰) + 2N(H₂) hydrogen column densities derived in this study and the visual extinction derived from dust continuum observations. The lines represent the total hydrogen column density predicted for a given $A_V$, assuming N(H)/$A_V$ = 7.17 × 10²¹ cm⁻² mag⁻¹ for the LMC and 1.68 × 10²² cm⁻² mag⁻¹ for the SMC (Section 4.2).</figcaption></figure>

from Hα, H I, [C II], [C I], and CO observations and the visual extinction derived from the dust continuum (Section 4.1). We include the N(H⁺) column density to account for the contribution from ionized gas to $A_V$. Given the uncertainties in the assumptions used to derive all quantities involved, we find a reasonable correspondence between N(H⁺) + N(H⁰) + 2N(H₂) and $A_V$, with a scaling that corresponds to the DGR of the LMC and SMC discussed in Section 4.2.

Analytic theories for the atomic-to-molecular transition have been presented by Sternberg et al. (2014) and McKee & Krumholz (2010) (see also Krumholz et al. 2008, 2009). Their work provides expressions for the molecular fraction as a function of total hydrogen column density (Equation (97) in Sternberg et al. 2014 and Equation (93) in McKee & Krumholz 2010) that can be used for comparison with our data. Both models assume a value of the strength of the FUV radiation field and that the gas metallicity and DGR scale by the same factor, $Z'_d$. These two parameters determine the total H column density at which the transition from H⁰ to H₂ takes place, with higher column densities when the metallicity is reduced and/or the FUV radiation field is increased. Hughes et al. (2010) studied the strength of the FUV field across the LMC using the dust temperature maps presented by Bernard et al. (2008). They found an average value of $\chi_0$ = 1.3, in units of the local FUV radiation field calculated by Draine (1978), which is similar to that found in the solar neighborhood. In the SMC, however, the strength of the FUV field is found to be larger than that in the solar neighborhood and the LMC (e.g., Lequeux 1989). Li & Draine (2002) used IRAS, DIRBE, and ISO observations to constrain a dust model of the SMC in which their best fit for the strength of the FUV field corresponds to $\chi_0$ = 1.9. We therefore adopted $\chi_0$ = 1.3 and 1.9 for the LMC and SMC, respectively. Note that because
<figure><img src="image_16.png" /><figcaption>**Figure 12.** Upper panels: column densities of H�0 and H₂ as a function of the total hydrogen column density, N(H⁰) + 2N(H₂), for LOSs in the LMC (left) and SMC (right). Lower panels: column densities of H⁺ and H⁰ decomposed in CNM and WNM gas as a function of the total hydrogen column density for LOSs in the LMC (left) and SMC (right). In all panels, the dashed-dotted lines indicate the predictions for N(H⁰), and the dashed lines indicate the predictions for N(H₂) from the models of Sternberg et al. (2014) and McKee & Krumholz (2010) for Z'_d = 0.26 and χ₀ = 1.3 for the LMC and Z'_d = 0.1 and χ₀ = 1.9 for the SMC (see Section 5.1).</figcaption></figure>

we have adopted a single value of the FUV radiation field in our model calculations, we do not expect an exact match between model and derived column densities, as the FUV field is likely to vary among our LOSs. Both Hughes et al. (2010) and Li & Draine (2002) show variations of a factor of ~2 in the derived FUV radiation field in the LMC and SMC, respectively. Because the model-predicted column density at which the H⁰-to-H₂ transition takes place is proportional to the FUV radiation field, variations in this quantity among our sample should result in a factor of ~2 of scatter around the model-predicted column density. This predicted scatter is consistent with that seen in our comparison between model and observed column densities (see Figures 12, 15, and 16). For the metallicity scaling, we adopted the scaling on the DGR rather than that of the carbon abundance, which is Z'_d = 0.26 for the LMC and 0.1 for the SMC. The use of the scaling on the DGR is justified because the Sternberg et al. (2014) and McKee & Krumholz (2010) models focus on describing the chemical

structure of the cloud (the H⁰-to-H₂ transition), which depends on the absorption of photons and the formation rate of H₂ on dust grains. These models, however, do not consider the thermal structure of the gas, which depends on the abundance of gas coolants. Thus, the most relevant parameter in these models is the DGR, rather than the gas-phase abundances of elements.

Sternberg et al. (2014; see their Figure 14) differed from McKee & Krumholz (2010) in their consideration of shielding by dust associated with both H⁰ and H₂, instead of only shielding by dust associated with H⁰ (McKee & Krumholz 2010), and their assumption of a larger dust absorption cross section ($\sigma_g = 1.9 \times 10^{-21} Z' \text{ cm}^{-2}$ versus $1.0 \times 10^{-21} Z' \text{ cm}^{-2}$). Additionally, Sternberg et al. (2014) assumed a cloud geometry of a slab illuminated from both sides instead of the spherical geometry assumed by McKee & Krumholz (2010). They showed that the differences between the two models are mostly explained by their consideration of absorption by dust associated with H₂ and the larger dust absorption cross section they assumed.
<figure><img src="image_17.png" /><figcaption>**Figure 13.** Column density of molecular hydrogen as a function of the total hydrogen column density, $N(H^0) + 2N(H_2)$, for the values derived here and those observed in UV absorption for $H_2$ and in optical absorption for $H^0$ compiled by Welty et al. (2012). The straight line represents a molecular fraction of $f(H_2) = 1$. We assume uncertainties of 10% for $N(H^0)$ and 20% for $N(H_2)$ in the Welty et al. (2012) data. We also show the predictions from the models of Sternberg et al. (2014) and McKee & Krumholz (2010) for $Z'_d = 0.26$ and $\chi_0 = 1.3$ for the LMC and $Z'_d = 0.1$ and $\chi_0 = 1.9$ for the SMC (see Section 5.1).</figcaption></figure>

In the top panels of Figure 12, we show the column densities of $H^0$ and $H_2$ that we have derived as a function of the $H^0 + 2H_2$ column density. We also show the $H^0$ and $H_2$ column densities predicted by McKee & Krumholz (2010) and Sternberg et al. (2014) as gray dashed and blue solid lines, respectively. We find that the Sternberg et al. (2014) model is in good agreement with the derived $H^0$ and $H_2$ column densities. The McKee & Krumholz (2010) model tends to predict lower $H_2$ column densities for a given value of $N(H^0) + 2N(H_2)$. Note that the column densities predicted by the Sternberg et al. (2014) and McKee & Krumholz (2010) models are functions of the shielding column density, which is dominated by the cold neutral medium gas in the case of atomic gas. Thus, their predicted $H^0$ column densities correspond to that of the CNM only, while the derived $H^0$ column densities presented in the top panel of Figure 12 correspond to the total $H^0$ along the LOS, which includes the contribution from both the WNM and the CNM.

In the bottom panels of Figure 12, we show a decomposition of the $H^0$ column density between WNM and CNM gas (see Section 3.2) and the contribution from $H^+$ (Section 3.1). The Sternberg et al. (2014) and McKee & Krumholz (2010) models tend to overestimate the CNM column densities. It is, however, possible that we have underestimated the CNM column density, as it can emit [C II] emission that is below the sensitivity of our observations. Most of the $H^0$ in our sample is in the form of WNM. We estimated that the CNM typically represents 28% and 14% of the total $H^0$ column density in the LMC and SMC, respectively. The CNM fractions in the LMC and SMC are in good agreement with those derived in H I absorption studies by Dickey et al. (2000) using H I absorption against background continuum sources (33% for the LMC and 13% for the SMC; Dickey et al. 2000). Note that the CNM fractions derived by Dickey et al. (2000) are proportional to the assumed kinetic temperature of the gas, which in their case is

55 K. While the CNM fraction in the LMC is somewhat smaller than that found in the solar neighborhood (40%; Heiles & Troland 2003; Pineda et al. 2013), this fraction is significantly smaller in the SMC. Dickey et al. (2000) argued that the low CNM fraction in the SMC is a result of the lower abundance of gas coolants, while the reduction of the photoelectric (PE) heating efficiency due to the lower DGR is offset by the enhanced FUV radiation field in the SMC (Lequeux 1989; Li & Draine 2002), resulting in warmer atomic gas. The $H^+$ component plays a small role in the derived column densities.

We found that the thresholds where $H_2$ formation becomes important are ~$10^{21.5}$ cm$^{-2}$ for the LMC and ~$10^{22}$ cm$^{-2}$ for the SMC. These correspond to 35 and 110 $M_{\odot}$ pc$^{-2}$ for the LMC and SMC, respectively, including the contribution from He. These values are higher than the Milky Way value of $10 M_{\odot}$ pc$^{-2}$ (Krumholz et al. 2008) and reflect the lower DGR of the LMC and SMC that reduces the rate of $H_2$ formation for a given total hydrogen column density. Thus, a larger hydrogen column is needed for $H_2$ self-shielding to balance $H_2$ photodestruction, allowing the gas to transition to a mostly molecular phase in a low-metallicity environment.

In Figure 13, we show the $H_2$ column densities we derived together with those compiled for the LMC and SMC by Welty et al. (2012) as a function of $N(H^0)+2N(H_2)$. For clarity, we exclude upper limits in the Welty et al. (2012) data. We also include the predictions of the model by Sternberg et al. (2014) and McKee & Krumholz (2010).<fnref n="19" /> In the LMC, the Sternberg et al. (2014) model is in good agreement with both the $H_2$ column densities derived here and most of those compiled by Welty et al. (2012). In the SMC, however, the Sternberg et al.

<footnote n="19">¹⁹ Note that Welty et al. (2012) also compared their data set with that of McKee & Krumholz (2010). The predictions plotted in their Figure 17 differ from those shown here, as they assumed different metallicity/DGR scaling factors of 0.5 and 0.2 for the LMC and SMC, respectively, and a FUV field of unity is implicitly assumed for both galaxies.</footnote>
<figure><img src="image_18.png" /><figcaption>Figure 14. Molecular fraction (Equation (12)) as a function of total hydrogen for the LMC and SMC. The lines are the predictions from the models of Sternberg et al. (2014) and McKee & Krumholz (2010) for $Z_d' = 0.26$ and $\chi_0 = 1.3$ for the LMC and $Z_d' = 0.1$ and $\chi_0 = 1.9$ for the SMC (see Section 5.1). The filled circles and squares correspond to LOSs in which we derived the H$_2$ column density from $N(C^+)$ only, while the open circles and squares represent LOSs in which $C^0$ and CO also contribute to the total carbon column density. We also include molecular fractions calculated from the sample of column densities observed in absorption in the FUV for H$_2$ and in the optical for H I compiled by Welty et al. (2012). For clarity, we do not show the error bars in the Welty et al. (2012) data. The uncertainties in the Welty et al. (2012) data are 30% for the total hydrogen column density and 36% for the molecular fraction based on the assumed uncertainties of 10% for $N(H^0)$ and 20% for $N(H_2)$.</figcaption></figure>

(2014) and McKee & Krumholz (2010) models are unable to reproduce the data points compiled by Welty et al. (2012) for $N(H_2) \le 10^{20} \text{ cm}^{-3}$. A larger metallicity factor and/or a lower FUV radiation field (e.g., $Z' = 0.2$ and $\chi_0 = 1$ for the Sternberg et al. 2014 model) would be required for the theoretical models to reproduce the data points. We see that the column densities derived from [C II], [C I], and CO provide a connection between the diffuse H$_2$ detected by FUV absorption and the dense H$_2$ in well-shielded regions. This connection is better illustrated in Figure 14, where we show the molecular

fraction,

$$f(H_2) = \frac{2N(H_2)}{N(H^0) + 2N(H_2)}, \quad (12)$$

calculated for both our sample and that of Welty et al. (2012) as a function of the total hydrogen column density. We again include the predictions from the models of Sternberg et al. (2014) and McKee & Krumholz (2010). We see that the FUV absorption observations mostly trace $f(H_2) \le 0.2$, while our determination using [C II], [C I], and CO observations traces the $0.1 \le f(H_2) \le 1$ range. In Figure 14, we also separate LOSs in which we derived the H$_2$ column density only from $N(C^+)$ from those in which $C^0$ and CO also contribute to the total carbon column density. We see that for molecular fractions $f(H_2) \gtrsim 0.45$, the $C^0$, $^{12}$CO, and $^{13}$CO column densities are large enough for the detection of [C I], $^{12}$CO, and $^{13}$CO lines. The [C II] line alone traces a larger range in the molecular fraction down to 0.1. This result suggests that [C II] observations, together with those of [C I] and CO, are important tools for the study of the transition from atomic to molecular gas in the ISM of galaxies.

## 5.2. The C$^+/C^0$/CO Transition

Another critical transition in interstellar clouds is that from ionized to neutral carbon and to the main molecular form of carbon, CO. The $C^+/C^0$/CO layer is the interface between the diffuse CO-dark H$_2$ gas and the dense molecular gas in which star formation takes place. Therefore, it has an important role in the evolution of molecular clouds and star formation. The formation of CO in the gas phase is balanced by photodestruction by FUV photons. The $C^+/C^0$/CO transition can take place if there is sufficient dust column density to efficiently shield the FUV photons that destroy CO. In low-metallicity environments, the lower DGR pushes the $C^+/C^0$/CO layer to higher hydrogen column densities. This effect has important implications for the ability of [C II], [C I], and CO to trace the molecular hydrogen mass in clouds.

In Figure 15, we show a comparison between the H$_2$ column density and the $C^+$, $C^0$, and CO column densities associated with H$_2$ derived in our sample in the LMC and SMC. The gray lines represent the total gas-phase carbon column density, $N(C^+)_{H_2} + N(C^0) + N(CO)$, that results by multiplying the H$_2$ column density by the [C]/[H$_2$] abundance ratio of the LMC and SMC (Section 3). Ionized carbon is the dominant form of gas-phase carbon that is associated with molecular gas in our sample. Both neutral carbon and CO represent a lower fraction of the total carbon along the LOS. The low column densities of CO and [C I] in our sample suggest that a large fraction of the gas column is associated with diffuse H$_2$ gas.

We were able to determine $C^0$ and CO column densities above $N(H_2) \sim 10^{21} \text{ cm}^{-2}$ in the LMC and $\sim 10^{21.5} \text{ cm}^{-2}$ in the SMC. The difference in these thresholds is likely influenced by the difference in metallicity and DGR between the LMC and SMC. As we discussed in the previous section, the $C^0$ and CO column densities have a measurable contribution to the total carbon abundance associated with H$_2$ for $f(H_2) \simeq 0.45$ in both the LMC and SMC.

In Figure 16, we show the fraction of $C^+$, $C^0$, and CO with respect to the total gas-phase carbon ($C^+ + C^0 + CO$) abundance associated with H$_2$ gas as a function of $N(H_2)$. Ionized
<figure><img src="image_19.png" /><figcaption>**Figure 15.** The C+, C°, and CO column densities associated with H₂ gas as a function of the H₂ column density derived in our sample in the LMC (top) and SMC (bottom). The gray lines represent the total gas-phase carbon column density, N(C⁺) + N(C°) + N(CO), that results by multiplying the H₂ column density by the [Cl]/[H₂] abundance ratio of 1.56 × 10⁻⁴ for the LMC and 2.8 × 10⁻⁵ for the SMC (Section 3).</figcaption></figure>

carbon is the dominant form of carbon in our LOS in the LMC and SMC. On average, C⁺, C°, and CO represent 89%, 9%, and 10% of the gas-phase carbon associated with molecular gas in the LMC and 77%, 6%, and 17% in the SMC, respectively.

### 5.3. The CO-to-H₂ Conversion Factor

The mass of molecular clouds in the Milky Way and external galaxies is often derived using observations of the ¹²CO J = 1 → 0 line together with an empirically derived CO-to-H₂ conversion factor, $X_{CO} \equiv \frac{N(\text{H}_2)}{I_{CO}}$ (Bolatto et al. 2013 and references therein). The X_{CO} factor is calibrated in the Milky Way as 2 × 10²⁰ cm⁻² (K km s⁻¹)⁻¹, but it is known to vary with environmental conditions—in particular with metallicity, which increases its value—as fainter CO emission per unit column density is observed (e.g., Rubio et al. 1993; Leroy et al. 2007; Hughes et al. 2010; Muller et al. 2010; Roman-Duval et al. 2014). In low-metallicity environments, the reduced DGR

results in enhanced photodestruction of CO, reducing the volume of H₂ that is traced by CO. However, the calibration in the Galaxy accounts for regions where CO is not the dominant form of gas-phase carbon (Pineda et al. 2010). With our estimate of N(H₂) and our ¹²CO J = 1 → 0 observations, we can study the variation of X_{CO} as a function of metallicity in the LMC and SMC.

In the left panel of Figure 17, we show the X_{CO} conversion factor as a function of the total H₂ column density in the LMC and SMC. We calculated X_{CO} by dividing the H₂ column density (Section 5.1) by the CO integrated intensity, both of which correspond to their total value along the LOS. In the figure, we show data points for all LOSs where CO was detected, including those where neither higher −J nor isotopic transitions were detected (Section 4.4.1). In these LOSs, the H₂ column density was derived from N(C⁺). While showing a large scatter, the values of X_{CO} in the LMC are consistent with and the values in the SMC are somewhat lower than those derived by Roman-Duval et al. (2014) of 6 × 10²⁰ and 4 × 10²¹ cm⁻² (K km s⁻¹)⁻¹ for the LMC and SMC, respectively, which are similar to values commonly found in the literature (e.g., Rubio et al. 1993; Leroy et al. 2007; Pineda et al. 2009; Hughes et al. 2010; Muller et al. 2010). We find a mean value of X_{CO} in the LMC of 9.1 × 10²⁰ cm⁻² (K km s⁻¹)⁻¹ and in the SMC of 2.2 × 10²¹ cm⁻² (K km s⁻¹)⁻¹. The correlation between X_{CO} and N(H₂) in Figure 17 is a result of the intensity of CO having a much smaller variation compared with the H₂ column density in our sample.

As mentioned above, the reduced DGR in the low-metallicity Magellanic clouds causes an enhanced photodestruction of CO, which results in CO tracing a smaller volume of the H₂ gas. If this is the case, the ¹²CO J = 1 → 0 intensity is reduced by a filling factor, which in turn results in a larger X_{CO} compared to that of a cloud with solar metallicity. We can test this effect in our sample, as we derived the filling factor of the CO-emitting gas as part of the excitation analysis in Section 4.4.1. For LOSs detected in ¹²CO J = 1 → 0 but where the H₂ column density was derived from N(C⁺), we assumed a filling factor of unity. In the right panel of Figure 17, we show X_{CO} as a function of the total H₂ column density in our sample, but this time we multiply X_{CO} by the filling factor derived in the excitation analysis, effectively correcting the I_{CO} intensity for filling factor effects.

While applying this correction increases the scatter in the observed values, we see that the values of X_{CO} in the LMC and SMC are closer to the value of X_{CO} in the Milky Way, where the filling factor is closer to unity. The mean values of X_{CO} in the LMC and SMC are reduced to 2.9 × 10²⁰ and 7.6 × 10²⁰ cm⁻² (K km s⁻¹)⁻¹, respectively, which are closer to the value derived in the Milky Way of 2 × 10²⁰ cm⁻² (K km s⁻¹)⁻¹ (Bolatto et al. 2013). The reduced values of X_{CO} resulting from applying a filling factor seem to confirm that the effect of metallicity on the value of X_{CO} is to reduce the filling factor of molecular gas traced by CO, resulting in a lower ¹²CO J = 1 → 0 intensity that in turn enhances the X_{CO} factor.

### 5.4. Photoelectric Heating and [CII] Cooling

PE heating is the main heating mechanism in atomic and diffuse molecular clouds. In this process, the gas is heated by energetic electrons that are expelled from dust grains after they absorb FUV photons (Spitzer 1948). The efficiency with which
<figure><img src="image_20.png" /><figcaption>Figure 16. Fraction of C⁺, C⁰, and CO with respect to the total carbon (C⁺ + C⁰ + CO) column density associated with H₂ gas as a function of the H₂ column density in the LMC (left) and SMC (right).</figcaption></figure>

PE heating works is dependent on the strength of the FUV radiation field, the gas temperature, and the electron density, as well as on grain charging, which depends on the type of dust grain involved.

The efficiency of PE heating can be studied by comparing the total IR intensity (tracing the energy absorbed by dust) and the [C II] emission, which is the main cooling line in regions where PE heating dominates. Such a comparison has been done in the LMC at an angular resolution of 14′.9 (217 pc spatial resolution) by Rubin et al. (2009), resulting in a PE heating efficiency that varies by a factor of 1.4 between diffuse and bright star-forming regions. They find that, overall in the LMC, [C II] constitutes 1.32% of the FIR luminosity. Similar contributions from [C II] to the observed FIR intensities were found by Israel & Maloney (2011) in several star-forming regions in the LMC and SMC with the Kuiper Airborne Observatory at a resolution of 55" (13.3 pc for the LMC and 16.2 pc for the SMC) and in a larger sample of low-metallicity dwarf galaxies (Cormier et al. 2015). The [C II]/FIR ratios found in the Magellanic clouds and other dwarf galaxies are higher than the 0.1%-0.2% value found in the Milky Way (Wright et al. 1991). This difference in the [C II]/FIR ratio has been attributed to the lower DGR in low-metallicity environments, which results in lower FIR intensity and an increased volume of C⁺-emitting regions. In the following, we compare our observations of [C II] in the LMC and SMC with that of the total IR emission derived from a combination of the *Spitzer* (Meixner et al. 2006) and *Herschel* (Meixner et al. 2013) dust continuum data.

Galametz et al. (2013) compared the resolved galaxy observations of several *Herschel* and *Spitzer* bands with dust SED models from Draine & Li (2007) to derive empirical relationships between the specific intensity at different bands and the TIR intensity integrated between 3 and 1100 µm. In irregular galaxies (such as the LMC and SMC), the total IR is related to the FIR intensity, defined as the intensity integrated between 42 and 122 µm, as TIR/FIR ≈ 2 (Hunter et al. 2001).

We adopted the calibration for the TIR intensity (in units of erg cm⁻² s⁻¹ sr⁻¹) presented by Galametz et al. (2013) that uses the *Spitzer* 24 µm and *Herschel* 100 µm specific intensities (in units of erg cm⁻² s⁻¹ sr⁻¹ Hz⁻¹),

$$I^{\text{TIR}} = 2.421\nu I_{\nu,24} + 1.410\nu I_{\nu,100}, \quad (13)$$

as it ensures that the resulting TIR intensity can be smoothed to the 12" resolution of the *Herschel* [C II] observations. Galametz et al. (2013) found that this calibration is a good representation of the TIR resulting from full SED modeling of a large number of galaxies observed by *Herschel* and *Spitzer*. At 12" resolution, our comparison between the TIR and [C II] intensities is at 3 and 3.5 pc scales in the LMC and SMC, respectively. The resulting TIR maps have a typical uncertainty of $2 \times 10^{-4}$ erg cm⁻² s⁻¹ sr⁻¹.

In Figure 18, we show the [C II] intensity as a function of the TIR estimated in the LMC and SMC. In the plot, we also include lines of constant [C II]/TIR going from 0.1% to 3.16% in steps of 0.25 dex, as well as the fit of the [C II] and TIR relationship obtained by Rubin et al. (2009). We find that [C II] intensities in both the LMC and SMC have a similar relationship with the total IR emission despite their different metallicity and DGR. On average, [C II] emits 0.6% and 0.8% of the TIR intensity in the LMC and SMC, respectively. In terms of the FIR intensity, these averages correspond to 1.2% and 1.6%; thus, they are in good agreement with the 1.32% found by Rubin et al. (2009) over the whole LMC and the values found in star-forming regions of the LMC and SMC by Israel & Maloney (2011).

Dense PDRs in the LMC tend to have lower [C II]/TIR ratios, with [C II] emitting about 0.17% of the TIR. A similar variation from diffuse to bright regions was suggested by Rubin et al. (2009). In Figure 18, we include data points that correspond to the 30 Doradus and N159W massive star-forming regions in the LMC ([C II] data from Boreiko & Betz 1991 at 43" resolution) that, together with the warm and dense PDRs in our sample, are in agreement with the lower
<figure><img src="image_21.png" /><figcaption>**Figure 17.** Left: the $X_{CO}$ conversion factor as a function of the total H$_2$ column density derived in Section 5.1 in our LMC and SMC sample. The horizontal lines correspond to $X_{CO} = 2 \times 10^{20} \text{ cm}^{-2} (\text{K km s}^{-1})^{-1}$ compiled from several estimates in the Galaxy by Bolatto et al. (2013) and $6 \times 10^{20}$ and $4 \times 10^{21} \text{ cm}^{-2} (\text{K km s}^{-1})^{-1}$ estimated by Roman-Duval et al. (2014) in the LMC and SMC, respectively. Right: the $X_{CO}$ conversion factor as a function of the total H$_2$ column density corrected by the filling factor of the CO-emitting gas derived in our sample (see Section 4.4.1).</figcaption></figure>

<figure><img src="image_22.png" /><figcaption>**Figure 18.** The [C II] intensity as a function of the total IR intensity in our LMC and SMC sample. We also include values for the 30 Doradus and N159W massive star-forming regions in the LMC, with the [C II] data taken from Boreiko & Betz (1991). The dashed lines represent constant values of PE efficiency ranging from 0.1% to 3.16% in steps of 0.25 dex. The solid line corresponds to the fit of the [C II]-TIR relationship obtained by Rubin et al. (2009).</figcaption></figure>

[C II]/TIR ratios at high [C II] and TIR intensities suggested by the fit presented by Rubin et al. (2009). The reduced [C II]/TIR ratio in dense PDRs can be understood as the combined effect of high volume densities and the high temperatures resulting from large FUV radiation fields. With increasing volume density and temperature, the excitation of the $^2\text{P}_{3/2}$ level becomes a weak function of these quantities, and the [C II] intensity is only a function of column density (Goldsmith et al. 2012). The TIR intensity, which is a function of column

density and dust temperature, is unaffected by the larger volume densities and FUV radiation fields.

### 5.5. Comparison with FUSE LOSs

As mentioned in Section 2, our sample contains seven LOSs in the LMC and three in the SMC that coincide with stars that show UV and optical absorption features that are used to determine the H$_2$ and H$^0$ column densities of their foreground gas (e.g., Tumlinson et al. 2002; Cartledge et al. 2005; Welty et al. 2012). In this subsample, we detected [C II] emission in five LOSs in the LMC and one in the SMC, with one LOS in the LMC also detected in [C I] and CO. Because the column densities determined in our analysis correspond to the average column density within the beam of our observations, and those from UV and optical absorption correspond to column densities in a pencil beam, a comparison between these column densities will be sensitive to substructure in the ISM of the LMC and SMC. Another difference between these methods for determining column densities is that the column densities in our analysis correspond to the gas along the full sight line through the LMC or SMC, while the UV and optical absorption only corresponds to the gas that is in the foreground of the target stars. Therefore, a comparison between the column densities determined with these different methods can also give us some information on the structure of the ISM along the LOS in the LMC and SMC.

In Figure 19, we compare the H$^0$ column densities derived using H I 21 cm observations ($N(H^0)_{21\text{cm}}$) and those derived using optical absorption of the Ly$\alpha$ line ($N(H^0)_{Ly\alpha}$); we also compare the H$_2$ column densities derived in our analysis ($N(H_2)_{HIFI}$) and those derived from UV absorption observations ($N(H_2)_{UV}$). The H$^0$ and H$_2$ column densities derived from absorption observations were compiled by Welty et al. (2012). We see that some LOSs are in reasonable agreement between the column densities derived using H I 21 cm observations and those derived using Ly$\alpha$ absorption, with some having slightly larger $N(H^0)_{Ly\alpha}$ than that expected from H I 21 cm observations.
<figure><img src="image_23.png" /><figcaption>**Figure 19.** Left: comparison of the H�0 column densities derived using 21 cm observations and those derived using optical absorption of the Lyα line. Right: comparison of the H₂ column densities derived in our analysis and those derived from UV absorption observations. The H⁰ and H₂ column densities derived from absorption measurements were compiled by Welty et al. (2012).</figcaption></figure>

However, there are several LOSs that have $N(H^0)_{Lyα}$ that is significantly lower than $N(H^0)_{21cm}$. We also see that the H₂ column densities in our sample are systematically larger than those derived with FUSE. A similar result is seen in the comparison of atomic hydrogen column densities in a larger sample presented by Welty et al. (2012) in the LMC and SMC. As explained by these authors, $N(H^0)_{Lyα} < N(H^0)_{21cm}$ could be a result of having the bulk of the atomic gas behind the target stars. Alternatively, this effect can be explained, as can the case in which $N(H^0)_{Lyα} > N(H^0)_{21cm}$, as a result of substructure within the 21 cm beam where the pencil beam of the optical absorption goes through a maximum or minimum in the gas distribution that has a much different column density compared with the average within the beam. The first explanation would require H₂ to be more extended along the LOSs than H⁰, which contradicts the spectral information that shows that H I is much more extended in velocity compared with [C II], [C I], and CO. A more likely explanation for our data is the existence of substructure within the beam of the observations. The larger difference for the H₂ column densities compared to that for the H⁰ column densities could be a reflection of the different structure of the atomic and molecular gas, with the latter concentrated into smaller clumps.

## 6. Conclusions

In this paper, we present a survey of the [C II], [C I], and CO emission observed with high-velocity resolution in the LMC and SMC. The paper aims to characterize the transition from diffuse atomic to dense molecular clouds in the low-metallicity environments of the Magellanic clouds. Our sample was selected based on the H I, CO, and dust continuum emission and includes regions in different stages of the transition from diffuse regions to dense PDRs associated with star formation. Our results can be summarized as follows:

1. We determined the contribution from different phases of the ISM to the observed [C II] emission in our sample. In LOSs associated with warm and dense PDRs, the [C II] emission from this ISM component tends to dominate.

The contribution from diffuse CO-dark H₂ to the observed [C II] shows a large scatter, ranging from ~10% to ~80%. This large scatter could be a reflection of clouds at different stages in the transition from diffuse to dense molecular gas, with different clouds having varying CO-dark H₂ fractions. We found that ionized and atomic gas have typically smaller contributions to the observed [C II] emission (Section 3).

2. Using LOSs where only H I and [C II] emission were detected and a derivation of the visual extinction from dust continuum emission, we determined a typical thermal pressure of the diffuse ISM to be $p_{th}/k_B = 3.4 \times 10^4$ K cm⁻³ in the LMC and $1 \times 10^5$ K cm⁻³ in the SMC (Section 4.3).

3. We used the [C I] and CO observations to determine the column density, kinetic temperature, and volume density of the CO- and [C I]-emitting gas. We found average values and standard deviations in the CO-emitting layer of 44 ± 19 K in the LMC and 50 ± 20 K in the SMC for the kinetic temperature, $10^{16.7 \pm 0.5}$ cm⁻² in the LMC and $10^{16.6 \pm 0.3}$ cm⁻² in the SMC for the beam-averaged CO column density, $10^{5.1 \pm 0.9}$ cm⁻³ in the LMC and $10^{4.0 \pm 0.3}$ cm⁻³ in the SMC for the H₂ volume density, and (0.2 ± 0.2) in the LMC and (0.20 ± 0.3) in the SMC for the filling factor. In the [C I]-emitting layer, we found 91 ± 60 K in the LMC and 73 ± 62 K in the SMC for the kinetic temperature, $10^{16.4 \pm 0.4}$ cm⁻² in the LMC and $10^{16.2 \pm 0.3}$ cm⁻² in the SMC for the beam-averaged C⁰ column density, and $10^{4.8 \pm 1.4}$ cm⁻³ in the LMC and $10^{3.5 \pm 0.5}$ cm⁻³ in the SMC for the H₂ volume density. This information, combined with the [C II] observations, allowed us to calculate the column densities of CO-dark H₂ gas in our sample (Section 4.4).

4. We studied the transition between atomic and molecular gas in the LMC and SMC by comparing the H₂ column densities of CO-dark H₂ and CO-traced H₂ with the total hydrogen, $N(H^0) + 2N(H_2)$, column density in our sample. We found reasonable agreement between our
observations and theoretical models of the H⁰-to-H₂ transition (Section 5.1).

5. We found that most of the observed H I is in the form of WNM. We estimated that CNM represents 28% and 14% of the total H⁰ column density in the LMC and SMC, respectively. The value in the LMC is similar to that in the Milky Way (40%; Heiles & Troland 2003; Pineda et al. 2013), while the value in the SMC represents a smaller fraction (Section 5.1).

6. We found that the H₂ column densities determined using [C II], [C I], and CO observations in our survey trace molecular fractions in the 0.1 $\le f(H_2) \le 1$ range. In contrast, UV absorption observations mostly trace $f(H_2) \le 0.2$. The C⁰ and CO column densities have a measurable contribution to the total gas-phase carbon column density for molecular fractions $f(H_2) \ge 0.45$, while [C II] alone traces a larger range in the molecular fraction, down to 0.1 (Section 5.1).

7. Studying the C⁺/C⁰/CO transitions in our sample in the LMC and SMC reveals that most of the molecular gas in our sample is traced by [C II]. Both column densities of neutral carbon and CO represent a lower fraction of the total carbon associated with molecular gas along the LOS. On average, C⁺, C⁰, and CO represent 89%, 9%, and 10% of the gas-phase carbon in the LMC and 77%, 6%, and 17% in the SMC, respectively. The [C I] and CO are detected above N(H₂) ~ $10^{21.5}$ cm⁻² in the LMC and ~$10^{21.8}$ cm⁻² in the SMC. The difference in the thresholds for [C I] and CO detection is possibly a result of the metallicity and DGR difference between the LMC and SMC (Section 5.2).

8. We found a mean value of the XCO conversion factor in the LMC of $9.1 \times 10^{20}$ cm⁻² (K km s⁻¹)⁻¹ and in the SMC of $2.2 \times 10^{21}$ cm⁻² (K km s⁻¹)⁻¹, both of which are larger than the value of $2 \times 10^{20}$ cm⁻² (K km s⁻¹)⁻¹ in the Milky Way. When applying a filling factor correction to the CO emission, we saw that the values of XCO in the LMC and SMC become closer to the value of XCO in the Milky Way, where the filling factor is close to unity. This result suggests that the effect of metallicity in the value of XCO is to reduce the filling factor of molecular gas traced by CO, which leads to a lower ¹²CO J = 1 → 0 intensity that in turn results in an enhanced XCO factor (Section 5.3).

9. We found that most of the LOSs in our sample are consistent with a linear relationship between the [C II] and TIR emission. The [C II] emission represents about 1% of the TIR in both the LMC and SMC, despite their difference in metallicity and DGR. This [C II]/TIR fraction is consistent with previous determinations using the [C II] map of the LMC observed with the BICE balloon (Section 5.4).

10. We compared the H⁰ and H₂ column densities derived in our analysis and those derived using optical and UV absorption in a subsample of our survey. We found significant discrepancies between the column density observations that can be explained by substructure within the beam used for our observations (Section 5.5).

In conclusion, our results show that [C II], [C I], and CO observations are important tools for characterizing the transition from atomic to molecular clouds in external galaxies. Future large-scale mapping of the LMC and SMC in these

spectral lines with current/future balloon, airborne (SOFIA, GUSTO), and space observatories will make significant progress in our understanding of the evolution of the ISM, star formation, and galaxy evolution.

We would like to thank Drs. Julia Roman-Duval and Cheoljong Lee for providing their dust continuum data sets and Dr. Franck Le Petit for his help in comparing our H I and H₂ data sets with theoretical predictions. We also thank an anonymous referee for a number of useful comments that significantly improved the manuscript. This research was conducted at the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We thank the staffs of the ESA and NASA *Herschel* Science Centers for their help. © 2016 California Institute of Technology. U.S. government sponsorship acknowledged.

# Appendix
## Beam Dilution Correction

### A.1. Absolute Intensities

Our survey contains a set of multiwavelength observations with different angular resolutions. When possible, our analysis was made with data at matching resolutions, but there are cases in which we had to rely on a combination of data with different angular resolutions. In order to quantify how beam dilution affects the results involving absolute intensities presented in this paper, we smoothed the 12" angular resolution 160 µm HERITAGE dust continuum maps of the LMC and SMC to the different angular resolutions of the data set used in our analysis (12", 17", 27", 33", 44", 48", and 60"). We assumed that the distribution of dust continuum emission is similar to that of the spectral lines used in our analysis. This assumption is motivated by the agreement between dust continuum and [C II] maps in star-forming regions in the LMC (see, e.g., Okada et al. 2015 and Galametz et al. 2013 for N159). We chose a reference angular resolution of 40", which is an intermediate value between 12" and 60". In Figure 20, we show the variation of the absolute 160 µm intensities for angular resolutions between 12" and 60" with respect to that at 40". Remarkably, most of our LOSs in the LMC and SMC show variations within ~40% in the 12"–60" range, even though the beam area varies by a factor of 25. There are, however, four exceptions—NT77, PDR3_NE, PDR2_NW, and LMC12_NW—where the intensities vary by factors of 6, 3.5, 2, and 1.7, respectively. In the SMC, AzV462, SMC_NE_1a, and SMC_LIRS36 have the largest variations, with factors of 2.3, 2, and 1.8, respectively. Note that we did not detect [C II] emission in AzV462.

We also studied the variations of the SHASSA Hα intensities for varying angular resolutions by using the 24 µm dust continuum maps of the LMC and SMC. The 24 µm continuum emission is a proxy for hot dust emission associated with H II regions. The angular resolution of the 24 µm images is 6". We convolved these images to 12" corresponding to [C II] observations and 48" corresponding to the Hα map. We then computed the ratio of the 24 µm intensities at 12" and 48" to determine a beam dilution correction factor for the Hα data, which we used to quantify the impact of beam dilution in our results. We found that the median flux variation between 12" and 48" is a factor of 1.15 in both the LMC and SMC. The
<figure><img src="image_24.png" /><figcaption>**Figure 20.** The 160 µm dust continuum intensity smoothed to different angular resolutions relative to that at 40" as a function of angular resolution for our sample in the LMC (left panel) and SMC (right panel).</figcaption></figure>

<figure><img src="image_25.png" /><figcaption>**Figure 21.** Fraction of the [C II] emission that we estimated arises from ionized gas, cold atomic gas, CO-dark H₂, and PDRs as a function of the observed [C II] emission in the LMC and SMC. All data points have been corrected by dilution effects.</figcaption></figure>

mean variations are factors of 1.6 in the LMC and 1.7 in the SMC. The differences between the median and mean factors are mostly driven by three LOSs in both the LMC and SMC that show larger variations in their fluxes. In the LMC, they are SK-67D2, PDR3_NE, and NT77, with fluxes varying by factors of 3.1, 3.4, and 8.3, respectively. (Note that SK-67D2 was not detected in [C II] in our survey.) In the SMC, the LOSs with the largest variations are SMC_LIRS36, SMC_NE_1a, and SMC_NE_4c_low, with fluxes varying by factors of 3.6, 3.7, and 5.2, respectively.

In Figure 21, we present the fraction of [C II] originating from different ISM phases as a function of the observed [C II] intensity in the case in which the intensities are corrected by beam dilution effects. All quantities involved in the calculation

of the contribution from different ISM phases to the observed [C II] emission have been corrected by a beam filling effect to correspond to a common angular resolution of 40". In general, beam dilution effects have a small impact on the fraction of [C II] emission originating from the different phases of the ISM.

## A.2. Line Ratios

Another instance in which we used a combination of pointed observations at different angular resolutions is the excitation analysis in Section 4.4, where we used line ratios of CO and [C I] lines to estimate the physical conditions of the CO- and [C I]-emitting gas. In the following, we describe a simple
model that we used to account for beam dilution effects in the line ratios. Because we assumed Gaussian sources, the model provides only a first-order correction to the observed line ratios, as the structure of the ISM is likely to be more complex.

The ratio of the antenna temperatures of two lines, $T_1^*$ and $T_2^*$, originating from a Gaussian source with FWHM size $\Theta_s$ when observed with the same FWHM beam size $\Theta_f$ centered in the source, is given by

$$ \frac{T_1^*}{T_2^*} = \frac{\frac{t_1 \Theta_s^2}{\Theta_f^2 + \Theta_s^2}}{\frac{t_2 \Theta_s^2}{\Theta_f^2 + \Theta_s^2}} = \frac{T_1}{T_2}, \qquad (14) $$

where $T_1$ and $T_2$ are the intrinsic peak antenna temperature of the source. When the source is convolved with Gaussians that have two different FWHM beam sizes $\Theta_a$ and $\Theta_b$, the integrated intensity ratio is given by

$$ \frac{T_1^*}{T_2^*} = \frac{T_1 \Theta_s^2 + \Theta_b^2}{T_2 \Theta_s^2 + \Theta_a^2}. \qquad (15) $$

Thus, the intrinsic ratio is related to the observed ratio as

$$ \frac{T_1}{T_2} = \frac{T_1^*}{T_2^{*f_{1,2}}}, \qquad (16) $$

where

$$ f_{1,2} \equiv \frac{\Theta_s^2 + \Theta_a^2}{\Theta_s^2 + \Theta_b^2}. \qquad (17) $$

The observed line ratio can therefore be corrected for beam dilution effects if an estimate of $f_{1,2}$ is available. We can estimate $f_{1,2}$ by comparing the line or continuum emission of the same source at different angular resolutions. In this case, $T_1 = T_2$ and, therefore, the dilution factor is given by the inverse of the ratio of the convolved peak intensities, $f_{1,2} = T_2^*/T_1^*$.

We used the 160 µm HERITAGE dust continuum map with 12" angular resolution to estimate $f_{1,2}$ for the sources in our survey. We smoothed the 160 µm map to the different angular resolutions of the spectral lines involved in our analysis (17", 27", 33", and 44") and calculated the dilution factors of each line ratio using Equation (17) by calculating the ratio of the 160 µm emission at the different angular resolution pairs. We used the derived correction factors to study the effect of beam dilution in the line ratios used in our analysis. Note that we did not correct the line ratios for beam dilution effects in our analysis, due to the uncertainties of whether dust continuum emission traces the distribution of gas. In Table 6, we present the effect of beam dilution in the line ratios that are calculated using observations at different angular resolutions. We find that these line ratios would typically vary by 10%. This small variation suggests that the observed structures are relatively extended at the resolution of our observations.

## References

* Alves, M. I. R., Calabretta, M., Davies, R. D., et al. 2015, MNRAS, 450, 2025
* Barinovs, G., van Hemert, M. C., Krems, R., & Dalgarno, A. 2005, ApJ, 620, 537

* Bernard, J.-P., Reach, W. T., Paradis, D., et al. 2008, AJ, 136, 919

* Bialy, S., & Sternberg, A. 2016, ApJ, 822, 83

* Blitz, L., & Rosolowsky, E. 2006, ApJ, 650, 933

* Bohlin, R. C., Savage, B. D., & Drake, J. F. 1978, ApJ, 224, 132

* Bolatto, A. D., Leroy, A. K., Jameson, K., et al. 2011, ApJ, 741, 12

* Bolatto, A. D., Wolfire, M., & Leroy, A. K. 2013, ARA&A, 51, 207

* Boreiko, R. T., & Betz, A. L. 1991, ApJL, 380, L27

* Boulares, A., & Cox, D. P. 1990, ApJ, 365, 544

* Carilli, C. L., & Walter, F. 2013, ARA&A, 51, 105

* Cartledge, S. I. B., Clayton, G. C., Gordon, K. D., et al. 2005, ApJ, 630, 355

* Chevance, M., Madden, S. C., Lebouteiller, V., et al. 2016, A&A, 590, A36

* Cohen, R. S., Dame, T. M., Garay, G., et al. 1988, ApJL, 331, L95

* Cormier, D., Madden, S. C., Lebouteiller, V., et al. 2015, A&A, 578, A53

* de Graauw, T., Helmich, F. P., Phillips, T. G., et al. 2010, A&A, 518, L6+

* Dickey, J. M., Mebold, U., Stanimirovic, S., & Staveley-Smith, L. 2000, ApJ, 536, 756

* Draine, B. T. 1978, ApJS, 36, 595

* Draine, B. T., & Li, A. 2007, ApJ, 657, 810

* Dufour, R. J., Shields, G. A., & Talbot, R. J., Jr. 1982, ApJ, 252, 461

* Dwek, E. 1998, ApJ, 501, 643

* Field, G. B., Goldsmith, D. W., & Habing, H. J. 1969, ApJL, 155, L149

* Fitzpatrick, E. L. 1985a, ApJ, 299, 219

* Fitzpatrick, E. L. 1985b, ApJS, 59, 77

* Galametz, M., Kennicutt, R. C., Calzetti, D., et al. 2013, MNRAS, 431, 1956

* Gaustad, J. E., McCullough, P. R., Rosing, W., & Van Buren, D. 2001, PASP, 113, 1326

* Gnedin, N. Y., Tassis, K., & Kravtsov, A. V. 2009, ApJ, 697, 55

* Goldsmith, P. F. 2013, ApJ, 774, 134

* Goldsmith, P. F., Langer, W. D., Pineda, J. L., & Velusamy, T. 2012, ApJS, 203, 13

* Goldsmith, P. F., Pineda, J. L., Langer, W. D., et al. 2016, ApJ, 824, 141

* Goldsmith, P. F., Yıldız, U. A., Langer, W. D., & Pineda, J. L. 2015, ApJ, 814, 133

* Gordon, K. D., Clayton, G. C., Misselt, K. A., Landolt, A. U., & Wolff, M. J. 2003, ApJ, 594, 279

* Gordon, K. D., Roman-Duval, J., Bot, C., et al. 2014, ApJ, 797, 85

* Grenier, I. A., Casandjian, J.-M., & Terrier, R. 2005, Sci, 307, 1292

* Güsten, R., Nyman, L. Å., Schilke, P., et al. 2006, A&A, 454, L13

* Haffner, L. M., Dettmar, R.-J., Beckman, J. E., et al. 2009, RvMP, 81, 969

* Heiles, C., & Troland, T. H. 2003, ApJ, 586, 1067

* Heyer, M., & Dame, T. M. 2015, ARA&A, 53, 583

* Higgins, D. R. 2011, PhD thesis, National Univ. Ireland Maynooth

* Higgins, R. D., & Kooi, J. W. 2009, Proc. SPIE, 7215, 721050L

* Hilditch, R. W., Howarth, I. D., & Harries, T. J. 2005, MNRAS, 357, 304

* Hughes, A., Wong, T., Ott, J., et al. 2010, MNRAS, 406, 2065

* Hunter, D. A., Kaufman, M., Hollenbach, D. J., et al. 2001, ApJ, 553, 121

* Israel, F. 2000, in Molecular Hydrogen in Space, ed. F. Combes & G. Pineau Des Forets (Cambridge: Cambridge Univ. Press), 293

* Israel, F. P. 1997, A&A, 328, 471

* Israel, F. P., & Maloney, P. R. 2011, A&A, 531, A19

* Israel, F. P., Maloney, P. R., Geis, N., et al. 1996, ApJ, 465, 738

* Jenkins, E. B., & Tripp, T. M. 2011, ApJ, 734, 65

* Kalberla, P. M. W., & Kerp, J. 2009, ARA&A, 47, 27

* Kim, S. 2006, PASP, 118, 94

* Kim, S., Staveley-Smith, L., Dopita, M. A., et al. 2003, ApJS, 148, 473

* Koornneef, J. 1982, A&A, 107, 247

* Krumholz, M. R., McKee, C. F., & Tumlinson, J. 2008, ApJ, 689, 865

* Krumholz, M. R., McKee, C. F., & Tumlinson, J. 2009, ApJ, 693, 216

* Ladd, N., Purcell, C., Wong, T., & Robertson, S. 2005, PASA, 22, 62

* Langer, W. D., Goldsmith, P. F., & Pineda, J. L. 2016, A&A, 590, A43

* Langer, W. D., Velusamy, T., Pineda, J. L., et al. 2010, A&A, 521, L17

* Langer, W. D., Velusamy, T., Pineda, J. L., Willacy, K., & Goldsmith, P. F. 2014, A&A, 561, A122

* Lee, C., Leroy, A. K., Schnee, S., et al. 2015, MNRAS, 450, 2708

* Lee, M.-Y., Madden, S., Lebouteiller, V., et al. 2016, A&A, 596, A85

* Lequeux, J. 1989, in Recent Developments of Magellanic Cloud Research, ed. K. S. de Boer, F. Spite, & G. Stasińska (Paris: Observatoire de Paris), 119

* Leroy, A., Bolatto, A., Stanimirovic, S., et al. 2007, ApJ, 658, 1027

* Li, A., & Draine, B. T. 2002, ApJ, 576, 762

* Madden, S. C., Poglitsch, A., Geis, N., Stacey, G. J., & Townes, C. H. 1997, ApJ, 483, 200

* Martin, N., Maurice, E., & Lequeux, J. 1989, A&A, 215, 219

* Marx-Zimmer, M., Herbstmeier, U., Dickey, J. M., et al. 2000, A&A, 354, 787

* Marx-Zimmer, M., Zimmer, F., Herbstmeier, U., & Dickey, J. M. 1999, in IAU Symp. 190, New Views of the Magellanic Clouds, ed. Y.-H. Chu et al., 124

* McKee, C. F. 1989, ApJ, 345, 782

* McKee, C. F., & Krumholz, M. R. 2010, ApJ, 709, 308

* Meixner, M., Gordon, K. D., Indebetouw, R., et al. 2006, AJ, 132, 2268

* Meixner, M., Panuzzo, P., Roman-Duval, J., et al. 2013, AJ, 146, 62
Minamidani, T., Mizuno, N., Mizuno, Y., et al. 2008, ApJS, 175, 485

Minamidani, T., Tanaka, T., Mizuno, Y., et al. 2011, AJ, 141, 73

Mochizuki, K., Nakagawa, T., Doi, Y., et al. 1994, ApJL, 430, L37

Muller, E., Ott, J., Hughes, A., et al. 2010, ApJ, 712, 1248

Okada, Y., Requena-Torres, M. A., Güsten, R., et al. 2015, A&A, 580, A54

Ostriker, E. C., McKee, C. F., & Leroy, A. K. 2010, ApJ, 721, 975

Ott, S., Bakker, J., Brumfitt, J., et al. 2006, in ASP Conf. Ser. 351, Astronomical Data Analysis Software and Systems XV, ed. C. Gabriel et al. (San Francisco, CA: ASP), 516

Paradis, D., Paladini, R., Noriega-Crespo, A., et al. 2011, ApJ, 735, 6

Parker, J. W., Garmany, C. D., Massey, P., & Walborn, N. R. 1992, AJ, 103, 1205

Pellegrini, E. W., Baldwin, J. A., & Ferland, G. J. 2011, ApJ, 738, 34

Pikel'Ner, S. B. 1968, SvA, 11, 737

Pilbratt, G. L., Riedinger, J. R., Passvogel, T., et al. 2010, A&A, 518, L1

Pineda, J. L., Goldsmith, P. F., Chapman, N., et al. 2010, ApJ, 721, 686

Pineda, J. L., Langer, W. D., Velusamy, T., & Goldsmith, P. F. 2013, A&A, 554, A103

Pineda, J. L., Mizuno, N., Röllig, M., et al. 2012, A&A, 544, A84

Pineda, J. L., Mizuno, N., Stutzki, J., et al. 2008, A&A, 482, 197

Pineda, J. L., Ott, J., Klein, U., et al. 2009, ApJ, 703, 736

Poglitsch, A., Krabbe, A., Madden, S. C., et al. 1995, ApJ, 454, 293

Requena-Torres, M. A., Israel, F. P., Okada, Y., et al. 2016, A&A, 589, A28

Reynolds, R. J. 1991, ApJL, 372, L17

Roman-Duval, J., Gordon, K. D., Meixner, M., et al. 2014, ApJ, 797, 86

Rubin, D., Hony, S., Madden, S. C., et al. 2009, A&A, 494, 647

Rubio, M., Garay, G., Montani, J., & Thaddeus, P. 1991, ApJ, 368, 173

Rubio, M., Lequeux, J., & Boulanger, F. 1993, A&A, 271, 9

Sandstrom, K. M., Bolatto, A. D., Draine, B. T., Bot, C., & Stanimirović, S. 2010, ApJ, 715, 701

Schaefer, B. E. 2008, AJ, 135, 112

Schöier, F. L., van der Tak, F. F. S., van Dishoeck, E. F., & Black, J. H. 2005, A&A, 432, 369

Sheffer, Y., Rogers, M., Federman, S. R., et al. 2008, ApJ, 687, 1075

Sofia, U. J., Cardelli, J. A., Guerin, K. P., & Meyer, D. M. 1997, ApJL, 482, L105

Spitzer, L., Jr. 1948, ApJ, 107, 6

Stanimirovic, S., Staveley-Smith, L., Dickey, J. M., Sault, R. J., & Snowden, S. L. 1999, MNRAS, 302, 417

Staveley-Smith, L., Kim, S., Calabretta, M. R., Haynes, R. F., & Kesteven, M. J. 2003, MNRAS, 339, 87

Sternberg, A., Le Petit, F., Roueff, E., & Le Bourlot, J. 2014, ApJ, 790, 10

Stutzki, J., Graf, U. U., Haas, S., et al. 1997, ApJL, 477, L33

Tang, N., Li, D., Heiles, C., et al. 2016, A&A, 593, A42

Tielens, A. G. G. M., & Hollenbach, D. 1985, ApJ, 291, 722

Tumlinson, J., Shull, J. M., Rachford, B. L., et al. 2002, ApJ, 566, 857

van der Tak, F. F. S., Black, J. H., Schöier, F. L., Jansen, D. J., & van Dishoeck, E. F. 2007, A&A, 468, 627

van Dishoeck, E. F., & Black, J. H. 1988, ApJ, 334, 771

Viallefond, F., & Goss, W. M. 1986, A&A, 154, 357

Wang, M., Chin, Y.-N., Henkel, C., Whiteoak, J. B., & Cunningham, M. 2009, ApJ, 690, 580

Weingartner, J. C., & Draine, B. T. 2001, ApJ, 548, 296

Welty, D. E., Lauroesch, J. T., Wong, T., & York, D. G. 2016, ApJ, 821, 118

Welty, D. E., Xue, R., & Wong, T. 2012, ApJ, 745, 173

Westerlund, B. E. 1997, The Magellanic Clouds (New York: Cambridge Univ. Press)

Whittet, D. C. B., Gerakines, P. A., Hough, J. H., & Shenoy, S. S. 2001, ApJ, 547, 872

Wiesenfeld, L., & Goldsmith, P. F. 2014, ApJ, 780, 183

Wilson, N. J., & Bell, K. L. 2002, MNRAS, 337, 1027

Wolfire, M. G., Hollenbach, D., & McKee, C. F. 2010, ApJ, 716, 1191

Wolfire, M. G., Hollenbach, D., McKee, C. F., Tielens, A. G. G. M., & Bakes, E. L. O. 1995, ApJ, 443, 152

Wolfire, M. G., McKee, C. F., Hollenbach, D., & Tielens, A. G. G. M. 2003, ApJ, 587, 278

Wong, T., Hughes, A., Ott, J., et al. 2011, ApJS, 197, 16

Wright, E. L., Mather, J. C., Bennett, C. L., et al. 1991, ApJ, 381, 200