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

# Resolved UV and [C II] Structures of Luminous Galaxies within the Epoch of Reionization

J. Matthee<sup>1,2,10</sup>①, D. Sobral<sup>3</sup>①, L. A. Boogaard<sup>2</sup>①, H. Röttgering<sup>2</sup>, L. Vallini<sup>2</sup>①, A. Ferrara<sup>4,5</sup>①, A. Paulino-Afonso<sup>6</sup>, F. Boone<sup>7</sup>, D. Schaerer<sup>7,8</sup>①, and B. Mobasher<sup>9</sup>

<sup>1</sup>Department of Physics, ETH Zürich, Wolfgang-Pauli-Strasse 27, 8093 Zürich, Switzerland; mattheej@phys.ethz.ch

<sup>2</sup>Leiden Observatory, Leiden University, P.O. Box 9513, NL-2300 RA, Leiden, The Netherlands

<sup>3</sup>Department of Physics, Lancaster University, Lancaster, LA1 4YB, UK

<sup>4</sup>Scuola Normale Superiore, Piazza dei Cavalieri 7, I-56126 Pisa, Italy

<sup>5</sup>Kavli Institute for the Physics and Mathematics of the Universe (WPI), University of Tokyo, Kashiwa 277-8583, Japan

<sup>6</sup>Max-Planck-Institut für Extraterrestrische Physik, Postfach 1312, Giessenbachstr., D-85748 Garching, Germany

<sup>7</sup>Université de Toulouse; UPS-OMP; IRAP; Toulouse, France

<sup>8</sup>Observatoire de Genève, Université Genève, 51 Ch. des Maillettes, 1290 Versoix, Switzerland

<sup>9</sup>Department of Physics and Astronomy, University of California, Riverside, CA 92521, USA

Received 2019 March 19; revised 2019 June 27; accepted 2019 July 2; published 2019 August 21

## Abstract

We present new deep ALMA and Hubble Space Telescope (HST)/WFC3 observations of MASOSA and VR7, two luminous Lyα emitters (LAEs) at $z = 6.5$, for which the UV continuum levels differ by a factor of four. No IR dust continuum emission is detected in either, indicating little amounts of obscured star formation and/or high dust temperatures. MASOSA, with a UV luminosity $M_{1500} = -20.9$, compact size, and very high Lyα EW$_0 \approx 145$ Å, is undetected in [C II] to a limit of $L_{[C II]} < 2.2 \times 10^7$ L$_\odot$, implying a metallicity $Z \lesssim 0.07$ Z$_\odot$. Intriguingly, our HST data indicate a red UV slope $\beta = -1.1 \pm 0.7$, at odds with the low dust content. VR7, which is a bright ($M_{1500} = -22.4$) galaxy with moderate color ($\beta = -1.4 \pm 0.3$) and Lyα EW$_0 = 34$ Å, is clearly detected in [C II] emission (S/N = 15). VR7's rest-frame UV morphology can be described by two components separated by $\approx 1.5$ kpc and is globally more compact than the [C II] emission. The global [C II]/UV ratio indicates $Z \approx 0.2$ Z$_\odot$, but there are large variations in the UV/[C II] ratio on kiloparsec scales. We also identify diffuse, possibly outflowing, [C II]-emitting gas at $\approx 100$ km s$^{-1}$ with respect to the peak. VR7 appears to be assembling its components at a slightly more evolved stage than other luminous LAEs, with outflows already shaping its direct environment at $z \sim 7$. Our results further indicate that the global [C II]–UV relation steepens at SFR $< 30$ M$_\odot$ yr$^{-1}$, naturally explaining why the [C II]/UV ratio is anticorrelated with Lyα EW in many, but not all, observed LAEs.

**Key words:** dark ages, reionization, first stars – galaxies: formation – galaxies: high-redshift – galaxies: ISM – galaxies: kinematics and dynamics

## 1. Introduction

The advent of the Atacama Large Millimeter/submillimeter Array (ALMA) has enabled the first detailed studies of the interstellar medium (ISM) in various kinds of star-forming galaxies and quasar hosts within the epoch of reionization at $z > 6$ (e.g., Wang et al. 2013; Maiolino et al. 2015; Willott et al. 2015; Inoue et al. 2016; Decarli et al. 2017; D'Odorico et al. 2018). ALMA also opened the opportunity of resolving internal structure at kiloparsec scales (e.g., Jones et al. 2017; Matthee et al. 2017b; Carniani et al. 2018b; Hashimoto et al. 2019). Currently, ALMA observations of galaxies at $z > 6$ focus on measuring the dust far-infrared (FIR) continuum and FIR fine-structure lines [C II]$_{158\,\mu m}$ and [O III]$_{88\,\mu m}$. The [C II] line partly traces H II regions and neutral gas (i.e., photo-dissociating regions [PDRs]; Vallini et al. 2013), while the [O III] line only traces H II regions. Combined with detailed photoionization modeling, these lines can be used to constrain gas properties such as the metallicity and ionization state (e.g., Olsen et al. 2017; Vallini et al. 2017).

While the dust continuum is typically undetected in galaxies selected through their UV or Lyα emission (e.g., Capak et al. 2015; Bouwens et al. 2016b; Pentericci et al. 2016;

Matthee et al. 2017b; but see Bowler et al. 2018; Hashimoto et al. 2019; Tamura et al. 2019 for counterexamples), the fine-structure lines are now routinely detected. Due to the low metallicity and higher ionization state of the ISM in high-redshift galaxies, the [O III] line may be the strongest FIR emission line (e.g., Ferkinhoff et al. 2010; Inoue et al. 2016), unlike at low redshift, where [C II] is typically the strongest (e.g., Stacey et al. 1991; Wolfire et al. 2003; Herrera-Camus et al. 2015). Indeed, recent studies have shown [O III] to be more luminous than [C II] in objects for which both lines are constrained (Inoue et al. 2016; Carniani et al. 2017; Hashimoto et al. 2019; Walter et al. 2018).

However, even though being potentially less luminous in high-redshift galaxies, the [C II] line is more easily observed for the vast majority of galaxies currently known at $6 \lesssim z \lesssim 7$. This is because [C II] redshifts into favorable frequencies in ALMA bands 6 and 7 at $z > 5$, while this only happens at $z \gtrsim 8$ for [O III] (with the exception of a few narrow redshift windows around $z \sim 7$, e.g., Carniani et al. 2017, where the sensitivity to [C II] is still a factor of $\approx 2$ better). Therefore, practically, observations of the [C II] line are still the least expensive for measuring systemic galaxy redshifts (see Lyα) and studying ISM properties and kinematics.

Some early ALMA observations of galaxies at $z \approx 5-6$ (e.g., Capak et al. 2015) indicated relatively luminous [C II] emission

<footnote n="10"><sup>10</sup> Zwicky Fellow.</footnote>
relative to the UV star formation rate (SFR), compared to the observed relation in the local universe (De Looze et al. 2014). Several other searches at $z \approx 6-7$ either found a moderate [C II] deficit (Pentericci et al. 2016; Bradač et al. 2017) or report stringent upper limits on the [C II] luminosity that would place these distant galaxies far below the local relation (e.g., Ouchi et al. 2013; Ota et al. 2014). Such a deficit could potentially be due to the prerequisite of a known Ly$\alpha$ redshift, which leads to a bias of observing young, metal-poor systems (e.g., Trainor et al. 2016). New observations of additional galaxies (e.g., Matthee et al. 2017b; Carniani et al. 2018b; Hashimoto et al. 2019; Smit et al. 2018) and reanalysis of earlier ALMA data (Carniani et al. 2018a) indicate that there is more scatter at high redshift (Carniani et al. 2018b) instead of a strongly preferred low or high [C II]/UV ratio. The [C II]/UV ratio correlates with the relative strength of Ly$\alpha$ emission (Carniani et al. 2018b; Harikane et al. 2018). These measurements are nonetheless complicated by the observation that [C II] is often offset spatially from the UV emission and multiple components of [C II] emission can be associated with UV components (e.g., Matthee et al. 2017b).

In this paper, we present new ALMA and *Hubble Space Telescope (HST)* observations of two very luminous Ly$\alpha$ emitters (LAEs) at $z = 6.5$, named VR7 and MASOSA (Sobral et al. 2015; Matthee et al. 2017a), combined with a multi-wavelength analysis of archival data. While these galaxies have similar Ly$\alpha$ luminosity, the UV continuum is very different, leading to a factor of $\approx 5$ difference in the equivalent width (EW). What causes these differences in their Ly$\alpha$ EWs? Are the SFRs, dust content, and/or ages of the galaxies different? Are there any differences in their metallicities? To answer these questions, we focus on characterizing the properties of the ISM (metallicity, kinematics, dust content) and the stellar populations (SFR and age) and compare these to other galaxies observed at $z \approx 6-7$.

Our targets resemble the Ly$\alpha$ luminosity of the well-known LAEs Himiko and CR7 (Ouchi et al. 2013; Sobral et al. 2015), which are clearly resolved into several components separated by $\approx 4$ kpc both in the rest-frame UV and in [C II] emission (e.g., Matthee et al. 2017b; Carniani et al. 2018b; Sobral et al. 2019). Do these luminous LAEs also consist of several components, which could indicate that mergers are ubiquitous in the strongest Ly$\alpha$ emitters? We address this with high spatial resolution observations in both rest-frame UV and [C II] emission.

The structure of this paper is as follows. We present the new ALMA and *HST/WFC3* observations and their data reduction in Section 2. Section 3 presents the rest-frame UV and Ly$\alpha$ properties of both galaxies and compares them to the observed galaxy population at similar look-back time. We then investigate the FIR continuum data from ALMA to constrain the galaxies’ dust continuum emission in Section 4. In Section 5 we present integrated measurements of the [C II] line, while we present the structure in VR7, resolved in both the spatial and spectral dimension, in Section 6. In Section 7, we discuss the nature of MASOSA and which physical properties determine the [C II]/UV ratio in high-redshift galaxies. Finally, we summarize our results in Section 8. We use a $\Lambda$CDM cosmology with $\Omega_{\Lambda} = 0.70$, $\Omega_{M} = 0.30$, and $H_{0} = 70 \text{ km s}^{-1} \text{ Mpc}^{-1}$ and a Salpeter (1955) initial mass function (IMF). Throughout the paper, UV and IR luminosities are converted to SFR following Kennicutt (1998).

<table><caption>**Table 1**
Overview of the Different ALMA Reductions Used in Various Parts of Our Analysis</caption><thead><tr><th>Reduction</th><th>Beam Size</th><th>Used in Section</th></tr></thead><tbody><tr><td>VR7</td><td></td><td></td></tr><tr><td>Natural, 650 k&lambda; taper</td><td>0&prime;64 &times; 0&prime;61</td><td>Sections 4, 5, 6.5.2</td></tr><tr><td>Briggs, Robust = 0.5</td><td>0&prime;47 &times; 0&prime;40</td><td>Section 6</td></tr><tr><td>MASOSA</td><td></td><td></td></tr><tr><td>Natural, 650 k&lambda; taper</td><td>0&prime;68 &times; 0&prime;60</td><td>Sections 4, 5</td></tr></tbody></table>

## 2. Observations and Data Reduction

### 2.1. ALMA

We performed ALMA observations in band 6 of our targets aiming at detecting [C II] line emission and FIR dust continuum emission through ALMA program 2017.1.01451.S in Cycle 5. We observed in four spectral windows: two centered around redshifted [C II]$_{158 \mu m}$ emission at $z \approx 6.5$ ($\approx 252$ GHz) and two around 235 GHz, each with a bandwidth of 1875 MHz and 7.8 MHz ($9-10 \text{ km s}^{-1}$) resolution. Both galaxies were observed with 43 antennas in configuration C43-4, allowing baselines from 15 to 783 m, leading to a natural resolution of $\approx 0.4$. Each galaxy was observed for three executions that consisted of 49 minutes of integration time on target, leading to a total on-target exposure of 147 minutes.

MASOSA (Matthee et al. 2015; Sobral et al. 2015) was observed during three executions on 2018 March 23–24 with precipitable water vapor (PWV) columns of 0.9, 1.9, and 2.0 mm, respectively. The quasar J1058+0133 was used as an atmospheric bandpass and flux calibrator, while quasar J0948+0022 was used as a phase calibrator. VR7 (Matthee et al. 2017a) was observed under good conditions (PWV columns of 1.0, 1.0, and 0.7 mm for each execution, respectively) on 2018 September 6 and 7. The quasars J2148+0657 and J2226+0052 were used for atmospheric bandpass and flux calibrator and phase calibrator, respectively.

The data have been reduced and calibrated using CASA version 5.1.1-5 following the standard pipeline procedures. The final imaging was performed using the CLEAN task, with different choices for weighting and tapering depending on the specific scientific question. In general, using natural weighting, we measure a background rms = 6 $\mu$Jy beam$^{-1}$ in the continuum and a background rms = 0.08–0.10 mJy beam$^{-1}$ in 18 km s$^{-1}$ channels.

Our analysis and reduction strategy was as follows: We first reduced the data with natural weighting and a UV tapering of 650 k$\lambda$ (corresponding to an on-sky FWHM of 0.3) and by averaging over two velocity channels, leading to a data cube with a beam FWHM $\approx 0.7 \times 0.7$ and $\approx 18 \text{ km s}^{-1}$ velocity resolution. The tapering was used to optimize detectability of extended emission, while the typical [C II] line with FWHM $\approx 150 \text{ km s}^{-1}$ would still be well resolved. This reduction was used for initial inspection. For detections, we then reimaged the data with higher spatial resolution (through Briggs weighting with robust parameter 0.5 and without tapering) and/or without channel averaging, with specific parameters chosen optimized for the science question (and motivated in each relevant section independently and summarized in Table 1). For nondetections, we test the robustness of not detecting the line/continuum by changing the reduction
method, and we use the tapered reduction motivated by the properties of similar galaxies.

## 2.2. HST/WFC3

We present HST/WFC3 observations of MASOSA observed as part of HST program 14699 (PI: Sobral) on 2018 March 6. Observations were performed for two orbits with the F110W and F160W filters. During each orbit, four exposures of ≈650 s were taken following the standard WFC3/IR dither pattern. This results in total exposure times of 5.2 ks. We acquire flatfielded and calibrated FLT images from the STScI server and fix the astrometric solution to the ground-based near-infrared data (which are fixed to GAIA DR2 astrometry). Then, after masking bad pixels and cosmic rays, we median-combine individual exposures to a final image with 0.′′064 pixel scale with bilinear interpolation using SWARP (Bertin 2010). We also apply this reduction strategy to the HST/WFC3 data on VR7 (originating from the same HST program) that were presented earlier in Matthee et al. (2017a). Compared to the previous reduction, the spatial resolution is increased slightly as the image is better sampled. We have verified that the integrated flux and sensitivity are consistent within 5%.

We use unsaturated, high signal-to-noise ratio (S/N) detections of stars to measure a resolution of FWHM = 0.′′25 and 0.′′28 in the F110W and F160W images, respectively. The depth is estimated by measuring the standard deviation of the total counts in 1000 apertures with 0.′′8 diameter placed in empty sky regions. We measure a 3σ limit of F110W = 27.2 and F160W = 26.6 AB magnitude after correcting for aperture losses.

# 3. The Properties of Targeted Galaxies

In this section, we present the general properties of the targeted galaxies based predominantly on Lyα measurements and HST/WFC3 data. The measurements are summarized in Table 2. Our targets have been selected on their high Lyα luminosity (>2 × L*; e.g., Matthee et al. 2015; Konno et al. 2018), a confirmed spectroscopic redshift z ~ 6.5, and their observability with ALMA. The Lyα luminosities of our targets are relatively high compared to the majority of earlier ALMA observations, but not extreme (Figure 1).

MASOSA, at z = 6.541 ± 0.001 (Sobral et al. 2015), is located in the COSMOS field, such that there are Spitzer/IRAC data available through the SPLASH program (in particular in the [3.6] and [4.5] bands; Steinhardt et al. 2014). VR7, at z = 6.534 ± 0.001 (Matthee et al. 2017a), is located in the CFHTLS-W4/SA22 field, where no deep Spitzer/IRAC data are available. Upper limits on rest-frame UV lines besides Lyα based on X-SHOOTER observations (i.e., C IV, He II, C III]) are presented for VR7 in Matthee et al. (2017a). The nondetections of high-ionization lines and the narrow Lyα line widths (<400 km s⁻¹) indicate that our targets are not powered by a strong active galactic nucleus (AGN; e.g., Sobral et al. 2018).

## 3.1. Lyα Luminosity and Rest-frame EW

We measure the Lyα flux using a combination of narrowband (NB921) and broadband (z) imaging from Subaru/Suprime-Cam as described in Matthee et al. (2015). The transmission of the NB921 filter is not a perfect top hat (which is assumed in the flux calculation in Matthee et al. 2015). Therefore, we correct the line

<caption>**Table 2**</caption>

The Integrated Properties of VR7 and MASOSA as Measured from HST and Ground-based Imaging Data, Optical Spectroscopy, and ALMA Observations

<table><thead><tr><th>Property</th><th>VR7</th><th>MASOSA</th></tr></thead><tbody><tr><td>R.A. (J2000)</td><td>22:18:56.36</td><td>10:01:24.80</td></tr><tr><td>Decl. (J2000)</td><td>+00:08:07.32</td><td>+02:31:45.34</td></tr><tr><td>M<sub>1500</sub></td><td>-22.37<sup>+0.05</sup><sub>-0.05</sub></td><td>-20.94<sup>+0.14</sup><sub>-0.13</sub></td></tr><tr><td>β<sub>obs</sub></td><td>-1.49<sup>+0.26</sup><sub>-0.26</sub></td><td>-1.49<sup>+0.56</sup><sub>-0.58</sub></td></tr><tr><td>β<sub>Lyα corr.</sub></td><td>-1.38<sup>+0.29</sup><sub>-0.27</sub></td><td>-1.06<sup>+0.68</sup><sub>-0.72</sub></td></tr><tr><td>r<sub>eff,UV</sub> (kpc)</td><td>1.56 ± 0.05</td><td>1.12<sup>+0.44</sup><sub>-0.19</sub></td></tr><tr><td>L<sub>Lyα</sub> (erg s<sup>-1</sup>)</td><td>2.4<sup>+0.2</sup><sub>-0.2</sub> × 10<sup>43</sup></td><td>2.4<sup>+0.4</sup><sub>-0.4</sub> × 10<sup>43</sup></td></tr><tr><td>EW<sub>Lyα,0</sub> (Å)</td><td>34<sup>+4</sup><sub>-4</sub></td><td>145<sup>+50</sup><sub>-43</sub></td></tr><tr><td>z<sub>Lyα</sub></td><td>6.534 ± 0.001</td><td>6.543 ± 0.003</td></tr><tr><td>ν<sub>FWHM, Lyα</sub> (km s<sup>-1</sup>)</td><td>340 ± 14</td><td>386 ± 30</td></tr><tr><td>[3.6]-[4.5]</td><td>...</td><td>-0.6 ± 0.2</td></tr><tr><td colspan="3"><strong>[C II] and IR properties</strong></td></tr><tr><td>S<sub>ν,fit</sub>Δν (mJy km s<sup>-1</sup>)</td><td>417.8 ± 80.3</td><td>&lt;12.7</td></tr><tr><td>L<sub>[C II],CASA</sub> (10<sup>8</sup> L<sub>☉</sub>)</td><td>4.76 ± 0.44</td><td>&lt;0.22</td></tr><tr><td>z<sub>[C II]</sub></td><td>6.5285<sup>+0.0007</sup><sub>-0.0004</sub></td><td>...</td></tr><tr><td>ν<sub>FWHM, [C II]</sub> (km s<sup>-1</sup>)</td><td>200<sup>+48</sup><sub>-32</sub></td><td>...</td></tr><tr><td>Δν<sub>Lyα</sub> (km s<sup>-1</sup>)</td><td>-217<sup>+29</sup><sub>-18</sub></td><td>...</td></tr><tr><td>r<sub>1/2, major, [C II]</sub> (kpc)</td><td>3.9 ± 0.4</td><td>...</td></tr><tr><td>r<sub>1/2, minor, [C II]</sub> (kpc)</td><td>1.9 ± 0.2</td><td>...</td></tr><tr><td>f<sub>ν,160 µm</sub> (µJy beam<sup>-1</sup>)</td><td>&lt;10.6</td><td>&lt;9.2</td></tr><tr><td colspan="3"><strong>Derived properties</strong></td></tr><tr><td>SFR<sub>UV, no dust</sub> (M<sub>☉</sub> yr<sup>-1</sup>)</td><td>54<sup>+3</sup><sub>-2</sub></td><td>15<sup>+2</sup><sub>-2</sub></td></tr><tr><td>SFR<sub>UV, Meurer</sub> (M<sub>☉</sub> yr<sup>-1</sup>)</td><td>279<sup>+128</sup><sub>-118</sub></td><td>208<sup>+193</sup><sub>-171</sub></td></tr><tr><td>SFR<sub>UV+IR</sub> (M<sub>☉</sub> yr<sup>-1</sup>)</td><td>54<sup>+5</sup><sub>-2</sub></td><td>15<sup>+4</sup><sub>-2</sub></td></tr><tr><td>EW(Hβ+[O III]) (Å)</td><td>...</td><td>≈1500</td></tr><tr><td>Z<sub>[C II]</sub>/Z<sub>☉</sub></td><td>≈0.2</td><td>≤0.07</td></tr><tr><td>L<sub>IR, T<sub>d</sub>=45 K</sub> / 10<sup>10</sup> L<sub>☉</sub></td><td>&lt;2.6</td><td>&lt;2.3</td></tr><tr><td>SFR<sub>IR, T<sub>d</sub>=45 K</sub> (M<sub>☉</sub> yr<sup>-1</sup>)</td><td>&lt;4.6</td><td>&lt;4.0</td></tr><tr><td>M<sub>dyn</sub>/(sin i)<sup>2</sup> (10<sup>10</sup> M<sub>☉</sub>)</td><td>3.0<sup>+0.7</sup><sub>-0.6</sub></td><td>...</td></tr></tbody></table>

**Note.** Errors show the 16th and 84th percentiles, while upper limits are at the 1σ level. See text in Section 5.2 that motivates the [C II] luminosity limit for MASOSA.

flux for the actual transmission at the specific wavelength where Lyα is observed, which is known thanks to spectroscopy. Our flux measurements are therefore not dependent on (assumptions on) slit losses or the flux calibration of spectra. VR7 and MASOSA have an almost identical integrated Lyα luminosity, while their UV continuum luminosity differs by a factor of ≈4 (see the next subsection). Measurements of the UV continuum luminosity and the UV slope are used to constrain the continuum around Lyα and to estimate the EW. The Lyα EW of VR7 is moderate (EW₀ = 34 ± 4 Å), while the EW of MASOSA is extremely high (albeit with large uncertainties; EW₀ = 145<sup>+50</sup><sub>-43</sub> Å). The emissions of both Lyα lines consist of a single, clearly asymmetric red peak, with FWHM ranging from 340 to 390 km s⁻¹ (Sobral et al. 2015; Matthee et al. 2017a).

## 3.2. UV Luminosity and Colors

We measure the rest-frame UV luminosity using the HST/WFC3 F110W and F160W data. Measurements are performed with 1.′′2 diameter apertures to account for the object sizes and include corrections for missing encapsulated flux based on the WFC3 manual. We correct the F110W photometry for the contribution from Lyα emission based on the measured Lyα
<figure><img src="image_2.png" /><figcaption>**Figure 1.** Observed Lyα and UV luminosities of MASOSA and VR7 compared to the galaxy population at z ≈ 6–7. Gray pentagons show a compilation of LAEs from Curtis-Lake et al. (2012), Jiang et al. (2013), Matthee et al. (2017a), Shibuya et al. (2018), and Matthee et al. (2018), while green diamonds show galaxies at z ≈ 6–7 targeted by previous ALMA programs as compiled in Appendix C. Solid lines indicate the Lyα escape fraction assuming ξ<sub>ion</sub> = 10<sup>25.4</sup> Hz erg<sup>-1</sup> (e.g., Bouwens et al. 2016a), negligible Lyman continuum escape fraction, and no dust or intergalactic medium attenuation. This ξ<sub>ion</sub> implies an intrinsic Lyα EW = 146 Å (Sobral & Matthee 2019). VR7 is among the most UV-luminous galaxies targeted by ALMA so far, while MASOSA is among the galaxies with the highest Lyα escape fraction.</figcaption></figure>

luminosity and measure the UV slope $\beta$ using the F110W
–F160W colors following Ono et al. (2010).

For MASOSA, we measure F110W = 25.74<sup>+0.13</sup><sub>-0.11</sub> and
F160W = 25.58<sup>+0.23</sup><sub>-0.20</sub>. These measurements translate into
an absolute UV luminosity M<sub>1500</sub> = -20.94<sup>+0.14</sup><sub>-0.13</sub> and a red
UV slope β = -1.06<sup>+0.68</sup><sub>-0.72</sub> (-0.14 mag brighter and β<sub>obs</sub> =
-1.49<sup>+0.56</sup><sub>-0.58</sub> if we would not correct for the Lyα contribution). The
potential red color is intriguing given the high Lyα EW, as the
Lyα escape fraction is typically lower for red galaxies (e.g.,
Matthee et al. 2016). We discuss this in more detail in Section 7.1.
MASOSA has Spitzer/IRAC magnitudes [3.6] = 24.0 ± 0.1 and
[4.5] = 24.6 ± 0.2 (Laigle et al. 2016), resulting in a relatively
blue color [3.6]–[4.5] = -0.6 ± 0.2, similar to other LAEs at
z ≈ 6.5 (e.g., Harikane et al. 2018).

VR7 is very luminous in the rest-UV, with F110W = 24.35<sup>+0.05</sup><sub>-0.05</sub> and F160W = 24.28<sup>+0.08</sup><sub>-0.07</sub>, which translates into a UV luminosity $M_{1500} = -22.37_{-0.05}^{+0.05}$. VR7 is also relatively red, with $\beta = -1.38_{-0.27}^{+0.29}$, yet in agreement with the observed relation between UV slope and UV magnitude extrapolated to the luminosity of VR7 (Bouwens et al. 2012); see Figure 2. Due to its lower EW, the correction for the Ly$\alpha$ line has less impact ($-0.04$ mag and $\beta_{\text{obs}} = -1.49_{-0.26}^{+0.26}$).

3.3. Rest-frame UV Sizes

We measure the deconvolved UV sizes of MASOSA and VR7 by modeling the light distribution with an exponential profile using the IMFIT software package (Erwin 2015). The PSF image is created by averaging the (normalized) light profiles of four nearby nonsaturated stars and can be modeled accurately with a Moffat profile with FWHM = 0.′′25. As the morphology in the F110W filter may be affected by Lyα emission, we measure morphology in the F160W filter (see Figure 3; see Appendix B for details on the fitting). For MASOSA, we measure $r_{\text{eff}} = 1.12_{-0.19}^{+0.44}$ kpc (corresponding to

<figure><img src="image_3.png" /><figcaption>**Figure 2.** Observed UV slope $\beta$ and UV luminosities of MASOSA and VR7 compared to galaxy compilations at $z \approx 5$ (from Capak et al. 2015) and $z \approx 6-7$ observed with ALMA and to the general Lyman break population from Bouwens et al. (2015). Red stars highlight sources with the Ly$\alpha$ EW$_0 > 50$ Å. The shaded region shows the fitted relation (Bouwens et al. 2015) and its uncertainty, while we show its extrapolation to higher UV luminosities with a dotted line.</figcaption></figure>

<figure><img src="image_4.png" /><figcaption>**Figure 3.** HST/WFC3 F160W thumbnail images of VR7 (left) and MASOSA (right) with matched scale and contrast. Contour levels are at 3$\sigma$, 5$\sigma$, and 7$\sigma$. VR7's UV emission (rest-frame wavelength $\lambda_0 \approx 2000$ Å) is more luminous and more extended than MASOSA. MASOSA is compact but has an indication of a second component that is only seen in F160W.</figcaption></figure>

FWHM = 0.41<sup>+0.16</sup><sub>-0.07</sub> arcsec, so marginally resolved in the data)
with an ellipticity 0.12<sup>+0.33</sup><sub>-0.12</sub> and a position angle 27° ± 27°,
meaning that the UV light is almost spherically symmetric. We
note that we would measure a slightly smaller size when
allowing the Sérsic index to vary ($r_{\text{eff}} \approx 0.9$ kpc for $n \approx 0.1$) or
when using the F110W filter ($r_{\text{eff}} = 0.87 \pm 0.05$ kpc, indicat-
ing relatively compact Ly$\alpha$ emission). While the F160W image
hints toward a faint second clump, a two-component
exponential does not provide a better fit to the data.

VR7 is well resolved, and a single exponential model
results in $r_{\text{eff}} = 1.56 \pm 0.05$ kpc with PA = 80 $\pm$ 2 and
ellipticity $0.55 \pm 0.02$. Relaxing the Sérsic index results again
in a slightly smaller size ($r_{\text{eff}} = 1.52 \pm 0.05$ kpc) with
$n = 0.4 \pm 0.08$. In Appendix B we show that the HST data
from VR7 are better described by two slightly smaller
exponential components with a separation of $\approx 0.′′35$
(1.9 kpc), but we list the single-component fit in Table 2 for
consistency with integrated [C II] and UV luminosity
measurements.
<figure><img src="image_5.png" /><figcaption>**Figure 4.** FIR continuum of VR7 (left) and MASOSA (right) in primary-beam-corrected data with natural weighting and 0.03 taper (beam major axes of ≈0.07). The black lines show the ±2σ, 3σ, and 4σ contours, where the 1σ level ranges from 9.2 to 10.6 µJy beam⁻¹. Dark-blue contours illustrate the morphologies in the HST/WFC3 F110W band, while dark red shows the F160W band. Both bands trace the rest-frame UV emission. No FIR continuum emission is detected at the locations of VR7 or MASOSA, independent of the data reduction method.</figcaption></figure>

<figure><img src="image_6.png" /></figure>

### 3.4. How Do These Galaxies Compare to the General Galaxy Population?

We compare the UV and Lyα luminosities of MASOSA and VR7 to other LAEs identified at $z \sim 6.5$ and to other UV-selected galaxies at $z \sim 6-7$ observed with ALMA at similar rest-frame frequency (see Appendix C). The latter sample mostly comprises galaxies for which the redshifts have previously been confirmed with (strong) Lyα emission, as ALMA observations require that the redshift is precise to within $\Delta z \approx 0.05$, which is smaller than the typical uncertainty of photometric redshifts. A few notable exceptions include the galaxy B14-6566 at $z = 7.15$ with faint Lyα emission (Hashimoto et al. 2019) and two Lyman break galaxies at $z \approx 6.8$ for which the photometric redshift could reliably be estimated (Smit et al. 2018); see Appendix C for a complete overview.

As shown in Figure 1, VR7 is among the most luminous galaxies within the epoch of reionization ever observed with ALMA. MASOSA has an $\approx M_{1500}^*$ luminosity and is among the objects with the highest Lyα EW. We note that the spatial resolution of the ALMA observations presented in this work is a factor of ~2 higher than most previous observations, allowing more detailed investigation of dynamics and kinematics. The UV sizes of MASOSA and VR7 are comparable to the sizes of other galaxies with similar luminosity (e.g., Shibuya et al. 2015; Bowler et al. 2017).

## 4. The FIR Continuum

The ALMA data in the four spectral windows are combined to create 235–250 GHz (1.2 mm) continuum images centered on both objects. These continuum images constrain the FIR continuum luminosity. For VR7, we mask the frequencies where [C II] is detected (see Section 5), while for MASOSA we use the full frequency range. Several foreground objects are identified in the continuum maps, which can all be associated with objects in HST, Spitzer, or VLT/MUSE data

(Matthee et al. 2019, in preparation). These objects lie in the foreground of our targets because of detections in images with filters blueward of the Lyman break at $z \approx 6.5$. The foreground objects confirm that the relative astrometry between the ALMA and HST data and ground-based imaging is accurate to within $\approx 0.01''$. The general properties of these foreground objects and estimates of their redshifts are summarized in Table 5.

### 4.1. Upper Limits at $\lambda_0 \approx 160$ µm

No FIR continuum emission is detected at the locations of VR7 or MASOSA (Figure 4). This result does not depend on the weighting or tapering applied when imaging the visibility data from ALMA, meaning that it is unlikely that flux is resolved out. As source sizes are $\approx 0.06''-1''$ (major axis) in the rest-frame UV, we use ALMA continuum images constructed with natural weighting and 650 kλ tapering (synthesized beam FWHM $\approx 0.07''$) to measure physically motivated and conservative upper limits. For VR7, we measure a $1\sigma$ limiting $f_{\lambda_0=160\mu m} < 10.6\mu Jy$ beam⁻¹, while for MASOSA the $1\sigma$ limit is $f_{\lambda_0} = 160\mu m < 9.2\mu Jy$ beam⁻¹.

We compare our measured upper limits on the continuum flux density around $\lambda_0 = 160 \mu m$ with those measured in other UV-selected galaxies at $z \approx 6-7$ in Figure 5. We show IR flux densities, instead of (temperature-dependent) IR luminosities, but note that we corrected the flux density for CMB heating assuming a dust temperature $T = 45$ K, leading to a median correction of $\approx 0.03$ dex (da Cunha et al. 2013). We also indicate the expected 1500 Å/160 µm flux density ratios for local galaxies with extremely low metallicities as inferred from a compilation of galaxies in the local universe by Maiolino et al. (2015). For comparison, we show the upper limiting flux density ratio for low-metallicity dwarf galaxy I Zw 18, which we determine using measurements from GALEX/near-UV (Gil de Paz et al. 2007) and Herschel/PACS $F_{160}$ (Rémy-Ruyer et al. 2015). This upper limit is comparable to the upper limit for VR7. The IR continuum sensitivity in both MASOSA and
<figure><img src="image_7.png" /><figcaption>**Figure 5.** UV continuum luminosity vs. the ratio of the UV ($\lambda_0$ = 0.15 µm) and IR ($\lambda_0$ = 160 µm) flux densities for the galaxies observed in this paper and compilations of galaxies observed with ALMA at $z \approx 5$ and $z \approx 6-7$. Dotted gray lines indicate the 160 µm/1500 Å ratio at which the unobscured SFRs are 0.2, 1, and 5 times the obscured SFR, assuming a dust temperature of 45 K and Kennicutt (1998) conversions. We indicate the typical gas-phase metallicities measured in galaxies in the local universe with comparable flux density ratios (based on a compilation by Maiolino et al. 2015) and the upper limit in the extremely low metallicity galaxy I Zw 18, but we stress that the flux ratio is not exclusively metallicity sensitive.</figcaption></figure>

VR7 is higher than other galaxies of comparable UV luminosity. VR7 is the most UV-luminous galaxy for which the continuum around $\lambda_0 = 160$ µm is currently undetected. This extremely low 160 µm/1500 Å ratio indicates that VR7 has a lower dust-to-gas ratio and/or higher dust temperature compared to other luminous galaxies. The only other galaxy with even stronger limits on the 160 µm/1500 Å flux density ratio is the luminous LAE CR7 (Matthee et al. 2017b).

## 4.2. IR Luminosity and Total SFR

In order to use these measurements to derive upper limits on the total IR luminosity ($L_{IR}$), one has to make assumptions on the distribution of dust temperatures that combined result in a modified blackbody spectral energy distribution (SED), as an upper limit at a single frequency cannot constrain the shape of this SED. Moreover, at $z \sim 7$ the temperature of the cosmic microwave background is $\approx 20$ K, which contributes to the heating of dust grains (da Cunha et al. 2013). The choice of (luminosity-weighted) dust temperature strongly impacts the integrated IR luminosity. For example, an increase in the dust temperature from 25 to 45 K decreases the limiting $L_{IR}$ by a factor of $\approx 4$ when the continuum is constrained around $\lambda_0 \approx 160$ µm (e.g., Schaerer et al. 2015).

While dust temperatures in typical star-forming galaxies in the $z < 2$ universe are $\approx 35$ K (e.g., Rémy-Ruyer et al. 2013), it is plausible that the dust temperature is higher at higher redshifts, in particular in galaxies with hard ionizing sources (i.e., early generations of stars) and a low metallicity (e.g., Cen & Kimm 2014; Maiolino et al. 2015). While direct constraints on the dust temperatures in high-redshift galaxies require large investments of ALMA time (e.g., Bouwens et al. 2016b), Faisst et al. (2017) show that the luminosity-weighted dust temperature in high-$z$ analogs at low redshift may be as high as 70 K. Behrens et al. (2018) furthermore simulate that intense radiation fields in young galaxies could even result in dust temperatures as high as 90 K.

<figure><img src="image_8.png" /><figcaption>**Figure 6.** Relation between the IRX and the observed UV slope $\beta$ for our observed galaxies (MASOSA in purple; VR7 in blue) and compilations of galaxies at $z \approx 5$ (yellow) and $z \approx 6-7$ (green). IR luminosity has been calculated assuming a dust temperature $T = 45$ K, and upper limits are shown at the $3\sigma$ level. The solid and dashed lines show the relations for the commonly used Calzetti and SMC extinction curves, respectively.</figcaption></figure>

With these caveats and considerations in mind, we convert our upper limits for VR7 to a limiting $L_{IR,T=45\,K} < 2.6 \times 10^{10} L_\odot$ and $L_{IR,T=45\,K} < 2.3 \times 10^{10} L_\odot$ ($1\sigma$) for MASOSA integrating an optically thin graybody between 8 and 1000 µm (see Table 2 and Appendix C for details). We assume a dust temperature of 45 K and a power-law exponent $\beta = 1.5$ and correct for the heating from the CMB following Ota et al. (2014), although note that these corrections are only a factor of $\approx 0.93$ and are less important for higher dust temperatures (see Appendix C for more details). The assumed dust temperature is similar to the measured dust temperature in local low-metallicity galaxy I Zw 18 (Rémy-Ruyer et al. 2015). Assuming a dust temperature of 35 K would result in $L_{IR,T=35\,K} < 1.3 \times 10^{10} L_\odot$ and $L_{IR,T=35\,K} < 1.2 \times 10^{10} L_\odot$ for VR7 and MASOSA, respectively.

As the observed UV continuum slope correlates with the level of dust attenuation (e.g., Meurer et al. 1999), it is expected that the IR/UV luminosity ratio (the IRX ratio) depends on the UV slope, with the relation being dependent on the shape of the attenuation curve and the dust geometry (e.g., Faisst et al. 2017) and the intrinsic stellar SED (Reddy et al. 2018). We show the locations of VR7 and MASOSA on the IRX-$\beta$ plane in Figure 6, together with our compilation of galaxies at $z \approx 6-7$, assuming a dust temperature of 45 K. VR7 and MASOSA both lie significantly below the commonly used SMC and/or Calzetti attenuation curves, which could indicate that the dust temperature is higher. Additionally, another explanation could be that dust particles reside closer to the star-forming regions and hence maximize the reddening at fixed IR luminosity (e.g., Ferrara et al. 2017), particularly in compact galaxies.

The constraints on the IR continuum can be converted to an upper limit on the level of obscured star formation. Following Kennicutt (1998), we constrain $SFR_{IR} < 4.6 M_\odot \text{yr}^{-1}$ and $SFR_{IR} < 4.0 M_\odot \text{yr}^{-1}$ for VR7 and MASOSA at $1\sigma$, respectively, for a dust temperature of 45 K. Additionally, we can use the observed UV continuum luminosity as an indicator of the unobscured SFR. Following again Kennicutt (1998), we measure $SFR_{UV} = 54^{+3}_{-2} M_\odot \text{yr}^{-1}$ for VR7 and $SFR_{UV} = 15 \pm 2 M_\odot \text{yr}^{-1}$ for MASOSA. If we were to use the classical method to correct the UV luminosity for dust attenuation based
<figure><img src="image_9.png" /><figcaption>**Figure 7.** Same as Figure 4, but now collapsing frequencies around the [C II] line with VR7 on the left and MASOSA on the right. Black contours show the ±3σ, 4σ, and 5σ levels, where σ = 8.8–12.7 mJy km s⁻¹ beam⁻¹. Frequencies are collapsed with a width of 200 km s⁻¹ around a central frequency that is blueshifted by 200 km s⁻¹ from the Lyα redshift, as motivated in Section 5. [C II] emission is clearly detected in VR7 (with peak S/N = 9.4 and resolved over multiple beams), while we do not detect [C II] emission from MASOSA (see Section 5.2). Note that the PSF FWHM of the HST imaging is a factor of ≈2.5 smaller than that of this ALMA reduction.</figcaption></figure>

<figure><img src="image_10.png" /></figure>

on the observed UV slope (e.g., Meurer et al. 1999), we
would find SFR<sub>dust,Meurer</sub> = 279<sup>+128</sup><sub>-118</sub> M<sub>☉</sub> and SFR<sub>dust,Meurer</sub> =
208<sup>+193</sup><sub>-171</sub> M<sub>☉</sub> yr<sup>-1</sup> for VR7 and MASOSA, respectively. Such
high SFRs are, however, ruled out by the nondetection of
continuum emission in our ALMA observations, which indicate
that the fraction of obscured SFR is ≲10%–30%.<sup>11</sup> Therefore,
combining the limiting obscured SFR as an upper bound
with the unobscured SFR, we find SFR<sub>UV+IR</sub> = 54<sup>+5</sup><sub>-2</sub> M<sub>☉</sub> yr<sup>-1</sup>
and SFR<sub>UV+IR</sub> = 15<sup>+4</sup><sub>-2</sub> M<sub>☉</sub> yr<sup>-1</sup> for VR7 and MASOSA,
respectively.

**5. Integrated [C II] Emission Properties**

We now focus on the ALMA data around the redshifted
[C II]158 µm frequency. While recent work showed that ALMA
resolves [C II] emission in multiple components for a large fraction
of high-redshift galaxies (e.g., Matthee et al. 2017b; Carniani et al.
2018b; Hashimoto et al. 2018), we focus first on the integrated
[C II] properties by using a low spatial resolution reduction of our
ALMA data (i.e., synthesized beam FWHM = 0.7) and analyze a
higher resolution reduction in Section 6.

In order to search for [C II] emission, we visually scanned
the data cube in slices of width 18 km s⁻¹ from −1000 to
+1000 km s⁻¹ with respect to the Lyα redshift and searched
for detections within a radius of 3" (15 pkpc at z ~ 6.5). We
illustrate the results of this scan in Figures 7 and 8. Figure 7
shows collapsed [C II] narrowband images zoomed at the
positions of VR7 and MASOSA, while Figure 8 shows 1D
spectra extracted at the source location. Normalized Lyα
spectra are shown on the same barycentric vacuum velocity

scale. We clearly detect [C II] emission in VR7, while no [C II]
is detected in MASOSA.

5.1. *Integrated [C II] Properties of VR7*

As shown in the left panels of Figures 7 and 8, the [C II]
emission that is detected in VR7 is roughly colocated with the
rest-frame UV emission and is slightly blueshifted with respect to
Lyα. The peak flux is detected with an S/N = 9.4 in the
collapsed [C II] image, while the integrated S/N in the 1D
extracted spectrum is ≈15.

We measure the redshift of peak [C II] emission and the FWHM of the [C II] line as follows. For 1000 iterations, we perturb each data point in the 1D extraction with its associated uncertainty and measure the frequency where the flux density is highest and the (linearly interpolated) frequencies where half of the peak flux density is observed. Then, we obtain the median and the 16th and 84th percentiles of those measurements, and we find that the peak [C II] flux is observed at z[C II] = 6.5285<sup>+0.0007</sup><sub>-0.0004</sub>. This means that the peak of the Lyα emission line is redshifted with respect to [C II] by Δv<sub>Lyα</sub> = 217<sup>+29</sup><sub>-19</sub> km s<sup>-1</sup>. The [C II] FWHM is v<sub>FWHM,[C II]</sub> = 200<sup>+48</sup><sub>-32</sub> km s<sup>-1</sup>, which is a factor of ≈1.7 smaller than the Lyα FWHM.

Next, we measure the integrated [C II] luminosity by collapsing
the primary-beam-corrected cube over a wider range of
frequencies (252.285–252.519 GHz, corresponding to −20 to
−320 km s⁻¹ w.r.t. Lyα; illustrated in Figure 8). We use the IMFIT
task in CASA, which fits an elliptical Gaussian profile to
the data. This results in a total flux of 447 ± 42 mJy km s⁻¹
and a deconvolved source size (major/minor axis) of
(1.45 ± 0.14) × (0.71 ± 0.08) with a position angle of
96.5 ± 5.2 ((1.58 ± 0.13) × (0.94 ± 0.06) before deconv-
lution). Following Carilli & Walter (2013), this flux translates into
a [C II] luminosity (4.76 ± 0.44) × 10⁸ L<sub>☉</sub>, which is emitted over
a region of (7.8 ± 0.8) × (3.8 ± 0.4) kpc². The integrated [C II]
measurements are summarized in Table 2. The [C II] luminosity,

<footnote n="11">11 We calculate that the obscured SFR would be similar to the unobscured SFR in the case in which the dust temperature is 65 K in MASOSA and 90 K in VR7. Compared to our fiducial calculations assuming T<sub>dust</sub> = 45 K, this would increase the upper limits on the IR luminosity by +0.4 and +1.05 dex, respectively.</footnote>
<figure><img src="image_11.png" /></figure>

<figure><img src="image_12.png" /><figcaption>**Figure 8.** Extracted 1D spectra centered on redshifted [C II] emission with VR7 on the left and MASOSA on the right. The black line shows the ALMA spectrum that is extracted at the rest-frame UV position (see Figure 7), while the blue line shows the arbitrarily normalized $Ly\alpha$ spectrum. For MASOSA we extract the spectrum within a circular aperture with diameter equal to the beam major axis, while for VR7 we use an aperture that is 1.5 times larger to include extended flux. The noise level of the ALMA data is indicated in gray. Data are binned by a factor of 2 in the frequency direction, resulting in a 38 km s⁻¹ velocity resolution. [C II] emission in VR7, detected with integrated S/N ≈ 15, is blueshifted with respect to $Ly\alpha$. The red shaded area highlights the frequency range used to create the integrated [C II] image. We do not detect [C II] emission in MASOSA within 2'' and ±1000 km s⁻¹ from $Ly\alpha$.</figcaption></figure>

line width, and velocity offset to $Ly\alpha$ are similar to other galaxies at $z \approx 6-7$ of comparable UV luminosity, and we will discuss this in more detail in Section 7.3. The major and minor axes of the [C II] emission are larger than the axes of the rest-frame UV emission, which are $(4.6 \pm 0.2) \times (2.1 \pm 0.1) \text{ kpc}^2$ when the UV is modeled as a single component. This indicates that [C II] is more extended, similar to observations in quasar hosts (Cicone et al. 2015) and a stack of star-forming galaxies (Fujimoto et al. 2019). Moreover, Rybak et al. (2019) find that extended [C II] emission may plausibly dominate the total [C II] emission in high-resolution observations of strongly star-forming galaxies at $z \sim 3$.

The line width and size can be converted to an estimate of the total dynamical mass following Wang et al. (2013), which results in $M_{dyn}/(\sin i)^2 = 1.4^{+0.4}_{-0.3} - 3.0^{+0.7}_{-0.6} \times 10^{10} M_{\odot}$ using the semiminor or semimajor axis, respectively. This mass is 1–2 times the stellar mass estimated in Matthee et al. (2017a).

## 5.2. A Physically Motivated Upper Limit for the [C II] Luminosity in MASOSA

With the presented observations, no [C II] emission is detected in the LAE MASOSA. Besides the [C II] image collapsed over the frequency range used in Figure 7, we have also inspected collapsed [C II] images ranging from −1000 to +1000 km s⁻¹ with respect to the $Ly\alpha$ redshift and with widths ranging from 50 to 500 km s⁻¹, but none reveal a detection within 2'' from MASOSA. We have also changed the reduction strategy including no UV tapering to a large UV taper of 200 kλ, but none of our attempts resulted in a detection.

We measure the sensitivity of our observations using the rms value after masking out the position of detected foreground objects. We measure an rms of 12.7 mJy km s⁻¹ beam⁻¹ in a collapse of width 200 km s⁻¹, which increases to 20.8 mJy km s⁻¹ beam⁻¹ collapsing channels over 400 km s⁻¹. The noise increases by a factor of ≈1.2 when a UV taper of 300 kλ is used instead of 650 kλ. As can be seen in Figure 8, the noise per channel is relatively invariant of frequency, meaning that the sensitivity does not depend on central frequency. These two sensitivities correspond to $1\sigma$ limiting [C II] luminosities of (0.14–0.22) × $10^8 L_{\odot}$.

Alternatively, we assess the significance of the nondetection of [C II] in MASOSA by simulating how fake sources with

known luminosity, source size, and width would appear in our data. We simulate point sources with a range of [C II] luminosities and line widths ranging from 100 to 400 km s⁻¹ in reductions with different beam sizes. We place these sources in 500 random positions in our collapsed images and measure the median peak S/N. For an object that is unresolved when imaged with a beam FWHM > 0.′′7 (corresponding to 3.8 kpc), we measure an S/N of 3 for a luminosity of (0.45, 0.65, 0.90) × $10^8 L_{\odot}$ for widths of (100, 200, 400) km s⁻¹, respectively. If an object is unresolved when imaged with a beam FWHM > 0.′′4, an S/N of 3 is measured for (0.38, 0.55, 0.78) × $10^8 L_{\odot}$ for each respective velocity width. Therefore, variations in the source size and line width significantly influence the luminosity for which [C II] emission can be detected.

We show in Figure 9 that the [C II] line FWHM is correlated with the UV luminosity (with a Spearman rank of 0.46, P-value 98.5% for $z \approx 5-7$ data on their own), particularly when combined with data at $z = 2-3$. This correlation resembles the SFR-stellar mass “main sequence,” as the [C II] line width is expected to be a tracer of gravitational potential and hence mass and the UV luminosity is a proxy for SFR. Increasing scatter in [C II] widths at the brightest UV luminosity could indicate dust attenuation, which could decrease the UV luminosity. MASOSA is among the galaxies with the lowest UV luminosity currently targeted by ALMA at $z \approx 6-7$ and is thus expected to have a [C II] line width of ≈120 ± 80 km s⁻¹. Moreover, in galaxies for which both [C II] and $Ly\alpha$ are detected, the [C II] width is typically a factor of ≈0.7 ± 0.3 smaller than the $Ly\alpha$ FWHM, likely because the $Ly\alpha$ line width is affected by radiative transfer effects in the ISM and CGM of galaxies. This would point toward a [C II] width of ≈250 km s⁻¹ for MASOSA.

The [C II] half-light radius observed in galaxies at $z \approx 6-7$ is typically ≈2.5 ± 1 kpc, a factor of ≈2 larger than the UV size of these galaxies measured with HST (Carniani et al. 2018b). As described above, the UV half-light radius of MASOSA is ≈1 kpc. Therefore, a [C II] half-light radius of about ≈2 kpc is expected for MASOSA, which corresponds to a major axis of ≈0.′′7.

Combining the expected [C II] extent and line width, the most realistic and conservative estimate of sensitivity of our
<figure><img src="image_13.png" /><figcaption>**Figure 9.** [C II] line FWHM vs. UV luminosity for VR7, compilations of galaxies at $z \approx 5$ and $z \approx 6-7$, and a sample of LAEs at $z \approx 2-3$ (where we use the line width of rest-frame optical nebular lines; Erb et al. 2014). UV luminosity, which traces unobscured star formation, is correlated to [C II] width, which traces gravitational potential. Therefore, this relation resembles the “main sequence” between SFR and stellar mass. We highlight the UV luminosity and its uncertainty from MASOSA with a horizontal gray line, indicating that a [C II] FWHM of $\approx 120 \pm 80$ km s$^{-1}$ is expected based on the UV luminosity.</figcaption></figure>

observations is obtained by assuming a line width of 200 km s$^{-1}$ and using a reduction with beam FWHM $\approx 0.7''$. This implies a $3\sigma$ limit for the [C II] luminosity from MASOSA of $0.65 \times 10^8$ L$_{\odot}$. We compare this upper limit with expectations from the UV luminosity and its implications in Section 7.1. For comparison, this value corresponds to $0.3 \times L_{[C II]}^*$ in the local universe (Hemmati et al. 2017).

## 6. The Resolved Structure of VR7

Thanks to the high S/N of the detected [C II] line in VR7, we can reduce the ALMA data with a higher spatial resolution and/or relax the velocity averaging compared to our analysis in Section 5. In this section, we first investigate the detailed 1D line profile extracted over the full area over which [C II] emission is detected and then focus on resolving structure spatially in both rest-frame UV and [C II] emission.

### 6.1. 1D Line Modeling

In Figure 10 we show the 1D line profile of VR7's [C II] line extracted with the native 9 km s$^{-1}$ resolution and extracted over the full [C II] spatial extent. The line profile appears asymmetrically skewed toward the red and is slightly better described by a two-component Gaussian ($\chi_{red}^2 = 0.59$) than a single Gaussian component ($\chi_{red}^2 = 0.69$). The best-fitted values of the two-component Gaussian fit are listed in Table 3. The line profile is characterized by a relatively broad (FWHM = $145_{-33}^{+41}$ km s$^{-1}$), luminous component and a narrower (FWHM = $93_{-24}^{+26}$ km s$^{-1}$) component that is a factor of $\approx 3$ fainter. The narrow component is redshifted by $124_{-17}^{+19}$ compared to the more luminous component, but still blueshifted by $\approx 90$ km s$^{-1}$ compared to the peak of $Ly\alpha$ emission.

### 6.2. Two Separate UV Components

VR7's rest-frame UV continuum emission is well resolved in the HST/WFC3 image, with a maximum extent of about $1''_1$ ($\approx 6$ kpc) within the $2\sigma$ contours. While in general the galaxy is

<figure><img src="image_14.png" /><figcaption>**Figure 10.** 1D spectrum of the [C II] line in VR7, as in the left panel of Figure 8, but without channel averaging, resulting in a $\approx 9$ km s$^{-1}$ resolution. The black line shows the measurement, while the gray region illustrates the noise level. We fit the skewed line profile with a single Gaussian component (red) and a combination of two Gaussian components (blue). The bottom panel shows the residuals of the best fits. The data marginally prefer a two-component fit, indicating resolved velocity structure.</figcaption></figure>

<table><caption>**Table 3**

Best-fitted Parameters to the Two-component Gaussian Fits of the 1D [C II] Line Profile of VR7 (Shown as the Solid Blue Line in Figure 10)</caption><thead><tr><th>Property ([C II])</th><th>Value</th></tr></thead><tbody><tr><td>Flux<sub>1</sub> (mJy km s<sup>-1</sup> beam<sup>-1</sup>)</td><td>139.8<sup>+29.0</sup><sub>-30.3</sub></td></tr><tr><td>Flux<sub>2</sub> (mJy km s<sup>-1</sup> beam<sup>-1</sup>)</td><td>43.1<sup>+22.9</sup><sub>-21.6</sub></td></tr><tr><td>Flux<sub>2</sub>/(Flux<sub>1</sub>+Flux<sub>2</sub>)</td><td>0.24<sup>+0.15</sup><sub>-0.12</sub></td></tr><tr><td>Δv<sub>Lyα,1</sub> (km s<sup>-1</sup>)</td><td>-211<sup>+11</sup><sub>-9</sub></td></tr><tr><td>Δv<sub>Lyα,2</sub> (km s<sup>-1</sup>)</td><td>-87<sup>+15</sup><sub>-15</sub></td></tr><tr><td>FWHM<sub>1</sub> (km s<sup>-1</sup>)</td><td>145<sup>+41</sup><sub>-33</sub></td></tr><tr><td>FWHM<sub>2</sub> (km s<sup>-1</sup>)</td><td>93<sup>+26</sup><sub>-14</sub></td></tr></tbody></table>

elongated in the east-west direction, it also has a tail extending somewhat to the southwest (see the left panel of Figure 11). Because of this tail, the rest-frame UV light profile is better fitted by a two-component exponential model than a single-component model (with reduced $\chi_r^2 = 1.10$ vs. $\chi_r^2 = 2.83$, detailed in Appendix B). The best-fitted double exponential model contains two components separated by $0.35 \pm 0.01''$ (1.88 $\pm$ 0.07 kpc). We show the individual components in the middle and right panels of Figure 11. The image for component 1 is created by subtracting the PSF-convolved model of component 2, and vice versa. We list their properties in Table 4. Component 1 lies in the east and is relatively compact, while component 2 coincides with the peak UV emission and is elongated in the southwestern direction. Component 2 is slightly larger and a factor of $1.8_{-0.5}^{+0.7}$ more luminous, implying that both clumps have very similar SFR surface density. We note that we refrain from overinterpreting the colors of modeled components because of $Ly\alpha$ emission possibly affecting the morphology in the F110W filter, but we refer the reader to Appendix B for more details.

### 6.3. Spatially Resolved [C II] and UV Emission

Now, we focus on spatially resolving the [C II] emission using a high-resolution reduction ($0.47 \times 0.40$ beam;
<figure><img src="image_15.png" /><figcaption>**Figure 11.** Zoomed-in HST/WFC3 image of VR7 in the F160W filter. The left panel shows the data and the size of the PSF FWHM. The middle and right panels show the individual components of the best-fitted model of two individual exponential components. The contours are drawn in the left panel and shown in the other panels for reference. Each panel lists the effective radius of the respective component, modeled as a single exponential light profile.</figcaption></figure>

<table><caption>**Table 4**

Best-fitted Parameters for the Two-component Exponential Models of the HST/WFC3 F160W Image of VR7 (Shown in Figure 11)</caption><thead><tr><th>Property</th><th>Component 1</th><th>Component 2</th></tr></thead><tbody><tr><td>ΔR.A.</td><td>+0.263</td><td>-0.079</td></tr><tr><td>Δdecl.</td><td>+0.104</td><td>+0.126</td></tr><tr><td>r<sub>eff</sub></td><td>0.84 ± 0.11 kpc</td><td>1.12 ± 0.06 kpc</td></tr><tr><td>PA</td><td>76 ± 9°</td><td>48 ± 4°</td></tr><tr><td>Ellipticity</td><td>0.52 ± 0.12</td><td>0.62 ± 0.08</td></tr><tr><td>Fraction of total flux</td><td>36 ± 7%</td><td>64 ± 7%</td></tr></tbody></table>

**Note.** ΔR.A. and Δdecl. are the relative positions with respect to the center listed in Table 2.

2.5 kpc × 2.2 kpc); see Figure 12. Extended [C II] emission is observed over an area of ≈1.2′′3 × 0.6′′6, with the largest elongation in the east-west direction, similar to what was found above in a lower-resolution reduction. Interestingly, higher surface brightness contours are rotated by about 45° clockwise, resembling component 2 in the rest-frame UV. Besides the peak [C II] emission in the west, a hint of a second peak is seen in the east at ≈−0.5′′5 from the image center. Figure 12 also compares the [C II] map with the rest-frame UV components identified in the HST/WFC3 F160W imaging (Figure 11). In general, the [C II] morphology is similar to the rest-frame UV, with the major elongation in the east-west direction following the centers of both components. Intriguingly, the peak [C II] emission extends toward the southwest, which is similar to the major direction of elongation of component 2 in the HST data.<fnref n="12" />

## 6.4. Resolved UV/[C II] Ratio

As a result of the differences in the UV and [C II] morphologies, the UV/[C II] ratio differs within the galaxy. We quantify the variation of the UV/[C II] ratio by converting the luminosities to (unobscured) SFR surface densities following Kennicutt (1998) for the UV and the observed

<figure><img src="image_16.png" /><figcaption>**Figure 12.** Zoomed-in map of the VR7 galaxy. [C II] image from our high-resolution ALMA reduction (beam axes 0.47 × 0.40), where images are oriented such that east is to the left and north is up. Black lines indicate the ±3σ, 4σ, 5σ, and 6σ contour levels. Colored dashed contours show the individual components in the rest-frame UV identified in HST/WFC3 imaging, with a PSF FWHM matched to the ALMA beam by convolving with a 2D asymmetric Gaussian. White and black crosses highlight the position of peak [C II] emission at the core velocity (±40 km s⁻¹ compared to z = 6.5285) and at the velocity of the redshifted shell (+100 to +160 km s⁻¹), respectively. Like the UV continuum, [C II] extends in the east-west direction following the centers of both components. Higher surface brightness [C II] emission is extended to the southwest with a position angle of ≈−45°, remarkably similar to component 2 in the rest-frame UV. We note that the relative astrometry may have uncertainties on the order of ≈0.1′′1−0.2′′2.</figcaption></figure>

relation between [C II] luminosity and SFR in the local universe (De Looze et al. 2014) and illustrate the results in Figure 13.

The average SFR<sub>UV</sub>/SFR<sub>[C II]</sub> ratio in the region where both UV and [C II] are detected at 3σ is normal, with a mean ratio of 1.3 ± 0.1, but the scatter on ~2 kpc scales is large (0.7). For example, the lowest SFR<sub>UV</sub>/SFR<sub>[C II]</sub> ratio, observed in the southwestern part of the galaxy, is ≈0.4, while the highest ratio of about ≈2.8 is only 2.8 kpc away. The region with peak [C II]

<footnote n="12">¹² The relative astrometries of the ALMA and HST data have been verified with only a single foreground object, meaning that uncertainties on the order of ≈0.1′′1−0.2′′2 exist. However, the result that the position angles of high surface brightness regions of [C II] and UV emission do not align is robust to these uncertainties.</footnote>
<figure><img src="image_17.png" /><figcaption>**Figure 13.** Spatial variation of the SFR<sub>UV</sub>/SFR<sub>[C II]</sub> ratio, where the first is not dust corrected and the latter is estimated from the scaling relation in the local universe by De Looze et al. (2014). HST data have been smoothed to match the ALMA beam. The gray horizontally hatched region shows where no flux is detected in the rest-UV at >3σ, while the cross-hatches highlight where no [C II] flux is detected at >3σ. The dotted black-and-white line shows the [C II] level as in Figure 12.</figcaption></figure>

emission has a relatively low UV luminosity, potentially indicating obscured star formation (although no dust is detected in our ALMA observations; see Section 4) and/or a higher metallicity. We note that the UV colors of component 2 also indicate that it is older and more evolved (see Appendix B), although we stress that these colors need to be interpreted cautiously. The central-eastern part of the galaxy has a relatively high UV/[C II] ratio, indicating little obscuration and a low metallicity. No UV emission is detected at the position of the tentative [C II]-emitting component in the southeast. The outskirts of the galaxy tend to have a relatively high [C II] luminosity, indicating a strong contribution from PDR emission instead of [C II] emission originating from H II regions. This is consistent with measurements in local low-metallicity galaxies (Cormier et al. 2019) and the Milky Way (Pineda et al. 2013). Most observed [C II] emission originates from PDR regions (~50%; Pineda et al. 2013), but with other significant contributions from molecular hydrogen and cold atomic gas. Only a small fraction of the [C II] (which decreases with decreasing metallicity; Cormier et al. 2019) originates from H II regions. However, note that local galaxies are not affected by a strong CMB background that is present at high redshift.

## 6.5. Velocity Structure

### 6.5.1. Moment Maps

Here we explore the velocity structure of VR7 using the first-, second-, and third-moment maps of the [C II] emission in the high-resolution reduction, which can be interpreted as maps showing the peak velocity, velocity width, and line skewness, respectively; see Figure 14.<fnref n="13" />

The first-moment map (left panel of Figure 14) reveals that the peak velocity varies along the 45° SW axis, the same position angle as component 2. While this pattern somewhat resembles rotation, the maximum velocity difference is only $\Delta v \approx 40$ km s⁻¹, which, combined with the velocity dispersion of $\sigma \approx 65$ km s⁻¹, results in a kinematic ratio $\Delta v/2\sigma \approx 0.3$. This indicates a dispersion-dominated system (e.g., Förster Schreiber et al. 2009), unlike two luminous LBGs observed with ALMA in lower spatial resolution by Smit et al. (2018).

The second-moment map (middle panel of Figure 14) shows variations in the velocity dispersion within a factor of ≈1.5. The region with highest flux density has relatively low dispersion, while the northeast region has slightly broader line emission. It should be noted that the lowest dispersions in the second-moment map are on the order of $\sigma \approx 60$ km s⁻¹, which is similar to the line width of the dominant (broader) component identified in our fit to the 1D spectrum. This could indicate that the second, narrow and slightly redshifted component identified from the 1D spectrum (Figure 10) originates mostly from the northeastern region of the galaxy. This is supported by the line skewness map (right panel of Figure 14). While the highest skewness values that are found in regions where [C II] is detected at $\approx 4\sigma$ significance should be interpreted with caution, the skewness still increases toward the northeast in the regions with high S/N.

### 6.5.2. Position–Velocity Diagrams

We also analyze position–velocity (PV) diagrams in order to study how the [C II] line varies within VR7. We construct PV diagrams by averaging over a slit width of 1′′5 around the center in the N–S direction (PA = 90°) using the CASA task IMPV. Figure 15 shows the results for a reduction optimized for high spectral resolution (top row) and high spatial resolution (bottom row), where the first is achieved by reducing the data with normal weighting and UV tapering, while the latter is achieved with Briggs weighting, but averaging over two velocity channels.

Figure 15 confirms the result from the 1D analysis, which is that the [C II] line profile is composed of (at least) two components: a central dominant component that is relatively compact and a redshifted extended component that has a narrower line width, which we name the *shell* in the rest of the paper. The significance of the *shell* is >4$\sigma$ in the reduction with low spatial resolution, but lower in the reduction with higher spatial resolution, which indicates that the [C II] emission is somewhat diffuse and resolved out in the latter reduction. The *shell* is redshifted by $\approx 100$ km s⁻¹ compared to the central component and is therefore consistent with the narrow Gaussian component identified from the integrated 1D spectrum (see Table 3). As the *shell* is chemically enriched, it is more likely to be outflowing than inflowing. While the flux from the *shell* extends over roughly the same spatial scales as the central component, the center is shifted slightly toward the northeast, consistent with the skewness map in Figure 14. It is challenging to simultaneously show collapsed images of the core [C II] emission and the *shell* because of the small separation and the extended nature of the *shell*. The positions of peak [C II] flux are therefore indicated with crosses in Figure 12, which shows that the separation is $\approx 0.′′35$, or $\approx 2$ kpc.

Isolating the dominant component in the bottom panel of Figure 15, we measure a size FWHM ranging from 0.′′35 to 0.′′50. As the beam major axis is 0.′′46, this means that the central component is unresolved and has half-light radius

<footnote n="13">¹³ We note that the simple interpretation of the first and second moments may be somewhat confused by strong line skewness. For example, the first moment of a line-emitting region with fixed central velocity throughout the system, but with a red asymmetry in only part of the system, will show a slight gradient following the skewness gradient, as a red skewed line will result in a higher first moment.</footnote>
<figure><img src="image_18.png" /><figcaption>**Figure 14.** First- (left), second- (middle), and third-moment (right) maps of VR7, based on the highest spatial resolution ALMA reduction and using frequencies within [-100, +200] km s⁻¹ with respect to the [C II] redshift z = 6.5285. These moments correspond to the velocity map, the dispersion map, and the skewness map, respectively. The moment maps show complex structure and suggest that VR7 is dispersion dominated with a typical Δv/2σ ≈ 0.3 in the highest-S/N regions.</figcaption></figure>

≈1.3 kpc. Combined with the line width of FWHM = 145±41
km s⁻¹, this results in an estimate of the dynamical mass
M<sub>dyn</sub>/(sin i)<sup>2</sup> ≲ (5.5 ± 3) × 10<sup>9</sup> M<sub>☉</sub> (Wang et al. 2013;
Matthee et al. 2017b). Besides the dominant component, the
bottom panel of Figure 15 also hints that there is a second
fainter compact [C II]-emitting component at an offset of about
4 kpc ≈0.7 to the east (see also Figure 12, which shows that
this component overlaps with one of the UV components). The
peak flux of this component is a factor of 2 lower than the
dominant component and is detected at only ≈3σ-4σ
significance.

7. Discussion

We now focus on interpreting the ALMA and HST
observations and place them into the context of other
observations. We first address the nondetections of [C II] and
IR continuum in MASOSA and discuss ways to reconcile this
with the potentially red color in the UV (Section 7.1). Second,
we discuss the nature of VR7 based on current data in
Section 7.2, and we finally discuss the relation between [C II]
and UV emission in Section 7.3.

7.1. *The Nature of MASOSA: Dusty? Metal-poor?*

MASOSA is not detected in [C II] or FIR continuum
emission. This may be explained by a higher dust temperature
combined with a low dust mass and a relatively low metallicity.
Such properties are indeed expected given MASOSA's high
redshift and very high Lyα EW (typically found in young,
metal-poor, and dust-free galaxies; e.g., Sobral et al.
2015, 2018; Nakajima et al. 2016; Trainor et al. 2016).

As illustrated in Figure 16, the [C II] luminosity limits imply
that MASOSA is significantly [C II] deficient given its SFR,
compared to the typical [C II] luminosity found for star-forming
galaxies in the local universe (e.g., De Looze et al. 2014). This
can be explained by a lower gas-phase metallicity (e.g., Vallini
et al. 2015; Olsen et al. 2017). Following the relations derived
by Vallini et al. (2015) based on radiative hydrodynamical
simulations, the [C II] upper limit in MASOSA implies a gas-
phase metallicity ≲0.07 Z<sub>☉</sub> (see Figure 16), providing further
evidence for the very low metallicity of this source as discussed
in Sobral et al. (2015). An additional rough metallicity estimate
can be derived based on *Spitzer*/IRAC photometry. For the
redshift of MASOSA, the [3.6] band is contaminated by Hβ
+[O III] line emission, while the [4.5] band includes Hα
emission (e.g., Schaerer & de Barros 2009; Raiter et al. 2010a).
Assuming a flat continuum in the rest-frame optical, the

MASOSA blue IRAC color ([3.6]–[4.5] = −0.6 ± 0.2) indi-
cates an EW(Hβ+[O III]) ≈ 1500 (e.g., Smit et al. 2014) and
hence a gas-phase metallicity of ≈0.01–0.4 Z⊙ (Castellano
et al. 2017). More detailed constraints on the gas-phase metal
abundances in the ISM of MASOSA can be obtained from
future deep (near-)infrared spectroscopy.

MASOSA's possible red UV slope, $\beta = -1.06_{-0.72}^{+0.68}$, is relatively unexpected given the UV luminosity ($\beta \approx -2.1$; e.g., Bouwens et al. 2012 and Figure 2), the high Ly$\alpha$ EW (Hashimoto et al. 2017), the blue IRAC colors, and the nondetection of dust continuum. Besides dust attenuation, which we discussed in Section 4,<fnref n="14" /> a red UV slope could be reconciled with a young age in the case of a combination of strong nebular continuum emission, strong UV emission lines besides Ly$\alpha$, and/or multiple clumps.

For example, Raiter et al. (2010b) show that strong nebular
continuum free-free emission can redden the observed UV
slope of low-metallicity galaxies with hard ionizing spectra up
to $\beta \approx -1.5$ at $\lambda_0 = 1500$ Å. The most likely emission line
boosting the F160W magnitude (and hence the observed $\beta$) is
C III]1909. C III]1909 is the strongest UV emission line in star-
forming galaxies besides Ly$\alpha$ (e.g., Shapley et al. 2003) and
can have EW$_0 \approx 10-20$ Å (e.g., Rigby et al. 2015; Stark et al.
2015; Maseda et al. 2017), particularly in galaxies with strong
Ly$\alpha$ emission (e.g., Le Fèvre et al. 2019). However, we
calculate that correcting for a C III] EW of 25 Å would imply a
slope of $\beta = -1.2 \pm 0.7$, meaning that exceptionally strong
lines would be required (see also Stroe et al. 2017). Moreover,
MASOSA has a slightly different morphology in the reddest
HST band (e.g., Figure 3 and Figure 19), indicating the
presence of a second, relatively faint red clump. Such a clump
may be physically disconnected from the youngest stellar
population. By fitting the F160W data with a combination of
two point sources,<fnref n="15" /> we find that the red clump contributes
≈20% of the flux. Reducing the F160W flux by 20% results in
$\beta = -1.5 \pm 0.6$, meaning that the main component of
MASOSA can plausibly be “normally” blue.

Concluding, MASOSA has common properties for galaxies
with strong Lyα emission (relatively compact, low metallicity
and dust content) and is likely a young starbursting galaxy,
with little evidence for being powered by an AGN. While

<footnote n="14">¹⁴ Any significant amount of dust obscuration in MASOSA would also imply a lower metallicity in order to explain the nondetection of the [C II] line.</footnote>

<footnote n="15">¹⁵ The F160W data of MASOSA are described equally well by a single component within the uncertainties (see Appendix B), but we use the results from a two-component fit to address the implication of this hypothesis.</footnote>
<figure><img src="image_19.png" /><figcaption>**Figure 15.** PV diagrams of VR7. The top row uses an ALMA reduction optimized for high spectral resolution but with low spatial resolution (natural weight, UV taper; as in Figure 7). The bottom row uses a high spatial resolution reduction (Briggs weight R = 0.5; as in Figure 12) but averages over two spectral elements. Flux is averaged over the north-south direction over a 1"/5 width. Note that the y-axis scale differs between the top and bottom rows. The top row shows that the [C II] line profile can be resolved in a dominant central component and a faint, diffuse component *shell* at a redshift of ≈100 km s⁻¹. This diffuse component is likely resolved out and at lower significance in the high-resolution reduction shown in the bottom row.</figcaption></figure>

MASOSA's red UV color needs to be verified with deeper
observations, a combination of a possible second component,
relatively strong rest-frame UV lines such as C III], and strong
nebular continuum emission could reconcile MASOSA's

relatively red color with being a young and relatively dust-
free galaxy.

7.2. *The Nature of VR7: Hot Dust? Merger?*

VR7 is a large, UV-luminous star-forming galaxy with relatively strong Lyα emission and relatively normal [C II] emission. Remarkably, VR7 is the most UV-luminous SFG at z ≈ 6–7 for which no dust emission is detected at 160 µm. The limiting 160 µm/1500 Å flux density ratio is extremely low and in the local universe only seen for extremely metal-poor systems as I Zw 18 (Figure 5). The Lyα EW implies an Lyα escape fraction of ≈15%–20% (Sobral & Matthee 2019), which indicates that a significant fraction of Lyα emission is destroyed by dust. Moreover, the moderate color β ≈ −1.4 also indicates dust attenuation. Therefore, it is likely that dust is present in VR7 at high temperatures (T > 60 K), which would also reconcile the limits on the 160 µm continuum flux density with common IRX-β relations (Figure 6 and, e.g., Faisst et al. 2017). The nondetection of IR continuum emission, the narrowness of the [C II] line, and the UV morphology all indicate that VR7 is not powered by an AGN (e.g., Ota et al. 2014).

In Section 6, we showed that the [C II] and rest-frame UV
images clearly resolve the light from VR7, with clear
indications of multiple components even though they are
separated by only ≈0.′′3–0.′′4. This separation is a factor of 2–3
smaller than the separation of individual components in other
luminous Lyα emitters at z ≈ 6.5 (e.g., Matthee et al. 2017b;
Carniani et al. 2018b). Whether the components that are
building up the VR7 galaxy are already at a smaller physical
separation compared to other luminous LAEs (i.e., VR7 is in a
more advanced stage of a merger) or this is a viewing angle/
inclination effect cannot be addressed with the current data.
Interestingly, among the luminous LAEs targeted with ALMA
so far, VR7 has the lowest Lyα EW, which could indicate that
it is in a slightly more evolved stage than LAEs with higher
Lyα EW. A slightly more evolved stage is also in line with its
moderate UV slope, particularly since the other luminous LAEs
have β ≈ −2.

7.3. What Determines the [C II]/UV Ratio

Here we address the relative [C II] and UV luminosities of
galaxies at z ≈ 6–7, compared to galaxies in the local universe.
Earlier work found relatively high (Capak et al. 2015),
moderately low (Pentericci et al. 2016), or extremely low
[C II] luminosities at fixed SFR<sub>UV</sub> (e.g., Ota et al. 2014). As
detailed in Appendix C, we homogenize the measurements by
converting all UV luminosities to SFR<sub>UV</sub> consistently, and we
revise upper limits following the method applied to MASOSA.
Our results are compared to relations between [C II] luminosity
and SFR based on observations of galaxies in the local universe
(De Looze et al. 2014). Like our observations, we rescale the
relation to the common SFR calibration of Kennicutt (1998).
The results are shown in Figure 16, where the left panel uses
SFR estimated from only the UV, while the right panel uses the
information on the dust continuum provided by ALMA to
correct for obscured SFR.

While VR7 has an integrated [C II]/SFR<sub>UV</sub> ratio that is very
similar to that of other galaxies with similar SFR,<sup>16</sup> MASOSA

<footnote n="16">¹⁶ We note that the global SFR<sub>UV</sub> and [C II] surface densities of VR7 also agree with galaxies in the local universe observed by De Looze et al. (2014).</footnote>
<figure><img src="image_20.png" /></figure>

**Figure 16.** Relation between SFR<sub>UV</sub> and [C II] luminosity at z ≈ 6–7. The left panel shows SFR<sub>UV</sub> estimated from the rest-frame UV luminosity following Kennicutt (1998). The right panel shows SFR<sub>UV+IR</sub>, where we have converted IR continuum detections to SFR<sub>IR</sub> assuming a dust temperature of 45 K and assign the 1σ limiting SFR<sub>IR</sub> as an upper error on the SFR as in Section 4.2. We compare our integrated luminosity measurement for VR7 (blue pentagon) and the upper limit for MASOSA (purple triangle) to other galaxies observed at z ≈ 5 and z ≈ 6–7 (yellow green diamonds; all computed consistently; see Appendix C) and to the relation observed in the local universe (De Looze et al. 2014; blue band; where the shaded region shows the dispersion, the line style changes to dashed at the luminosities where relations are extrapolated, and the relation is shifted to the same IMF as used here). Upper limits are at the 1σ level. We indicate galaxies with strong Lyα lines with red stars. Dotted lines show the relation between [C II] luminosity and SFR at different gas-phase metallicities inferred from hydrodynamical simulations by Vallini et al. (2015), which used the same IMF as our analysis.

is [C II] deficient, similar to several other galaxies with SFR<sub>UV</sub> ≈ 10–20 M<sub>☉</sub> yr<sup>-1</sup>. There is significant dispersion in [C II] luminosities at fixed SFR<sub>UV</sub>, particularly in the range SFR ≈ 30–60 M<sub>☉</sub> yr<sup>-1</sup>, which includes most of the objects currently observed. A part of this dispersion is related to variation in relative strengths of Lyα emission, since the SFR/[C II] ratio is anticorrelated with the observed Lyα EW (e.g., Carniani et al. 2018b; Harikane et al. 2018). Indeed, a large fraction of objects with [C II] deficits are strong Lyα emitters (Figure 16). Furthermore, among the sample of luminous LAEs known at z ≈ 6–7, MASOSA is at the higher end of the Lyα EW distribution (e.g., Figure 1), in agreement with this picture. On the other hand, Figure 16 shows that several other strong Lyα emitters, such as CR7 (Matthee et al. 2017b), have rather typical [C II] luminosities given the SFR.

In the range SFR<sub>UV+IR</sub> ≈ 30–200 M<sub>☉</sub> yr<sup>-1</sup>, the detected galaxies at z ≈ 6–7 are roughly colocated with galaxies in the local universe and follow a similar slope log<sub>10</sub>(L<sub>[C II]</sub>) ∝ log<sub>10</sub>(SFR) once obscured SFR is accounted for. At SFRs below 30 M<sub>☉</sub> yr<sup>-1</sup> the results are less clear, as the [C II] line is undetected in a large fraction of galaxies, with a 1σ limiting luminosity well below the local relation.

If relatively faint [C II] lines are confirmed in these systems, it would imply that the relation between [C II] luminosity and SFR steepens at low SFRs and that a larger fraction of the ISM is ionized. This potentially indicates a stronger correlation between gas-phase metallicity and SFR in lower-mass galaxies. Alternatively, it could also indicate more bursty star formation in fainter galaxies (leading to a higher ionization parameter) or other differences in ISM properties. Such a steepening is observed in the local universe using a sample of low-metallicity dwarf galaxies (De Looze et al. 2014). Furthermore, such steepening would also explain why relatively low [C II] luminosities are found among most strong Lyα emitters, as the fraction of galaxies with strong Lyα emission is higher at low SFRs in the observed sample (Figure 16). This picture simultaneously allows higher [C II] luminosities in the most

luminous Lyα emitters found in the large wide-field surveys (e.g., Matthee et al. 2017b; Carniani et al. 2018b).

## 8. Conclusions

In this paper, we have presented new deep follow-up observations of two luminous Lyα emitters at z ≈ 6.5 with HST/WFC3 and ALMA with a resolution of 1.5–2 kpc. The HST data are used to characterize the strength and morphology of rest-frame UV emission, constraining the unobscured SFR. ALMA data are used to constrain the level of FIR dust continuum emission and measure the strength and dynamics of the atomic [C II]158 µm fine-structure line. The targeted galaxies, VR7 and MASOSA (Sobral et al. 2015; Matthee et al. 2017a), have a high ≈ 2 × L* Lyα luminosity, while their UV luminosity differs by a factor of 4. Table 2 summarizes our measurements. Our main results are the following:

1. VR7 is among the most luminous and largest star-forming galaxies known at z ≈ 6–7, with an unobscured SFR = 54<sub>-2</sub><sup>+3</sup> M<sub>☉</sub> yr<sup>-1</sup> and an extent of ≈ 6 kpc at the 2σ level in the UV. VR7 has a typical moderate color given its high UV luminosity, β = -1.38<sub>-0.27</sub><sup>+0.29</sup>, and moderate Lyα EW<sub>0</sub> = 34 ± 4 Å. The HST/WFC3 rest-frame UV data of VR7 strongly prefer a two-component exponential model over a single-component model. The individual components are separated by 1.88 ± 0.07 kpc and have sizes r<sub>eff</sub> = 0.84 ± 0.11 kpc and r<sub>eff</sub> = 1.12 ± 0.06 kpc. The larger component is a factor of 1.8<sub>-0.5</sub><sup>+0.7</sup> more luminous and is elongated in the southwestern direction.

2. MASOSA, which has a ≈ L* UV luminosity (unobscured SFR = 15<sub>-2</sub><sup>+2</sup> M<sub>☉</sub> yr<sup>-1</sup>) and a normal size (r<sub>eff</sub> = 1.12<sub>-0.19</sub><sup>+0.44</sup> kpc), is among the galaxies with the highest Lyα EW known at z ≈ 6–7 (EW = 145<sub>-43</sub><sup>+50</sup> Å). Surprisingly, MASOSA seems to have a relatively red UV continuum color β = -1.06<sub>-0.72</sub><sup>+0.68</sup>, which could indicate strong C III line emission, nebular continuum emission, and/or a multiple component nature.
3. No dust continuum emission is detected in either VR7 or MASOSA (Section 4.2). Unless the dust temperature is very high ($T_{\text{dust}} \gtrsim 60$ K), this indicates little amounts of dust present in both systems. VR7 has an extremely low $L_{\text{IR}}/L_{\text{UV}}$ ratio, similar to the ratio of the extremely metal-poor dwarf galaxy I Zw 18 and the luminous LAE CR7 (Matthee et al. 2017b). Assuming $T_{\text{dust}} = 45$ K, we find that the obscured SFR in both galaxies is a small fraction of the unobscured SFR ($\lesssim 10\%$ and $\lesssim 30\%$, respectively).

4. Our ALMA data reveal a strong [C II] emission detection in VR7 (S/N = 15), but none in MASOSA (Section 5). VR7 has a typical integrated [C II] luminosity relative to its SFR, indicating a metallicity of ≈0.2 $Z_{\odot}$. For MASOSA, we derive a physically motivated [C II] upper limit that results in a [C II] deficit given the SFR$_{UV}$ and that implies a very low metallicity of <0.07 $Z_{\odot}$. Combining our measurements with a large compilation, we find indications that the relation between [C II] luminosity and SFR steepens at SFRs $\lesssim 30\ M_{\odot}$ yr$^{-1}$, which could be explained by a tighter correlation between SFR and metallicity in these galaxies compared to more luminous systems. Such a correlation naturally explains why most galaxies with high Ly$\alpha$ EW have relatively low [C II] luminosity.

5. [C II] emission in VR7 is extended (major/minor axes ($7.8 \pm 0.8$) × ($3.8 \pm 0.4$) kpc$^2$) a factor of ≈1.8 larger than the UV emission. The faint and low surface brightness [C II] emission extends mostly in the east-west direction. The high surface brightness [C II] emission also extends 45° toward the southwest (Section 6). We find that the UV/[C II] ratio varies significantly within the galaxy (up to factors of ≈7 on ≈3 kpc scales). The dominant [C II]-emitting component is compact, with an estimated half-light radius $\lesssim 1.3$ kpc and being slightly offset from the peak UV emission. We find that some faint, diffuse [C II] emission is redshifted by +100 km s$^{-1}$ with respect to the central component, possibly indicating an outflow and explaining the skewness and asymmetry seen in the 1D spectrum. Finally, we find indications for a second compact component at a 4 kpc separation.

Our observations show that there is little dust present in luminous Ly$\alpha$ emitters at $z \approx 6-7$, unless the dust has very high temperatures ($\gtrsim 60$ K). Future ALMA surveys should therefore aim at observationally constraining these temperatures by performing deep observations at higher frequencies. There are large variations in the [C II] luminosities of galaxies at $z \approx 6-7$ that are strongly related to the SFR$_{UV}$, albeit with significant scatter. Deep ALMA observations of galaxies with SFR$_{UV} < 10\ M_{\odot}$ yr$^{-1}$ ($M_{1500} \gtrsim -20$) should test whether the relation between [C II] luminosity and SFR$_{UV}$ indeed steepens at low SFRs. ALMA and HST are capable of resolving the most luminous systems at $z \approx 6-7$ and reveal that they are actively assembling from multiple components with strong variations in the SFR$_{UV}$/[C II] ratio, while their environment is already being influenced by galactic outflows. In the future, resolved rest-frame optical spectroscopy using integral field unit spectroscopy on JWST will provide crucial complementary measurements of the ISM metallicity, the contribution from

nonthermal emission, and the masses, ages, and metallicities of stellar populations.

We thank the anonymous referee for constructive comments and suggestions. We thank Max Gronke for comments on an earlier version of this paper. L.V. acknowledges funding from the European Union's Horizon 2020 research and innovation program under the Marie Skłodowska-Curie grant agreement No. 746119. This paper makes use of the following ALMA data: ADS/JAO.ALMA#2017.1.01451.S. ALMA is a partnership of ESO (representing its member states), NSF (USA), and NINS (Japan), together with NRC (Canada), NSC and ASIAA (Taiwan), and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO, and NAOJ. Based on observations obtained with the Very Large Telescope, programs 294.A-5018, 097.A-0943, and 99.A-0462. Based on observations made with the NASA/ESA Hubble Space Telescope, obtained (from the Data Archive) at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with program No. 14699.

*Facilities:* ALMA, HST (WFC3), VLT (FORS2, MUSE, X-SHOOTER).

*Software:* astropy, (Astropy Collaboration et al. 2013), imfit (Erwin 2015), matplotlib (Hunter 2007), Topcat (Taylor 2013), Scipy (Jones et al. 2001), Swarp (Bertin 2010).

## Appendix A
Continuum Detections in the ALMA Data

Besides our targets, our continuum maps reveal a number of detections; see Figure 17. Three objects are detected in the MASOSA pointing with S/N ranging from 5 to 56, while two objects with S/N of 3.5–4.2 are (tentatively) detected in the VR7 pointing; see Table 5 for details. In the MASOSA field, the highest-S/N detection is likely a dusty and highly obscured star-forming galaxy, as it is not detected in optical imaging (limiting AB magnitudes ~28) or in the HST/WFC3 data presented here. The object is detected in the MIR Spitzer/IRAC [3.6] and [4.5] bands, showing that it is very red. No [C II] or any other emission line is detected in the ALMA data, indicating that it is not at similar redshift to MASOSA. The second most luminous continuum detection coincides with the center of a face-on spiral galaxy with a Sloan Digital Sky Survey redshift of $z = 0.131$. The third continuum detection (with an S/N = 5.0) is relatively near MASOSA (at ≈2"/5; see also the right panel of Figure 4) but coincides with a galaxy with photometric redshift $z = 1.46 \pm 0.05$ in the Laigle et al. (2016) catalog. The continuum detections in the VR7 field are of lower significance. The detection with S/N = 4.2 is shown in the left panel of Figure 17 and is colocated with a faint galaxy in the HST data. No emission line is detected in our MUSE data (Matthee et al. 2019, in preparation), suggesting that the galaxy is below $z < 3$ (otherwise Ly$\alpha$ could have been detected) and above $z > 1.5$ (otherwise [O II] could have been detected). The other continuum detection may be associated with an LAE at $z = 3.062$, although it is offset by 0"/25 to the northwest.
<figure><img src="image_21.png" /></figure>

<figure><img src="image_22.png" /><figcaption>**Figure 17.** Continuum maps of foreground sources in the field of VR7 (left) and MASOSA (right). For VR7 we zoom in on the foreground source detected with S/N = 4.2 and show that the IR continuum contours coincide well with the HST/WFC3 position. For MASOSA we show our full primary-beam-corrected image and highlight the positions of the three foreground sources, where contour levels are at the 4σ, 8σ, and 12σ level.</figcaption></figure>

<table><caption>**Table 5**
Global Properties of Foreground Objects</caption><thead><tr><th>R.A.<br>(J2000)</th><th>Decl.<br>(J2000)</th><th>f<sub>v,250 GHz</sub><br>(μJy beam<sup>-1</sup>)</th><th>S/N</th><th>Comment</th></tr></thead><tbody><tr><td>10:01:24.705</td><td>+2:31:32.73</td><td>1212</td><td>56.0</td><td>Dusty, highly obscured star-forming galaxy with unknown photo-z. Detected in the Spitzer/IRAC [3.6] and [4.5] bands, but not detected in optical bands or our HST/WFC3 imaging data.</td></tr><tr><td>10:01:25.427</td><td>+2:31:45.23</td><td>138</td><td>8.2</td><td>In the center of a relatively nearby face-on spiral galaxy at z = 0.131 with identifier SDSS J100125.41 +023145.2.</td></tr><tr><td>10:01:24.629</td><td>+2:31:45.62</td><td>64</td><td>5.0</td><td>Galaxy with photo-z = 1.46 ± 0.05 and M<sub>star</sub> = 10<sup>10.37±0.08</sup> M<sub>☉</sub>, ID 810182 in Laigle et al. (2016) ≈2" away from MASOSA.</td></tr><tr><td>22:18:57.117</td><td>+0:08:03.79</td><td>55</td><td>4.2</td><td>Detected in HST/WFC3 with F110W = 24.4. Marginally detected in a combined griz image. No emission line detected in MUSE (Matthee et al. 2019, in preparation). Likely photo-z ≈ 1.5–2.8.</td></tr><tr><td>22:18:57.143</td><td>+0:08:18.29</td><td>79</td><td>3.5</td><td>Associated (?) with an LAE identified in MUSE z<sub>Lyα</sub> = 3.062, offset 0.25 to the northwest</td></tr></tbody></table>

**Note.** Continuum flux is measured using the primary-beam-corrected map, while the S/N is measured before the primary beam correction.

## Appendix B
### Fitting Rest-frame UV Sizes

Here we provide more details on the fitting of the HST/WFC3 data, in particular on fitting VR7 as a single-component or a two-component model and the fitting results for MASOSA.

#### B.1. Single Component versus Two Component in VR7's HST Data

As described in the main text (Section 3.3), we find that the light profile of VR7 in the HST/WFC3 F160W data is better described by a combination of two exponential profiles than by a single exponential (or Sérsic) profile. The result of the fitting using IMFIT (Erwin 2015) is shown in Figure 18. The residual map of the single-component fit shows clear structure at the ≈3σ level, while residuals of such strength are not seen in the more flexible two-component fit. We note that the F110W data (whose morphology is possibly affected by Lyα emission) are also better described by a two-component fit, although the relative flux and sizes of the two components reverse (i.e., a larger and more

luminous component in the east and a smaller fainter component southwest). While this implies a color gradient, we note that these (and particularly their uncertainties) are challenging to quantify without strong priors such as assuming that the relative sizes, position angles, and ellipticities are the same and assuming a homogeneously distributed Lyα surface brightness. If we use the morphology of the individual components in the F160W data as a prior and assume that Lyα emission is homogeneously distributed, we measure that the flux fraction in component 1 (Figure 11) is 42% ± 6% and 36% ± 7% in the F110W and F160W data, respectively. This implies that component 1 has a UV luminosity and color of $M_{1500} = -21.41 \pm 0.17$ and $\beta = -2.04_{-0.95}^{+0.89}$, respectively, while component 2 has $M_{1500} = -21.77 \pm 0.12$ and $\beta = -1.14 \pm 0.62$. Without subtracting Lyα emission from the F110W flux density, the UV slopes would be bluer, with $\Delta\beta \approx -0.3$.

#### B.2. Single Component in MASOSA

We show our best-fit single exponential model to the F160W from MASOSA in Figure 19. Due to the limited S/N of the
<figure><img src="image_23.png" /><figcaption>**Figure 18.** Zoomed-in HST/WFC3 image of VR7 in the F160W filter. The left panels show the data and the size of the PSF FWHM. The middle panels show the models, and the right panels show the residuals. The top row shows a single exponential light profile, while the bottom row shows a combination of two exponential light profiles. For reference, the dotted contours illustrate the outline of the galaxy. Reduced χ² values are computed within the dotted contour level. The rest-frame UV light profile of VR7 is better described by two individual components.</figcaption></figure>

<figure><img src="image_24.png" /><figcaption>**Figure 19.** Same as Figure 18, but showing the single-component exponential fit to MASOSA.</figcaption></figure>

galaxy, the data cannot reliably distinguish between a single component and a combination of two components, although visually a hint of a second component can be seen in the northwest of the galaxy. Deeper HST data are required to confirm this component.

## Appendix C
### Compilation and Homogenization of Measurements from Galaxies at $z \approx 6-7$

In order to compare our measurements to earlier observations of galaxies at $z \approx 6-7$, we compile a sample of spectroscopically confirmed galaxies at $z \approx 6-7$ for which
<table><caption>**Table 6**

Observed Properties of the Galaxy Compilation at $z \approx 6-7$</caption><thead><tr><th>ID</th><th>z</th><th>M<sub>1500</sub></th><th>β</th><th>f<sub>160 µm</sub> (µJy)</th><th>L<sub>[C II]</sub>/10<sup>8</sup> (L<sub>☉</sub>)</th><th>v<sub>FWHM,[C II]</sub> (km s<sup>-1</sup>)</th><th>References</th></tr></thead><tbody><tr><td>A383-5.1</td><td>6.029</td><td>-19 ± 0.3</td><td>-2 ± 0.7</td><td>&lt;11</td><td>0.089 ± 0.031</td><td>100 ± 23</td><td>K16</td></tr><tr><td>NTTDF2313</td><td>6.07</td><td>-21.2</td><td>...</td><td>&lt;18</td><td>&lt;0.21</td><td>...</td><td>C18</td></tr><tr><td>WMH5</td><td>6.075</td><td>-22.8</td><td>...</td><td>218 ± 41</td><td>6.6 ± 0.7</td><td>251 ± 15</td><td>W15</td></tr><tr><td>BDF2203</td><td>6.12</td><td>-21.5</td><td>...</td><td>&lt;23</td><td>1.3 ± 0.3</td><td>150 ± 50</td><td>C18</td></tr><tr><td>CLM1</td><td>6.176</td><td>-22.7</td><td>...</td><td>44 ± 26</td><td>2.4 ± 0.3</td><td>162 ± 23</td><td>W15</td></tr><tr><td>GOODS3203</td><td>6.27</td><td>-21.6</td><td>...</td><td>&lt;41</td><td>&lt;0.59</td><td>...</td><td>C18</td></tr><tr><td>COSMOS20521</td><td>6.36</td><td>-21.13</td><td>-2.17</td><td>&lt;20</td><td>&lt;0.23</td><td>...</td><td>C18</td></tr><tr><td>VR7</td><td>6.534</td><td>-22.4 ± 0.05</td><td>-1.4 ± 0.3</td><td>&lt;10.6</td><td>4.76 ± 0.14</td><td>203 ± 40</td><td>This paper</td></tr><tr><td>MASOSA</td><td>6.543</td><td>-20.9 ± 0.1</td><td>-1.1 ± 0.7</td><td>&lt;9.2</td><td>&lt;0.22</td><td>...</td><td>This paper</td></tr><tr><td>UDS4821</td><td>6.56</td><td>-21.3</td><td>...</td><td>&lt;24</td><td>&lt;0.32</td><td>...</td><td>C18</td></tr><tr><td>Himiko</td><td>6.595</td><td>-21.8 ± 0.1</td><td>-2 ± 0.4</td><td>&lt;27</td><td>1.2 ± 0.2</td><td>180 ± 50</td><td>C18</td></tr><tr><td>CR7</td><td>6.604</td><td>-22.2 ± 0.1</td><td>-2.2 ± 0.4</td><td>&lt;7</td><td>3.87 ± 0.65</td><td>259 ± 24</td><td>M17</td></tr><tr><td>COSMOS24108</td><td>6.629</td><td>-21.67</td><td>-1.76 ± 0.2</td><td>&lt;18</td><td>1.1 ± 0.15</td><td>150 ± 40</td><td>P16</td></tr><tr><td>UDS16291</td><td>6.638</td><td>-20.97</td><td>-2.4 ± 0.4</td><td>&lt;20</td><td>0.7 ± 0.15</td><td>50 ± 15</td><td>P16</td></tr><tr><td>NTTDF6345</td><td>6.701</td><td>-21.57</td><td>-1.3 ± 0.5</td><td>&lt;16</td><td>1.8 ± 0.3</td><td>250 ± 70</td><td>P16</td></tr><tr><td>RXJ13471216</td><td>6.765</td><td>-19.1 ± 0.4</td><td>-2.5 ± 0.7</td><td>&lt;15</td><td>0.15 ± 0.03</td><td>75 ± 25</td><td>B17</td></tr><tr><td>COS-2987030247</td><td>6.807</td><td>-22.1 ± 0.1</td><td>-1.22 ± 0.5</td><td>&lt;25</td><td>3.6 ± 0.5</td><td>124 ± 18</td><td>L17</td></tr><tr><td>SDF-46975</td><td>6.844</td><td>-21.5</td><td>...</td><td>&lt;19.2</td><td>&lt;0.27</td><td>...</td><td>M15</td></tr><tr><td>COS-3018555981</td><td>6.854</td><td>-22 ± 0.1</td><td>-1.18 ± 0.5</td><td>&lt;29</td><td>4.7 ± 0.5</td><td>232 ± 30</td><td>S17</td></tr><tr><td>IOK-1</td><td>6.96</td><td>-21.4</td><td>-2.07 ± 0.26</td><td>&lt;21</td><td>&lt;0.25</td><td>...</td><td>O14</td></tr><tr><td>BDF-521</td><td>7.008</td><td>-20.3</td><td>...</td><td>&lt;17.4</td><td>&lt;0.22</td><td>...</td><td>M15</td></tr><tr><td>BDF-3299</td><td>7.109</td><td>-20.9</td><td>...</td><td>&lt;7.8</td><td>0.49 ± 0.1</td><td>75 ± 20</td><td>M15,C17</td></tr><tr><td>COSMOS13679</td><td>7.145</td><td>-21.46</td><td>-1.54 ± 0.4</td><td>&lt;14</td><td>0.7 ± 0.15</td><td>90 ± 35</td><td>P16</td></tr><tr><td>B14-65666</td><td>7.152</td><td>-22.3 ± 0.05</td><td>-1.85 ± 0.5</td><td>130 ± 25</td><td>13.3 ± 1.3</td><td>349 ± 31</td><td>B18,H18</td></tr><tr><td>SXDF-NB1006-2</td><td>7.212</td><td>-21.5</td><td>-2.6</td><td>&lt;14</td><td>&lt;0.28</td><td>...</td><td>I16</td></tr></tbody></table>

**Note.** Upper limits are at the $1\sigma$ level and in the case of [C II] based on the expected line width from the UV Luminosity (Figure 9).

**References.** B17: Bradač et al. 2017; B18: Bowler et al. 2018; C17: Carniani et al. 2017; C18: Carniani et al. 2018b; H18: Hashimoto et al. 2019; I16: Inoue et al. 2016; K16: Knudsen et al. 2016; L17: Laporte et al. 2017; M15: Maiolino et al. 2015; M17: Matthee et al. 2017b; O14: Ota et al. 2014; P16: Pentericci et al. 2016; S17: Smit et al. 2018; W15: Willott et al. 2015.

<table><caption>**Table 7**

Derived Properties of the Galaxy Compilation at $z \approx 6-7$, Based on the Measurements Listed in Table 6</caption><thead><tr><th>ID</th><th>SFR<sub>UV</sub>(M<sub>☉</sub> yr<sup>-1</sup>)</th><th>L<sub>IR,Td=35 K</sub>/10<sup>10</sup> L<sub>☉</sub></th><th>L<sub>IR,Td=45 K</sub>/10<sup>10</sup> L<sub>☉</sub></th><th>SFR<sub>IR,Td=45 K</sub>(M<sub>☉</sub> yr<sup>-1</sup>)</th></tr></thead><tbody><tr><td>A383-5.1</td><td>3.5 ± 0.9</td><td>&lt;1.2</td><td>&lt;2.4</td><td>&lt;4.2</td></tr><tr><td>NTTDF2313</td><td>19.0</td><td>&lt;2.0</td><td>&lt;4.0</td><td>&lt;6.9</td></tr><tr><td>WMH5</td><td>63.6</td><td>24.0 ± 4.5</td><td>48.2 ± 8.9</td><td>83.9 ± 15.6</td></tr><tr><td>BDF2203</td><td>23.8</td><td>&lt;2.6</td><td>&lt;5.1</td><td>&lt;9.0</td></tr><tr><td>CLM1</td><td>59.0</td><td>5.0 ± 2.7</td><td>10.0 ± 5.6</td><td>17.4 ± 9.8</td></tr><tr><td>GOODS3203</td><td>25.7</td><td>&lt;4.8</td><td>&lt;9.6</td><td>&lt;16.6</td></tr><tr><td>COSMOS20521</td><td>18.0</td><td>&lt;2.4</td><td>&lt;4.8</td><td>&lt;8.3</td></tr><tr><td>VR7</td><td>47.1 ± 1.7</td><td>&lt;1.3</td><td>&lt;2.6</td><td>&lt;4.6</td></tr><tr><td>MASOSA</td><td>15.0 ± 2.0</td><td>&lt;1.2</td><td>&lt;2.3</td><td>&lt;4.0</td></tr><tr><td>UDS4821</td><td>20.5</td><td>&lt;3.1</td><td>&lt;6.1</td><td>&lt;10.5</td></tr><tr><td>Himiko</td><td>29.9 ± 2.2</td><td>&lt;3.5</td><td>&lt;6.9</td><td>&lt;12.0</td></tr><tr><td>CR7</td><td>40.7 ± 3.2</td><td>&lt;0.9</td><td>&lt;1.8</td><td>&lt;3.1</td></tr><tr><td>COSMOS24108</td><td>27.1</td><td>&lt;2.3</td><td>&lt;4.6</td><td>&lt;8.1</td></tr><tr><td>UDS16291</td><td>16.0</td><td>&lt;2.6</td><td>&lt;5.2</td><td>&lt;9.0</td></tr><tr><td>NTTDF6345</td><td>25.1</td><td>&lt;2.1</td><td>&lt;4.2</td><td>&lt;7.3</td></tr><tr><td>RXJ13471216</td><td>4.0 ± 1.2</td><td>&lt;2.0</td><td>&lt;4.0</td><td>&lt;7.0</td></tr><tr><td>COS-2987030247</td><td>37.4 ± 3.0</td><td>&lt;3.4</td><td>&lt;6.7</td><td>&lt;11.7</td></tr><tr><td>SDF-46975</td><td>23.8</td><td>&lt;2.7</td><td>&lt;5.2</td><td>&lt;9.1</td></tr><tr><td>COS-3018555981</td><td>34.7 ± 2.6</td><td>&lt;4.0</td><td>&lt;7.9</td><td>&lt;13.8</td></tr><tr><td>IOK-1</td><td>22.1</td><td>&lt;3.0</td><td>&lt;5.9</td><td>&lt;10.3</td></tr><tr><td>BDF-521</td><td>9.6</td><td>&lt;2.5</td><td>&lt;4.9</td><td>&lt;8.6</td></tr><tr><td>BDF-3299</td><td>15.1</td><td>&lt;1.2</td><td>&lt;2.3</td><td>&lt;4.0</td></tr><tr><td>COSMOS13679</td><td>23.1</td><td>&lt;2.1</td><td>&lt;4.1</td><td>&lt;7.2</td></tr><tr><td>B14-65666</td><td>43.3 ± 1.6</td><td>19.8 ± 3.9</td><td>38.3 ± 7.6</td><td>66.7 ± 13.3</td></tr><tr><td>SXDF-NB1006-2</td><td>23.8</td><td>&lt;2.2</td><td>&lt;4.2</td><td>&lt;7.3</td></tr></tbody></table>

**Note.** SFRs are based on the calibrations presented in Kennicutt (1998). Infrared luminosities have been corrected for CMB heating following da Cunha et al. (2013). Upper limits are at the $1\sigma$ level.
<table><caption>**Table 8**

Lyα Properties of the Galaxy Compilation at $z \approx 6-7$</caption><thead><tr><th>ID</th><th>L<sub>Lyα</sub>/10<sup>42</sup> erg s<sup>-1</sup></th><th>EW<sub>0,Lyα</sub> (Å)</th><th>Δv<sub>Lyα</sub> (km s<sup>-1</sup>)</th></tr></thead><tbody><tr><td>A383-5.1</td><td>7</td><td>138</td><td>90 ± 20</td></tr><tr><td>NTTDF2313</td><td>...</td><td>...</td><td>...</td></tr><tr><td>WMH5</td><td>...</td><td>13</td><td>265 ± 52</td></tr><tr><td>BDF2203</td><td>...</td><td>3</td><td>...</td></tr><tr><td>CLM1</td><td>...</td><td>50</td><td>430 ± 69</td></tr><tr><td>GOODS3203</td><td>...</td><td>5</td><td>...</td></tr><tr><td>COSMOS20521</td><td>3</td><td>10</td><td>...</td></tr><tr><td>VR7</td><td>24</td><td>34</td><td>217 ± 25</td></tr><tr><td>MASOSA</td><td>24</td><td>145</td><td>...</td></tr><tr><td>UDS4821</td><td>...</td><td>48</td><td>...</td></tr><tr><td>Himiko</td><td>43</td><td>65</td><td>145 ± 15</td></tr><tr><td>CR7</td><td>83</td><td>211</td><td>167 ± 27</td></tr><tr><td>COSMOS24108</td><td>10</td><td>27</td><td>240 ± 0</td></tr><tr><td>UDS16291</td><td>0.8</td><td>6</td><td>110 ± 0</td></tr><tr><td>NTTDF6345</td><td>4</td><td>15</td><td>110 ± 0</td></tr><tr><td>RXJ13471216</td><td>6</td><td>26</td><td>20 ± 90</td></tr><tr><td>COS-2987030247</td><td>5</td><td>16</td><td>326 ± 30</td></tr><tr><td>SDF-46975</td><td>15</td><td>43</td><td>...</td></tr><tr><td>COS-3018555981</td><td>...</td><td>...</td><td>...</td></tr><tr><td>IOK-1</td><td>...</td><td>42</td><td>...</td></tr><tr><td>BDF-521</td><td>...</td><td>64</td><td>...</td></tr><tr><td>BDF-3299</td><td>7</td><td>50</td><td>64 ± 10</td></tr><tr><td>COSMOS13679</td><td>13</td><td>28</td><td>135 ± 0</td></tr><tr><td>B14-65666</td><td>2.6</td><td>4</td><td>772 ± 45</td></tr><tr><td>SXDF-NB1006-2</td><td>...</td><td>33</td><td>...</td></tr></tbody></table>

the [C II] line and adjacent continuum have been observed with ALMA. This compilation includes LAEs (Ota et al. 2014; Inoue et al. 2016; Matthee et al. 2017b; Carniani et al. 2018a), UV-selected galaxies confirmed through Lyα emission (Maiolino et al. 2015; Willott et al. 2015; Pentericci et al. 2016; Carniani et al. 2018b; Hashimoto et al. 2019), lensed UV-selected galaxies (Knudsen et al. 2016; Bradač et al. 2017), and UV-selected galaxies confirmed through [C II] emission (Smit et al. 2018). We also include the sample of LBGs, LAEs, and galaxies selected through their photometric redshift at $z \approx 5$ from Capak et al. (2015) with updated measurements from Barisic et al. (2017). For a proper comparison, we collect observed properties (i.e., UV, [C II], and Lyα luminosity, line widths, and IR flux densities) and self-consistently convert these to derived properties (i.e., SFRs and IR luminosities). Table 6 shows the observed properties of the galaxies at $z \approx 6-7$, while Table 7 lists derived properties. The Lyα properties are listed in Table 8.

We use published values of $M_{1500}$ and $\beta$ wherever possible, but for several objects we calculated $M_{1500}$ based on published photometry. When available, Lyα luminosities have been computed using published line flux measurements. For galaxies in which no [C II] is detected, we measure the physically motivated upper limit similar to the procedure in Section 5.2. This means that we estimate the expected line width based on the UV luminosity (typically 150–200 km s$^{-1}$) and convert published noise values to this width assuming rms$_{[C II]} \propto \sqrt{\Delta v_{channel}}$, where $\Delta v_{channel}$ is the channel width over which the noise is measured. Measurements of upper limits on the continuum flux at $\lambda_0 = 160 \mu m$ (typically $\lambda_{obs} = 1.2$ mm) have been converted to an FIR ($\lambda = 42.5-122.5 \mu m$) and IR ($\lambda = 8-1000 \mu m$) luminosity following Ota et al. (2014), after correcting the contribution from CMB heating following da Cunha et al. (2013). We

assume a modified blackbody of the form $B_{\nu} = \frac{M_{\text{dust}} \kappa}{D_L^2 c^2 e^{k_B T_d}}$, where $D_L$ is the luminosity distance; $k_B$ is the Boltzmann constant; $T_d$ is the dust temperature; $M_{\text{dust}}$ is the dust mass; $\kappa = \kappa_{850} (\frac{\nu}{\nu_{850}})^{\beta}$, where $\kappa_{850}$ is the dust opacity at 850 $\mu m$, which we assume is 0.077 $m^2 kg^{-1}$, similar to Ota et al. (2014); and $\beta$ is the graybody power-law exponent, which we assume to be $\beta = 1.5$ following Ota et al. (2014). Then, the obscured SFR is calculated following Kennicutt (1998): $\log_{10}(SFR_{IR}/M_{\odot} \text{yr}^{-1}) = \log_{10}(L_{IR}/\text{erg s}^{-1}) - 43.48$. $M_{1500}$ is converted to an unobscured SFR$_{UV}$ following Kennicutt (1998): $\log_{10}(SFR_{UV}/M_{\odot} \text{yr}^{-1}) = -0.328 (M_{1500}+17.3)$. The SFRs can be converted to the calibrations in Kennicutt & Evans (2012) that follow a Kroupa IMF by multiplying SFR$_{UV}$ by a factor of 0.63 and multiplying SFR$_{IR}$ by a factor of 0.86.

## ORCID iDs

J. Matthee https://orcid.org/0000-0003-2871-127X

D. Sobral https://orcid.org/0000-0001-8823-4845

L. A. Boogaard https://orcid.org/0000-0002-3952-8588

L. Vallini https://orcid.org/0000-0002-3258-3672

A. Ferrara https://orcid.org/0000-0002-9400-7312

D. Schaerer https://orcid.org/0000-0001-7144-7182

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