# The Source of Leaking Ionizing Photons from Haro11: Clues from HST/COS Spectroscopy of Knots A, B, and C

Göran Östlin, T. Emil Rivera-Thorsen, Veronica Menacho, Matthew Hayes, Axel Runnholm, Genoveva Micheva, M. S. Oey, Angela Adamo, Arjan Bik, John M. Cannon, et al.

► To cite this version:

Göran Östlin, T. Emil Rivera-Thorsen, Veronica Menacho, Matthew Hayes, Axel Runnholm, et al. The Source of Leaking Ionizing Photons from Haro11: Clues from HST/COS Spectroscopy of Knots A, B, and C. The Astrophysical Journal, 2021, 912, 10.3847/1538-4357/abf1e8. insu-03748210

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# The Source of Leaking Ionizing Photons from Haro 11: Clues from HST/COS Spectroscopy of Knots A, B, and C*

Göran Östlin<fnref n="1" />, T. Emil Rivera-Thorsen<fnref n="1" />, Veronica Menacho<fnref n="1" />, Matthew Hayes<fnref n="1" />, Axel Runnholm<fnref n="1" />, Genoveva Micheva<fnref n="2" />, M. S. Oey<fnref n="3" />, Angela Adamo<fnref n="1" />, Arjan Bik<fnref n="1" />, John M. Cannon<fnref n="4" />, Max Gronke<fnref n="5" />,<fnref n="12" />, Daniel Kunth<fnref n="6" />, Peter Laursen<fnref n="7" />,<fnref n="8" />, J. Miguel Mas-Hesse<fnref n="9" />, Jens Melinder<fnref n="1" />, Matteo Messa<fnref n="10" />, Mattia Sirressi<fnref n="1" />, and Linda Smith<fnref n="11" />

Received 2020 September 11; revised 2021 March 19; accepted 2021 March 23; published 2021 May 17

## Abstract

Understanding the escape of ionizing (Lyman continuum) photons from galaxies is vital for determining how galaxies contributed to reionization in the early universe. While directly detecting the Lyman continuum from high-redshift galaxies is impossible due to the intergalactic medium, low-redshift galaxies in principle offer this possibility but require observations from space. The first local galaxy for which Lyman continuum escape was found is Haro 11, a luminous blue compact galaxy at $z=0.02$, where observations with the FUSE satellite revealed an escape fraction of 3.3%. However, the FUSE aperture covers the entire galaxy, and it is not clear from where the Lyman continuum is leaking out. Here we utilize Hubble Space Telescope/Cosmic Origins Spectrograph spectroscopy in the wavelength range 1100–1700 Å of the three knots (A, B, and C) of Haro 11 to study the presence of Ly$\alpha$ emission and the properties of intervening gas. We find that all knots have bright Ly$\alpha$ emission. UV absorption lines, originating in the neutral interstellar medium, as well as lines probing the ionized medium, are seen extending to blueshifted velocities of 500 km s$^{-1}$ in all three knots, demonstrating the presence of an outflowing multiphase medium. We find that knots A and B have large covering fractions of neutral gas, making LyC escape along these sightlines improbable, while knot C has a much lower covering fraction (≤50%). Knot C also has the the highest Ly$\alpha$ escape fraction, and we conclude that it is the most likely source of the escaping Lyman continuum detected in Haro 11.

*Unified Astronomy Thesaurus concepts: Blue compact dwarf galaxies (165); Starburst galaxies (1570); Photoionization (2060); Lyman-alpha galaxies (978); Interstellar line absorption (843)*

## 1. Introduction

Understanding how galaxies have contributed to the reionization of the universe is a major theme in contemporary astrophysics. In the quest for this, it is vital to know under what conditions ionizing (Lyman continuum, LyC) radiation can make their way through the interstellar medium (ISM) and escape from galaxies. Directly observing LyC escape in the epoch of reionization (EoR; $z>6$) and even at somewhat lower redshifts ($z>4$) is impossible due to the long path lengths and increasing probability of intersecting neutral clouds in the intergalactic medium (IGM; Inoue & Iwata 2008). Photometric surveys have been used to find candidate LyC leakers at $z \sim 3$ (e.g., Iwata et al. 2009; Micheva et al. 2017; Mostardi et al. 2013), but such surveys are prone to contamination by lower-redshift sources and follow-up studies are typically able to confirm only very few candidates selected by photometry as real LyC leakers (Vanzella et al. 2012; Siana et al. 2015; Mostardi et al. 2015). At lower redshifts ($z \lesssim 2$), IGM absorption is less of an obstacle, but with the exception of

rare objects strongly magnified by gravitational lensing (e.g., the Sunburst arc at $z=2.4$; Rivera-Thorsen et al. 2019), sources are still largely unresolved, preventing a detailed study of the escape mechanism. Galaxies in the local universe, on the other hand, offer the possibility of detailed studies of processes that could lead to LyC escape, but here the challenge is to actually measure the escape of photons with a rest-frame wavelength in the far-ultraviolet, $\lambda \sim 900$ Å.

Sensitivity to such wavelengths was provided with the FUSE satellite, and the first local galaxy for which a net escape of LyC radiation was claimed was Haro 11 (Bergvall et al. 2006). This claim was disputed by Grimes et al. (2007), but a thorough reanalysis confirmed the leakage (Leitet et al. 2011), and an escape fraction of 3.3% was derived. However, the large (30" × 30") aperture of the FUSE spectrograph covers the entire galaxy, leaving the question “from where in the galaxy is LyC leaking out?” unanswered. The Hubble Space Telescope/Cosmic Origins Spectrograph (HST/COS; Green et al. 2012) is also capable of measuring spectra down to the rest-frame Lyman limit (McCandless et al. 2010) but has very poor sensitivity below 1100 Å, making observations quite costly in terms of exposure time, but some leaking galaxies have

<footnote n="1">¹Department of Astronomy, Oskar Klein Centre; Stockholm University; SE-106 91 Stockholm, Sweden; ostlin@astro.su.se</footnote>

<footnote n="2">²Leibniz-Institute for Astrophysics Potsdam, An der Sternwarte 16, D-14482 Potsdam, Germany</footnote>

<footnote n="3">³University of Michigan, Department of Astronomy, 323 West Hall, 1085 South University Avenue, Ann Arbor, MI 48109-1107, USA</footnote>

<footnote n="4">⁴Department of Physics & Astronomy, Macalester College, 1600 Grand Avenue, Saint Paul, MN 55105, USA</footnote>

<footnote n="5">⁵Department of Physics & Astronomy, Johns Hopkins University, Bloomberg Center, 3400 N. Charles Street, Baltimore, MD 21218, USA</footnote>

<footnote n="6">⁶Institut d'Astrophysique de Paris, 98 bis Boulevard Arago, 75014 Paris, France</footnote>

<footnote n="7">⁷Cosmic Dawn Center (DAWN), Denmark</footnote>

<footnote n="8">⁸Niels Bohr Institute, University of Copenhagen, Jagtvej 128, DK-2200, Copenhagen N, Denmark</footnote>

<footnote n="9">⁹Centro de Astrobiología (CSIC/INTA), Departamento de Astrofísica, P.O. Box 78, E-28692 Villanueva de la Cañada, Madrid, Spain</footnote>

<footnote n="10">¹⁰University of Massachusetts, Department of Astronomy, 710 North Pleasant Street, Amherst, MA 01003-9305, USA</footnote>

<footnote n="11">¹¹European Space Agency (ESA), ESA Office, Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA</footnote>

<footnote>* Based of observations with the NASA/ESA Hubble Space Telescope.</footnote>

<footnote n="12">¹² Hubble fellow.</footnote>
<figure><img src="image_3.png" /><figcaption>**Figure 1.** RGB composite of Haro 11 with Lyα in blue, UV continuum (λ ~ 1500 Å) in green, and Hα in red (Östlin et al. 2009; Rivera-Thorsen et al. 2017). The white circle shows the COS aperture for the previously published spectrum of knot C (Alexandrov et al. 2015; Rivera-Thorsen et al. 2017), and the orange circles the COS apertures for the new spectra presented in this paper.</figcaption></figure>

nevertheless been confirmed (Leitherer et al. 2016; Puschnig et al. 2017). Subsequently, HST/COS has been successfully used to measure LyC escape from compact galaxies at $z \gtrsim 0.2$ (Izotov et al. 2018a, 2018b) where the COS sensitivity is much higher. These sources, mostly of the green pea galaxy type (Cardamone et al. 2009) selected to have high ionization (emission-line ratios $[O III]/[O II] > 5$) are, however, barely resolved, preventing an analysis of the details of the escape mechanism.

Haro 11 is a luminous blue compact galaxy (BCG) at $z=0.0206$ with a starburst triggered by a dwarf galaxy merger (Östlin et al. 2001, 2015). It features an irregular morphology with three bright knots: A, B, and C (see Figure 1 and Vader et al. 1993). It has a low metallicity (oxygen abundance ~20% solar, Bergvall & Östlin 2002) but with indications that the oxygen abundance differs significantly between the knots by more than 0.2 dex (Guseva et al. 2012; James et al. 2013; Menacho et al. 2021). Its star formation rate is in the range $20–30 \ M_{\odot} \ yr^{-1}$ depending on the tracer used (Östlin et al. 2015). Haro 11 is remarkably devoid of neutral hydrogen, with only $5 \times 10^{8} M_{\odot}$ (MacHattie et al. 2014; Pardy et al. 2016) detected, and the molecular mass (albeit quite uncertain) is of the order $\gtrsim 10^{9} M_{\odot}$ (Cormier et al. 2014). The gas consumption timescale is hence on the order of 50 Myr, making Haro 11 one of the most extreme starbursts known. Bright X-ray point sources are associated with knots B and C (Grimes et al. 2007; Prestwich et al. 2015; Basu-Zych et al. 2016; Dittenber et al. 2020). Lyα imaging of Haro 11 showed bright emission from knot C, with extensions to the north and south indicating a possible outflow origin and large-scale diffuse emission on a galaxy-wide scale, but nothing from knots A and B (Hayes et al. 2007; Östlin et al. 2009). Because the escape of Lyα, just like that of LyC, is sensitive to the amount of neutral hydrogen, Hayes et al. (2007) suggested that the escaping LyC radiation was likely to come from knot C. Being the nearest known LyC emitter, Haro 11 offers unique opportunities to study the details

of LyC escape and how indirect LyC diagnostics behave in a spatially resolved fashion.

Knot C was observed with HST/COS in gratings G130M and G160M (Alexandrov et al. 2015; Heckman et al. 2015) showing evidence of outflows of neutral gas probed by low ionization state (LIS) metal lines, e.g., from Si II. These spectra were reanalyzed by Rivera-Thorsen et al. (2017) using the apparent optical depth (AOD) method (Savage & Sembach 1991; Jones et al. 2013; Rivera-Thorsen et al. 2015) applied to Si II lines of different oscillator strengths to solve for both the column density and covering fraction of the absorbing material, as a function of velocity. Because silicon has a lower ionization energy (8.15 eV) than hydrogen (16.6 eV), Si II traces gas where hydrogen is mostly neutral, while Si IV traces highly ionized gas (comparable to $[O III]$). These results showed that knot C has a low covering (~50%) of neutral gas clouds over a wide velocity range along the sight line, corroborating it as a good candidate for the source of ionizing photons observed with FUSE.

In the classical picture, a H II region that is radiation bounded will exhibit a highly ionized inner $[O III]$ zone and an outer $[O II]$ zone (where rarer photons capable of doubly ionizing oxygen have been exhausted; Osterbrock & Ferland 2006). If, on the other hand, there are enough photons to completely ionize the local ISM, the H II region is said to be density bounded, and ionizing photons will leak out into the galactic-scale ISM. In this case, the ionization state of the gas, as traced by $[O III]/[O II]$, will remain high (Pellegrini et al. 2012; Jaskot & Oey 2013). This is a factor for the successful COS observations at $z \gtrsim 0.2$ of LyC escape from galaxies selected to have high global $[O III]/[O II]$ ratios from SDSS (Izotov et al. 2018a, 2018b). Haro 11 was imaged by the HST in the $[O II]$ and $[O III]$ lines, with the scope to use this diagnostic to pinpoint likely sites of escaping LyC radiation. With knot C showing low ionization, knot B moderate, and knot A large-scale high ionization, Keenan et al. (2017) proposed knot A as the most probable source of leaking LyC.

Here we present new HST/COS observations with gratings G130M and G160M of knots A and B. We characterize the Lyα emission and the Si II and Si IV transitions, along with knot C, and include optical emission-line diagnostics from ESO/VLT/MUSE to provide further insight into the question on which source(s) that may dominate the observed LyC escape

Throughout the paper, we assume a luminosity distance of 88 Mpc ($m - M = 34.7$) and a scale of $0.42 \ kpc \ arcsec^{-1}$.

## 2. Observations and Reductions

### 2.1. COS Data

We have used archival HST/COS observations in G130M and G160M of knot C from program 13017 (PI Heckman), and new COS observations of knots A and B with the same gratings (from program 15352, PI Oestlin). Knot C was used for acquisition, using “mirror A” (see COS instrument handbook, IHB, Hirschauer et al. 2021), followed by blind offsets to knots A and B. The reason for this was that knot C is a bright and relatively isolated point-like source in the UV, while knots A and B have a more complicated structure encompassing several star clusters. Hence, acquiring on knot C, followed by blind offsets, was judged to produce a more predictable location of the apertures on knots A and B. Each of knots A and B was
<figure><img src="image_4.png" /><figcaption>**Figure 2.** COS/NUV acquisition images for (from left to right) knots A, B, and C, with logarithmic intensity scaling. The position of the nominal COS aperture ($\ominus = 2''/5$) is shown as a green circle. The orientations on sky are indicated by the green compasses in the upper-right corners. The dispersion direction for the FUV spectra is from left to right. The images of knots A and B were obtained with the primary science aperture (PSA) and “mirror A,” while knot C (from GO 13017) was obtained with PSA and “mirror B,” and shows the characteristic (for mirror B) secondary ghost image indicated by the red arrow. Note that this is just a characteristic of “mirror B” and does not produce a secondary FUV spectrum. Moreover, it should not be confused with the secondary source (visible as a small green dot, $\sim 1''$ W of knot C) in aperture C in Figure 1, which is a bona fide star cluster. For knot A (left), the yellow circle indicates a source discussed in Section 4.</figcaption></figure>

observed in single visits of two-orbit durations. Acquisition images of 51 and 60 s were obtained for knots A and B, respectively. The new observations use COS lifetime position 4 (see COS IHB, Hirschauer et al. 2021). Four different wavelength settings (denoted CENWAVE and FP-POS in the nomenclature of the COS IHB) per observation were used to cover detector gaps, optimized to get the best coverage of the features of interest. In order to protect the COS detector (e.g., from geocoronal Ly$\alpha$), not all possible combinations were allowed. For knot A, we used for G130M the wavelength setting CENWAVE=1222 with FP-POS 1 and 2, and setting 1291 with FP-POS 3 and 4, for a total exposure time of 1929 s. For knot B, setting 1222 with FP-POS 2 and 3, and 1291 with FP-POS 3 and 4, was used for a total exposure time of 1869 s. For G160M, both knots were observed with settings 1577 (FP-POS 2 and 3) and 1611 (FP-POS 3 and 4), for a total exposure time of 2558 s per knot. All observations were obtained through the primary science aperture (PSA) using “mirror A.” See Figure 1 for the location of the COS apertures on Haro 11.

The spectral resolution of COS is degraded for extended sources. To mitigate this spectral degradation, we restricted the spacecraft ORIENT range for knots A and B such that each of the sources would have its minimal spatial extent in the dispersion direction. For knot C, which consists of a single point-like source, this is not an issue, and the (archival) observations were obtained at a PA\_APER (the position angle on sky, measured from N to E, of the cross-dispersion axis, i.e., y-coordinate in the detector frame) of 258°6. To get the position angle of the dispersion direction in the sky frame, 90° is added to yield 348°6. In order to determine suitable ORIENT ranges for knots A and B, we used available ACS/SBC/F140LP far-UV images from program 9470 (PI Kunth) and ACS/HRC/F220W near-UV images from program 10575 (PI Östlin). For knot A, the COS observations were obtained at PA\_APER = -90.0 so the dispersion direction is 0°, i.e., directed to N. For knot B, the observations were done at PA\_APER = 94.2, so the dispersion direction is 184°, approximately south; see Figure 2.

Individual observations were reduced and extracted with CALCOS version 3.3.4; extracted spectra were combined using

custom splicing scripts. Spectra were first rebinned by a factor of 6 because of the oversampling of the spectrograph and to Nyquist-sample the nominal line-spread function.

## 2.2. MUSE Data

ESO/VLT/MUSE: Haro 11 has been observed with the MUSE integral field spectrometer (Bacon et al. 2010) on ESO/VLT in the extended setting covering a wavelength of $\sim 4650-9500$ Å. The reduction of these data is described in Menacho et al. (2019), and we here summarize the most important steps.

The sky subtraction was performed in two steps. In the first step, a sky mask was created by the MUSE pipeline “muse_scipost” from the 20% darkest area of each science exposure. These sky regions, however, also enclose areas in the halo with real but faint $H\alpha$ and [O III]5007 emission and as consequence produced an oversubtraction of the sky. Thus, in a second step, we customized the sky mask by removing these emission areas. The sky residuals are $<2 \times 10^{-20}$ erg s$^{-1}$ cm$^{-2}$ per pixel, or a maximum $10^{-18}$ erg s$^{-1}$ cm$^{-2}$ over a COS aperture, while $H\alpha$ and $H\beta$ line fluxes are four to five orders of magnitude greater. Hence, uncertainties in the sky level do not incur significant uncertainties in the $H\alpha/H\beta$ values or the derived extinctions. The $H\alpha$ and $H\beta$ spectra were corrected for underlying stellar absorption with the pPXF package (Cappellari 2017) using the MILES stellar spectra library (Sánchez-Blázquez et al. 2006; Vazdekis et al. 2010). We use these data to measure optical emission-line fluxes in the three knots.

# 3. Results

Here we describe the results and measurements done for each of knots A, B, and C.

## 3.1. MUSE

We used the COS acquisition images to find the exact location of the spectroscopic apertures in the MUSE image plane. Because the throughput of the $\ominus = 2''/5$ COS aperture is not a top-hat function, but subject to vignetting and also transmits some flux from $1''/25 < r < 2''$, we multiplied the
<table><caption>**Table 1**
Measurements for Knots A, B, and C</caption><thead><tr><th>Quantity</th><th>Knot A</th><th>Knot B</th><th>Knot C</th><th>Units/Comment</th></tr></thead><tbody><tr><td>F<sub>Hα</sub></td><td>2.27</td><td>5.19</td><td>1.40</td><td>(10<sup>-13</sup> erg s<sup>-1</sup> cm<sup>-2</sup>)</td></tr><tr><td>ν<sub>Hα</sub></td><td>+57</td><td>-24</td><td>-54</td><td>(km s<sup>-1</sup>)</td></tr><tr><td>ν<sub>star</sub></td><td>+51</td><td>-24</td><td>-21</td><td>(km s<sup>-1</sup>)</td></tr><tr><td>σ<sub>Hα</sub></td><td>71.1</td><td>89.2</td><td>64.2</td><td>(km s<sup>-1</sup>)</td></tr><tr><td>EW(Hα)</td><td>581</td><td>827</td><td>153</td><td>(Å)</td></tr><tr><td>[O III]/Hβ</td><td>4.47</td><td>4.06</td><td>3.19</td><td>5007/Hβ</td></tr><tr><td>[S III]/[S II]</td><td>3.18</td><td>3.63</td><td>1.55</td><td>(9531+9069)/(6717+6730)</td></tr><tr><td>(Hα/Hβ)<sub>raw</sub></td><td>3.52</td><td>4.38</td><td>4.24</td><td>observed</td></tr><tr><td>(Hα/Hβ)<sub>corr</sub></td><td>3.48</td><td>4.33</td><td>4.19</td><td>corrected for MW ext</td></tr><tr><td>E(B-V)<sub>n</sub></td><td>0.18</td><td>0.38</td><td>0.35</td><td>from (Hα/Hβ)<sub>corr</sub></td></tr><tr><td>β</td><td>-1.73</td><td>-0.86</td><td>-1.89</td><td>UV continuum slope</td></tr><tr><td>F<sub>Lyα</sub></td><td>0.349</td><td>0.656</td><td>1.515</td><td>(10<sup>-13</sup> erg s<sup>-1</sup> cm<sup>2</sup>)</td></tr><tr><td>W<sup>red</sup><sub>Lyα</sub></td><td>302</td><td>338</td><td>195</td><td>(km s<sup>-1</sup>) FWHM red peak</td></tr><tr><td>Δν<sub>Lyα</sub></td><td>532</td><td>409</td><td>(400<sup>a</sup>)</td><td>(km s<sup>-1</sup>) peak separation</td></tr><tr><td>A<sup>red</sup></td><td>1.5</td><td>2</td><td>2.7</td><td>Lyα red peak asymm</td></tr><tr><td>f<sup>Lyα</sup><sub>esc</sub></td><td>0.012</td><td>0.0065</td><td>0.059</td><td>Lyα escape fraction</td></tr><tr><td>R</td><td>0.072</td><td>0.29</td><td>2.0</td><td>“ext corrected” f<sup>Lyα</sup><sub>esc</sub></td></tr></tbody></table>

**Notes.** $F_{H\alpha}$ is the Hα flux, corrected for foreground reddening. $v_{H\alpha}$ and $v_{star}$ are velocities derived from Hα emission and UV photospheric absorption lines, respectively, and are with respect to the systemic velocity of 6194 km s⁻¹. Uncertainties are <3 km s⁻¹ and <10 km s⁻¹, respectively. The Hα line width, $\sigma_{H\alpha}$, has uncertainties <4 km s⁻¹, and EW(Hα) is accurate to <3%. (Hα/Hβ)raw is the observed line ratio, and (Hα/Hβ)corr the one corrected for foreground extinction (accurate to ~0.03%). $F_{Lyα}$ is the Lyα flux, corrected for foreground extinction; the uncertainties are dominated by the intrinsic COS calibration and the continuum determination and is accurate to ~10%. $W_{Lyα}^{red}$ has uncertainties of ~10% while peak separations have uncertainties of ~20 km s⁻¹. Red peak asymmetries are accurate to ~5%. The uncertainties on $f_{Lyα}^{esc}$ and R are dominated by the Lyα accuracy and is hence ~10%.

a We report the peak separation in knot C, assuming the small blue bump is the blue peak, although it is more likely due to a P Cygni profile. See the main text for further details.

MUSE spectral cube with the COS aperture vignetting function (Goudfrooij et al. 2012). Then, 1D MUSE spectra were extracted by integrating the flux over the vignetted aperture. Hence, these MUSE extractions spatially match the COS spectra for each of the knots and were used to measure emission-line fluxes of Hα, Hβ, [O III], [S II], and [S III]. To determine the Hα velocity and line width, we fitted the Hα profile with a single Gaussian and for the line width ($\sigma_{H\alpha} = \text{FWHM}/2.355$) subtracted the MUSE line-spread function in quadrature (Menacho et al. 2019). The line fluxes were corrected for foreground Milky Way (MW) reddening of $E(B-V) = 0.010$ (NED<fnref n="13" />; Schlafly & Finkbeiner 2011). The statistical uncertainties on the line fluxes are <0.03% for Hα, Hβ, [O III]; and <0.5% for the [S II] and [S III] lines. Various line ratios (all accurate to <1%) are given in Table 1. For [S III]/[S II] = ([S III]9531 + 9069)/([S II]6717 + 6730), the 9531 transition was not in the wavelength region observed, and we used [S III]9532 = 2.469 [S III]9069 from Osterbrock & Ferland (2006). The total uncertainties are dominated by the absolute MUSE flux calibration (good to a few percent). The MUSE spectra for the three knots are shown in Figure 3.

From the corrected Hα/Hβ ratio, we derived the nebular reddening expressed as $E(B-V)_n$ by adopting the

Fitzpatrick (1999) extinction law (with $R_V = 3.1$ and $k_{H\beta} - k_{H\alpha} = 1.23$) and an intrinsic Hα/Hβ ratio of 2.84 (appropriate for case B, $T_e \gtrsim 10$ K; Osterbrock & Ferland 2006). The results are presented in Table 1. All three knots have appreciable nebular extinction, but both knots A and C have blue UV continuum slopes, suggesting low stellar reddening (see Section 4.2 for further discussion).

Knot B has the highest Hα flux, the highest EW(Hα), and highest reddening. This is obviously where the ionizing photon production rate is highest. Knot A also appears very young and contains many massive young star clusters (Adamo et al. 2010). Both knots A and B have high ionization as judged by [O III]/Hβ and [S III]/[S II]. Knot C has lower EW(Hα) and ionization but a higher electron temperature and lower metallicity (James et al. 2013; Menacho et al. 2021). While compact, with $r_{eff} \approx 40$ pc (Östlin et al. 2015), it is clearly more extended than the compact star clusters making up knots A and B, and resembles more a nuclear star cluster and likely has experienced an extended star formation history.

## 3.2. Lyα Emission

In addition to the previously known emission from knot C (Rivera-Thorsen et al. 2017), we detect prominent, though fainter, Lyα also from knots A and B. The Lyα profiles are presented in Figure 4 (where the Lyα flux is given in units of the local continuum, normalized to 1 at rest wavelength $\lambda_0 = 1230$ Å). The three knots differ in their Lyα properties. The Lyα profile for knot C (previously published in Rivera-Thorsen et al. 2017) shows a strong red peak and a hint of a blue peak, which, however, resembles a P Cygni profile more. Knots A and B show distinct blue secondary peaks and, in addition, broad continuum absorption near Lyα. We measure the velocities of the red and blue peaks and the width of the red peak, and determine the peak separation (see Table 1, where we also include the asymmetry of the red peak as defined by Kakiichi & Gronke 2021 for future reference). For the purpose of measuring the Lyα escape fraction locally from the three knots, we integrated the Lyα flux using the local continuum located at approximately ±6 Å, as we want to evaluate the amount of Lyα photons that are directly emerging from each knot. Hence not all broad continuum absorption is accounted for. The flux uncertainties are dominated by the calibration accuracy of COS, which is quoted to be 5% (Hirschauer et al. 2021, although the accuracy for extended sources is somewhat worse), and the continuum placement and is overall estimated to be <10%. The Lyα fluxes were corrected for MW reddening with the same recipe as above (Fitzpatrick 1999; Schlafly & Finkbeiner 2011), resulting in an upwards correction of 9.8%. The effective resolution of COS depends on the spatial structure of the source, and because Lyα is in general extended, the resolution is degraded to $R \sim 2000$ (FWHM 150 km s⁻¹; Rivera-Thorsen et al. 2017), but may differ somewhat from source to source depending on the structure of the emission. In any case, the nonzero Lyα emission seen at zero velocity in knots A and B is likely a product of the limited spectral resolution and not a signature of Lyα photons at line center directly escaping without being scattered (see Figure 3 in Verhamme et al. 2015).

Knot A has a broad red Lyα peak and a peak separation of >500 km s⁻¹, suggesting scattering at relatively high H I column densities. Knot B has a similarly broad red peak but a smaller peak separation compared to knot A. Hence, the H I

<footnote n="13">¹³ NASA/IPAC Extragalactic Database: http://ned.ipac.caltech.edu.</footnote>
<figure><img src="image_5.png" /><figcaption>Figure 3. MUSE spectra for the three knots, matched to the COS apertures. The abscissa shows the rest wavelength over the range 4500–6900 Å. Prominent lines and the location of the Wolf-Rayet bumps are indicated. On the right, we show zoom-ins for the regions around Hα and [NII].</figcaption></figure>

<figure><img src="image_6.png" /><figcaption>Figure 4. Lyα and ISM line spectra of (from left to right) knots A, B, and C. The upper panels show the Lyα profiles (normalized such that the continuum at 1230 Å is 1). The middle panels show the average profile of the Si II absorption lines in black, normalized to 1 in the continuum, and in orange “1 − f$_{c}$” (1 minus the covering fraction, i.e., the noncovering fraction) resulting from the AOD analysis. The bottom panel shows the average Si IV profile in black, and in green “1 − f$_{c}$” for Si IV. The shaded regions in all panels represent the 1σ uncertainties, and for “1 − f$_{c}$,” the confidence intervals as described in Rivera-Thorsen et al. (2015). The velocity zero point is that derived from Hα for each knot.</figcaption></figure>

optical depth is likely comparable to or somewhat smaller than that in knot A. Knot C has the narrowest Lyα profile of the three knots, suggesting a lower H I scattering column density (Verhamme et al. 2006, 2015).

### 3.3. Lyα Escape Fractions

We use the MUSE Hα flux and reddening together with the COS Lyα flux measurements to calculate the Lyα escape fraction for each knot within the COS aperture according to

$$f_{\text{esc}}^{\text{Ly}\alpha} = F_{\text{Ly}\alpha} / (8.4 \cdot F_{\text{H}\alpha} \cdot 10^{0.4 \cdot E(B-V)_n \cdot k_{6563}})$$

(see Hayes et al. 2011), where 8.4 is the case B Lyα/Hα value for $T \gtrsim 10^4$ K and $n_e \sim 100$ cm⁻³ (Brocklehurst 1971; Hayes 2019), $E(B-V)_n$ is the nebular reddening derived from Hα/Hβ, and $k_{6563} = 2.32$ is the relative extinction coefficient at Hα. The derived escape fractions for knots A, B, and C are presented in Table 1. Knots A and B have low escape fractions ($\lesssim 1\%$), while knot C has $f_{\text{esc}}^{\text{Lyα}} = 6\%$. Because, due to resonant scattering, a fraction of the Lyα photons originating in the knots may escape outside the aperture rather than being absorbed by dust, the actual escape fractions are likely higher than the numbers given above.
We also compute $\mathcal{R}$ as in Östlin et al. (2009):

$$ \mathcal{R} = \frac{1}{8.4} \frac{F_{Ly\alpha}}{F_{H\alpha}} \cdot 10^{0.4 \cdot E(B-V)_n \cdot (k_{6563} - k_{1216})}, $$

which measures how the extinction-corrected $Ly\alpha/H\alpha$ ratio deviates from the prediction from a simple dust screen geometry where $Ly\alpha$ and $H\alpha$ are exposed to the same amount of dust. $\mathcal{R}$ can be viewed as the extinction-corrected escape fraction: If $Ly\alpha$ is only attenuated by dust and subject to the same dust optical depth as $H\alpha$, then $\mathcal{R}$ should be equal to 1. We take $(k_{6563} - k_{1216} = 8.6)$ from the starburst law of Calzetti et al. (2000), and the $\mathcal{R}$ values are given in Table 1. In contrast to $H\alpha$, the extinction correction for $Ly\alpha$ is sensitively dependent on the chosen extinction law: the Fitzpatrick (1999), Cardelli et al. (1989), and Calzetti et al. (2000) laws all give similar results (to within 20%), while the Small Magellanic Cloud (SMC) law (Prevot et al. 1984) gives $\mathcal{R}$ values that are several times higher (0.2, 2, and 12 for knots A, B, and C, respectively). This suggests that the SMC law (derived for stars) does not provide a fair estimate of the attenuation suffered by $Ly\alpha$ photons.

The uncertainties on $f_{esc}^{Ly\alpha}$ and $\mathcal{R}$ are dominated by the $Ly\alpha$ flux uncertainties and the accuracy of the COS vignetting function (i.e., how well matched the effective $Ly\alpha$ and $H\alpha$ apertures are).

### 3.4. Broad H I Absorption and Voigt Fitting

Both knots A and B show a decreased continuum level close to $Ly\alpha$. In order to assess whether this could be due to damped neutral hydrogen absorption, we performed Voigt profile fitting with a single component to the absorption wings of $Ly\alpha$. We fix the absorption velocity for the atomic gas to that measured from the Si II 1260 absorption line but leave the Doppler parameter and column density free. Because of many (sometimes blended) absorption components in the vicinity of $Ly\alpha$ (NV P Cyg, Si II ISM lines, and the same ISM features in the MW), we pay careful attention to the continuum placement, isolating regions we believe to be clean. On the red side, we adopt a region between the emission peak of the N V 1240 P cyg and the Si II 1260 ISM line, while on the blue side, we adopt a narrow region between the stellar CIII 1175 line in Haro 11 and the Si III 1206 Å absorption line in the MW (shifted back to 1182 Å in the rest-frame representation of Figure 5). For the Voigt profile fitting itself, we also adopt regions that we believe to be free from discrete interstellar absorption lines and the extended blue wing of the N V P Cyg, which when blending with the $Ly\alpha$ absorption may appear relatively flat as an artificially reduced continuum at ~1230 Å.

The results are shown in Figure 5. For knot A, we find a column density of $log(N_{H1}) = 20.7$, and for knot B, $log(N_{H1}) = 21.0$. No sign of damping wings is seen in knot C, and we could not produce a viable fit.

### 3.5. Stellar Velocities

We used photospheric lines from C III (1247 and 1776 Å) to measure the velocities of the UV-bright stellar components, i.e., the ionizing sources. For knots A and B, these agree with the velocities inferred from $H\alpha$, while for knot C, the stellar velocities are 33 km s⁻¹ higher; see Table 1. Hence, the $H\alpha$-emitting gas toward knot C is on average blueshifted with

<figure><img src="image_7.png" /><figcaption>Figure 5. Voigt profile fitting of the $Ly\alpha$ absorption wings in knots A (upper) and B (lower). Orange shading represent the regions for continuum normalization and green shading the region for fitting of the Voigt profiles. These fits provide column densities of 20.7 and 21.0 cm⁻² for knots A and B, respectively.</figcaption></figure>

respect to the ionizing stars, but the difference is just one-fifth of the $H\alpha$ line width ($FWHM_{H\alpha} = 2.355 \sigma_{H\alpha} = 150$ km s⁻¹; see Table 1) and does not affect any of the analysis in this paper.

### 3.6. ISM Absorption Lines, Covering Fraction, Column Densities

The spectra show the presence of metal absorption lines arising both from the Milky Way and Haro 11, notably several transitions of singly and triply ionized silicon. Because the ionization potential for silicon (8.2 eV) is lower than for hydrogen (13.6 eV), Si II primarily probes the part of the ISM where hydrogen is neutral. Other probes of the neutral ISM in the COS spectra of Haro 11 are O I 1302 and C II 1334 Å. We show the absorption profiles of these so called Low Ionization State (LIS) lines in Figure 6. Conversely, to triply ionize silicon, photons with energy 33.5 eV are required and Si IV hence traces highly ionized gas (similar to [O III]).

We analyze the UV absorption-line spectrum using the same methods as in Rivera-Thorsen et al. (2015, 2017) based on the AOD method (Savage & Sembach 1991; Jones et al. 2013) utilizing three Si II lines (1260, 1304, 1526 Å) with the same lower state but different oscillator strengths to solve for the column density and covering fraction as a function of velocity. The use of three lines with different strengths also improves the estimates of column densities if the lines are saturated. For the warm ionized gas, we perform the same analysis using two Si IV transitions (1393, 1402 Å). The results from knot C have been presented in Rivera-Thorsen et al. (2017). Here, we perform the same analysis for knots A and B (and for consistency redo the one for knot C, using the updated $H\alpha$ velocity in Table 1).

In each velocity bin (i), the residual intensity (in units of the continuum intensity) is given by

$$ \frac{I_i}{I_0} = 1 - f_{c,i} (1 - e^{-\tau_i}), $$
<figure><img src="image_8.png" /><figcaption>**Figure 6.** Selected LIS lines on a common velocity scale, for (from top to bottom) knots A, B, and C (respectively). The velocity zero point for each knot is that determined from the Hα emission line and given in Table 1. The LIS lines are color coded and shown as solid lines, while the dashed parts of the lines show adjacent features. E.g., the O I 1302 transition (solid green) shows the neighboring Si II 1304 line (dashed green, to the right) and vice versa (O I 1302 line shown as a dashed red line to the left of Si II 1304). The feature (dashed red) seen on the right side of Si II 1304 (solid red) in all panels is due to C II 1334 from the Milky Way. The shaded regions close to zero intensity show the error spectrum but is too small to be seen in knot A.</figcaption></figure>

where $f_{c,i}$ is the covering fraction, and $\tau_i$ is the optical depth, given by

$$ \tau_i = f\lambda \frac{\pi e^2}{m_e c} N_i = f\lambda \frac{N_i}{3.768 \times 10^{14}}, $$

where f is the oscillator strength for the transition in question and $N_i$ the column density for the given velocity bin and ion. The fact that the Si II line with the highest oscillator strength (1260 Å) has in general a deeper profile than the other two

indicates that the latter have lower optical depths and hence are not completely saturated; see Figure 6.

Because absorption occurs from the bright compact UV emission, the spectral resolution is only moderately degraded; Rivera-Thorsen et al. (2017) estimate $R \approx 10,000$ or $30 \text{ km s}^{-1}$ for knot C, but we note that knots A and B have more complicated UV morphologies, and hence possibly somewhat lower resolution. We investigated the COS/NUV acquisition images (see Figure 2) and collapsed them along the cross-dispersion direction to find the FWHM (and half-light width, HLW) in the dispersion direction and obtained 0.′34 and 0.′31 (HLW = 0.′34 and 0.′53) for knots A and B, respectively.

The COS/FUV beam is considerably aberrated, and point sources have FWHM in the cross-dispersion direction of 0.′5–1.′5 (COS ISR 2018-08, Fox 2018), and sources with FWHM $\ge 0.′6$ should be regarded as extended (COS IHB; Hirschauer et al. 2021). Hence, the continuum spectral resolution of knots A and B should be only marginally degraded by the low-surface-brightness wings. On the other hand, the spectral resolution at LP4 is degraded to approximately 13,500 (COS ISR 2018-07; Fox 2018). Hence, a spectral resolution of $R \sim 10,000$ is a fair approximation for each of the knots. Inspection of some Milky Way lines (e.g., Si II 1260 and C II 1335) indeed verifies that the spectra toward all three knots appear to have the same resolution. Our velocity sampling of $17 \text{ km s}^{-1}$ hence oversamples the effective resolution by a factor of $\sim 2$ (i.e., approximately Nyquist sampling) and is close to the thermal line width. The results of the AOD analysis are shown in Figures 4 and 7 and Table 2. The derived covering fractions vary with velocity and show a maximum close to the deepest Si II absorption.

The neutral ISM traced by Si II, which is sensitive to H I gas with column density $\ge 10^{17} \text{cm}^{-2}$ (Rivera-Thorsen et al. 2017), displays a velocity width ($W_{90\%}$ encompassing 90% of the absorption) of more than $400 \text{ km s}^{-1}$ for all three knots. This gas is very optically thick to Lyα but has $\tau_{LyC} \sim 1$ for a single-velocity bin. The deepest absorption is consistently seen at velocities that are blueshifted relative to Hα. Notably, the SiII profiles mimic the shape of the Lyα profiles between the two peaks, a behavior also seen in the LARS sample (Rivera-Thorsen et al. 2015), i.e., Lyα intensity quickly drops at velocities where there is a large covering of neutral gas optically thick to Lyα. The fact that we see neutral ISM absorption over a wide range of velocities and with varying (nonunity) covering fractions demonstrates that we have an outflowing clumpy neutral medium along each sight line, rather than a single neutral shell (Rivera-Thorsen et al. 2015). Knot A shows a flat covering fraction distribution for $v \le -70 \text{ km s}^{-1}$ with $f_c > 0.9$ in 12 consecutive velocity bins (i.e., over a range $200 \text{ km s}^{-1}$), and with 9 of these having $f_c > 0.95$. The covering fraction distribution in knot B is less flat, but has $f_c > 0.9$ in seven consecutive bins for $v \le -50 \text{ km s}^{-1}$, with the first three having $f_c > 0.95$. The highest covering fractions in knot C is seen at $v \sim -90 \text{ km s}^{-1}$ with two bins having $f_c \ge 0.6$.

The Si IV absorption has approximately the same shape as the Si II profile, and the covering fraction distribution mimics that of Si II, demonstrating that the outflowing medium is multiphase, consisting of both neutral and ionized gas. However, for all knots, there is a Si IV low-level absorption tail extending to larger blueshifted velocities than seen in Si II, although the flat shape of the spectrum as it blends into the continuum noise makes an exact measurement of the maximum
<figure><img src="image_9.png" /><figcaption>**Figure 7.** The upper panels shows the comparison of the Si II column densities derived from the AOD analysis of knots A (green), B (orange), and C (blue). The middle panels show the inferred Si II covering fractions. The bottom panels show the profiles of the strongest Si II transition ($\lambda_{rest} = 1260$ Å). Left: the reference velocity is here the individual values for each of knots A, B, and C. Right: the reference velocity for all knots is here the systemic velocity of knot C.</figcaption></figure>

<table><caption>**Table 2**
Results for Absorption Lines of Knots A, B, and C</caption><thead><tr><th>Quantity</th><th>Knot A</th><th>Knot B</th><th>Knot C</th><th>Unit</th><th>Comment</th></tr></thead><tbody><tr><td>v<sub>min</sub></td><td>-159 <sup>+17</sup><sub>-1</sub></td><td>-85 <sup>+17</sup><sub>-17</sub></td><td>-139 <sup>+17</sup><sub>-51</sub></td><td>km s<sup>-1</sup></td><td>Si II</td></tr><tr><td>v<sub>int</sub></td><td>-228 <sup>+1</sup><sub>-1</sub></td><td>-137 <sup>+17</sup><sub>-1</sub></td><td>-169 <sup>+1</sup><sub>-1</sub></td><td>km s<sup>-1</sup></td><td>Si II</td></tr><tr><td>v<sub>5%</sub></td><td>-57 <sup>+1</sup><sub>-17</sub></td><td>51 <sup>+17</sup><sub>-17</sub></td><td>37 <sup>+1</sup><sub>-17</sub></td><td>km s<sup>-1</sup></td><td>Si II</td></tr><tr><td>v<sub>95%</sub></td><td>-468 <sup>+17</sup><sub>-17</sub></td><td>-377 <sup>+17</sup><sub>-51</sub></td><td>-425 <sup>+34</sup><sub>-34</sub></td><td>km s<sup>-1</sup></td><td>Si II</td></tr><tr><td>W<sub>90%</sub></td><td>411 <sup>+17</sup><sub>-34</sub></td><td>428 <sup>+51</sup><sub>-34</sub></td><td>462 <sup>+34</sup><sub>-34</sub></td><td>km s<sup>-1</sup></td><td>Si II</td></tr><tr><td>f<sub>c,vmin</sub></td><td>0.95 <sup>+0.05</sup><sub>-0.04</sub></td><td>0.96 <sup>+0.04</sup><sub>-0.14</sub></td><td>0.55 <sup>+0.05</sup><sub>-0.04</sub></td><td></td><td>Si II</td></tr><tr><td>f<sub>c,vsys</sub></td><td>0.32 <sup>+0.06</sup><sub>-0.10</sub></td><td>0.87 <sup>+0.13</sup><sub>-0.18</sub></td><td>0.51 <sup>+0.02</sup><sub>-0.07</sub></td><td></td><td>Si II</td></tr><tr><td>log(N<sub>Si II</sub>, v<sub>min</sub>)</td><td>13.0 <sup>+∞</sup><sub>-0.2</sub></td><td>12.8 <sup>+∞</sup><sub>-0.2</sub></td><td>12.4 <sup>+0.1</sup><sub>-0.1</sub></td><td>cm<sup>-2</sup>/km s<sup>-1</sup></td><td>Si II</td></tr><tr><td>log(N<sub>Si II</sub>, v<sub>sys</sub>)</td><td>12.6 <sup>+∞</sup><sub>-0.4</sub></td><td>12.6 <sup>+∞</sup><sub>-0.3</sub></td><td>12.1 <sup>+0.1</sup><sub>-0.1</sub></td><td>cm<sup>-2</sup>/km s<sup>-1</sup></td><td>Si II</td></tr><tr><td>log(Σf<sub>c</sub> N<sub>Si II</sub>)</td><td>15.2</td><td>15.3</td><td>14.7</td><td>cm<sup>-2</sup></td><td>Si II</td></tr><tr><td>f<sub>c,vmin</sub></td><td>0.93 <sup>+0.07</sup><sub>-0.04</sub></td><td>1.0 <sup>+0</sup><sub>-0.22</sub></td><td>0.94 <sup>+0.06</sup><sub>-0.1</sub></td><td></td><td>Si IV at v<sub>min</sub> for Si II</td></tr><tr><td>f<sub>c,vsys</sub></td><td>0.40 <sup>+0.19</sup><sub>-0.07</sub></td><td>0.77 <sup>+0.23</sup><sub>-0.11</sub></td><td>0.98 <sup>+0.02</sup><sub>-0.15</sub></td><td></td><td>Si IV at v<sub>sys</sub></td></tr><tr><td>log(N<sub>Si IV</sub>, v<sub>min</sub>)</td><td>12.8 <sup>+∞</sup><sub>-0.4</sub></td><td>12.3 <sup>+∞</sup><sub>-0.4</sub></td><td>12.2 <sup>+0.2</sup><sub>-0.1</sub></td><td>cm<sup>-2</sup>/km s<sup>-1</sup></td><td>Si IV at v<sub>min</sub> for Si II</td></tr><tr><td>log(N<sub>Si IV</sub>, v<sub>sys</sub>)</td><td>13.6 <sup>+∞</sup><sub>-1.7</sub></td><td>12.6 <sup>+∞</sup><sub>-0.5</sub></td><td>12.0 <sup>+0.2</sup><sub>-0.1</sub></td><td>cm<sup>-2</sup>/km s<sup>-1</sup></td><td>Si IV at v<sub>sys</sub></td></tr><tr><td>log(Σf<sub>c</sub> N<sub>Si IV</sub>)</td><td>15.0</td><td>14.8</td><td>14.8</td><td>cm<sup>-2</sup></td><td>Si IV</td></tr></tbody></table>

**Note.** The first four lines give kinematic properties for the Si II lines, where $v_{min}$ is the velocity at the deepest absorption and $v_{int}$ the absorption-weighted average velocity, with respect to the Hα velocity for each knot. Summing all Si II absorption, $v_{5%}$ and $v_{95%}$ is the velocity where 5% and 95% of the total absorption are encompassed, respectively. $w_{90\%} = v_{95\%} - v_{5\%}$, is our measure of the total Si II absorption velocity width. The 6th and 11th rows give the Si II and Si IV covering fractions, respectively, both at $v_{min}$ for Si II, while rows 7 and 12 give $f_c$ for the systemic velocity of each knot. The errors represent the confidence intervals based on Monte Carlo simulations (Riviera-Thorsen et al. 2015, 2017). For the column densities, the unit is cm$^{-2}$ per km s$^{-1}$, and the upper limit is sometimes unconstrained (represented by $+\infty$). $\Sigma f_c N_{[Si II]}$ is the sum of the covering fraction times the column density and velocity bin width for Si II, and $\Sigma f_c N_{[Si IV]}$ is the same for Si IV. Given the very asymmetric confidence intervals for the column densities, it is not meaningful to present uncertainties on these quantities, which should be regarded to be quite approximate.

velocity difficult. The Si IV doublet is, however, also known to be strong in early-type (O and B) stars, where high-velocity, highly ionized winds are common, revealing blueshifted

Si IV absorption extending from $\sim -100$ to $\lesssim -1000$ km s$^{-1}$ (Grady et al. 1987; Howarth & Prinja 1989; Shepard et al. 2020). All three knots show P Cygni profiles in C IV and N V,
and though not obvious, could be there for Si IV as well. Hence, the Si IV absorption profiles toward the knots in Haro 11 may plausibly also have contributions from circumstellar absorption, and interpreting differences with respect to Si II should have this caveat in mind.

There are some subtle details worthy of mention here:

1. In knot B, the Si II and Si IV profiles follow each other closely, but there is significant absorption in both profiles at redshifted (compared to Hα) velocities up to +150 km s⁻¹. This can be understood in terms of bulk motions in the multiphase ISM; the Hα velocity dispersion (Table 1) corresponds to an FWHM of 210 km s⁻¹ so blueshifted components of this magnitude could be foreseen, even when accounting for the fact that Hα also samples gas on the far side of knot B. Another possibility is that parts of the clumpy medium are falling into knot B. The Si IV absorption extends to blueshifted velocities of at least -700 km s⁻¹.

2. In knot A, we see no neutral gas at redshifted velocities, but a redshifted Si IV component up to $v \lesssim 100$ km s⁻¹ (again consistent with the Hα velocity dispersion), which on the blue side extends to $v \lesssim -750$ km s⁻¹. The low covering of neutral gas at zero velocity ($f_c \approx 0.3$) leads to a narrow window, on the blue side of the Lyman limit, where LyC photons will not be absorbed. However, this window is just 30 km s⁻¹, corresponding to 0.1 Å, and hence insignificant for the escape of LyC.

3. In knot C, we see redshifted Si II absorption up to $v \sim +100$ km s⁻¹, but the Si IV absorption extends to +250 and -800 km s⁻¹ on the red and blue sides, respectively.

The inferred neutral ISM covering is significantly smaller in knot C than in A and B, but the ionized gas covering, while still smaller (the Si IV absorption never goes to zero), is not markedly different. Hence, while a similar amount of gas may be moving toward knot C as in knot B, the ionization/ pressure balance is such that the once neutral clouds have largely evaporated. The tail of blueshifted Si IV gas with $v < -500$ km s⁻¹ is seen in all knots but is not detected in Si II, suggesting that neutral clouds at those velocities have been ionized but could also be due to circumstellar absorption from hot stars with fast winds.

In Figure 7, we compare the column densities, covering fractions, and Si II 1260 line profiles of knots A, B, and C. In the left panel, the reference velocity is the systemic Hα velocity for each individual knot (6251, 6170, 6140 km s⁻¹ for knots A, B, and C, respectively). In the right panel, the zero point of the velocity scale is 6140 km s⁻¹, the systemic Hα velocity of knot C (we will return to this in Section 4.5). The derived column densities of Si II at $v_{min}$ (left panels of Figure 7) is $\log(N_{Si II}) = 13.0, 12.8, 12.4 (\text{cm}^{-2}/(\text{km s}^{-1}))$ for knots A, B, and C, respectively, or $\log(N_{Si II}) = 14.2, 14.0, 13.6 (\text{cm}^{-2})$ when integrated over a 17 km s⁻¹ velocity bin. However, the AOD method saturates at high optical depth, with the consequence that for several velocity bins in knots A and B, the upper limit on the column density is unconstrained, and there may also be smaller high-column-density clouds immersed in the moderately thick medium here probed. Following Rivera-Thorsen et al. (2017), we adopt a Si/H abundance ratio of $3.2 \times 10^{-6}$ implying that the neutral H column at minimum Si II intensity is of the order of

$\log(N_{H1}) \sim 19.7, 19.5, \text{ and } 19.1 (\text{cm}^{-2})$, for knots A, B, and C, respectively. The LyC becomes optically thick ($\tau_{LyC} = 1$) for $\log(N_H) = 17.2$, and the covered fractions of these velocity bins are completely opaque to both Hα and LyC. The covering fractions at $v_{min}$ are $f_c = 95\%, 96\%, \text{ and } 55\%$, and the LyC transmissions through the $v_{min}$ velocity bin hence 5%, 4%, and 45% for knots A, B, and C, respectively.

We can in principle follow the same procedure for each velocity bin. However, because we do not know how the covered versus noncovered regions correlate spatially between velocity bins, we are unable to estimate the LyC transmission integrated along each sight line. If there is no correlation, the integrated transmission approaches zero after just a small number of velocity bins.

We can estimate the number of absorbing clouds per line of sight required to produce the observed absorption and $f_c$ profiles. We first note that with the expected cloud sizes on the order of $\lesssim 1$ pc (McCourt et al. 2018) and effective aperture diameters of $\sim 0''3$ (area $\sim 10^4$ pc²), each aperture would contain $\gtrsim 10^4$ independent unresolved lines of sight that could hit a cloud. If one assumes optically thick absorbers and a reasonable transonic velocity dispersion inside the cold $\sim 10^4$ K medium in which Si II is embedded,<fnref n="14" /> one can estimate the expectation value for the number of absorbers for a sight line toward the knots as

$$N_{\text{abs.}} \approx \sum_{i} f_{c,i} \Delta v / w_{\text{abs}},$$

where $\Delta v$ is the spectral sampling (17 km s⁻¹), and $w_{abs}$ the velocity width of an optically thick cloud. For $w_{abs}$, we take the absorption equivalent width $w$ in velocity units which in the optically thick regime is $w = 2b\sqrt{\ln \tau_0}$ (Petitjean 1998), where $b$ is the Doppler width (13 km s⁻¹ for $10^4$ K) and $\tau_0$ the optical depth at line center. To estimate $\tau_0$, we take the average Si II column density ($\log N = 12.7 \text{ cm}^{-2}/\text{km s}^{-1}$) times the spectral resolution ($\sim 2 \times \Delta v$) and the oscillator strength for Si II$_{1260}$ to get $\tau_0 \sim 30$ and $w_{abs} \sim 50$ km s⁻¹ ($\sim 0.2$ Å). This yields $N_{abs.} \{8, 8, 4\}$ for clump A, B, and C, respectively. Assuming furthermore an uncorrelated (Poissonian) distribution of the absorbers in space, the probability of not encountering one is simply $e^{-N_{abs.}}$ which yields $\langle 0.1\% \rangle, \langle 0.1\% \rangle, \sim 2\%$ free lines of sight for the three clumps, which can be regarded as lower limits for the escape fractions.

If, on the other hand, the noncovered part of the aperture is spatially correlated between velocity channels, there can be a significant net LyC transmission. Spatially correlated “holes” would be expected if these are produced by either mechanical feedback (winds) or carved out by ionizing photons, so this is more likely in the current context, but the correlation is unlikely to be perfect. In the extreme case of perfect correlation, the escape fraction would be equal to that of the most covered velocity bin ($f_c = 0.98_{-0.05}^{+0.02}$, $1.0_{-0.15}^{+0.15}$, $0.65_{-0.07}^{+0.02}$), hence escape fractions of $2_{-2}^{+5}$, $0_{-0}^{+15}$, and $35_{-2}^{+7}\%$, respectively for knots A, B, and C.

With this caution in mind, and noting that some channels have unbounded upper limits on the column densities, we can

<footnote n="14">¹⁴ Note that while this is a fair assumption for the ISM, absorption-line studies of the circumgalactic medium find routinely much smaller Doppler parameters consistent with a purely thermal broadening (e.g., Churchill et al. 2020).</footnote>
make a rough estimate for the effective total H I columns contained in the solid angle toward each knot by taking

$$N = \sum_{i} N_{i} \times f_{c,i} \times \Delta v,$$

i.e., the sum of the column densities times the covering fraction over the range of velocities with reliable values (range [-500:0], [-440:150], and [-430:120] km s⁻¹ for knots A, B, and C, respectively; interpolating over unreliable bins, such as +10 for knot A and -220 for knot C). The results<fnref n="15" /> are presented in Table 2. The integrated $f_c \times N$ of knots A and B corresponds to $\log(f_c \times N_H) = 20.7$ (in rough agreement with the estimates from the Voigt profile fitting) and $20.2 \log(cm^{-2})$ for knot C. The estimate for knot C could be on the low side by ~0.3 dex due to the lower O/H; there is, however, no sign of broad H I absorption near $Ly\alpha$, and we do not know whether the lower O/H applies over the full effective line-of-sight extent (~10 kpc) or is local to knot C.

We also summed $N_{i,Si IV} \times f_{c,i}$ over the observed velocity width to arrive at similar, but somewhat lower, values as for Si II. However, this slightly underestimates the column density, because the ionized medium also contains some Si III and a little Si V. With a residual neutral hydrogen fraction in the ionized outflow component of $<10^{-5}$ (Gray et al. 2019), the integrated paths through the Si IV gas would be optically thick for $Ly\alpha$; hence, it would scatter also in this medium, removing $Ly\alpha$ from line center (see Verhamme et al. 2015), but thin ($\tau < 0.1$) to ionizing photons.

In addition to the Si II and Si IV transitions used for the AOD analysis, we detect clear LIS absorption in O I 1302 and C II 1334+1335, and for knots A and B, also S II 1250+1253. The O I line is of particular interest as the ionization potentials of H I and O I are both 13.6 eV, and the matter probed must come from the neutral hydrogen phase. For knots A and B, the O I line has the same shape, width, and depth (going to ~0 intensity) as the neighboring Si II 1304 line (see Figure 6), while for knot C, the width of O I is similar to Si II, although the absorption is much shallower (~0.8 times the local continuum for O I compared to ~0.6 for Si II). This lends support to the conclusion that the covering fraction toward knot C is significantly smaller than for knots A and B.

The O I 1302 line has an oscillator strength of 0.049 compared to 0.086 for the neighboring Si II 1304 transition. However, oxygen is significantly more abundant than silicon (typically $\log(Si/O) = -1.7$; Garnett et al. 1995) so one would expect the oxygen line to be stronger (as seen in knot B), unless both are totally optically thick in which case they should be equally strong (as seen in knot A). The fact that the O I line is weaker in knot C might indicate that Si II here partly arises in ionized gas and hence overestimates the neutral gas column and covering fraction. Because the ionization ([S III]/[S II]) is much lower in knot C, a significant fraction of [S II] and Si II may come from a phase where hydrogen is ionized. Thus, the neutral hydrogen covering fraction in knot C could be closer to 20%.

# 4. Discussion

## 4.1. Ionization Structure and Results by Keenan et al.

Keenan et al. (2017) used HST narrowband images to study the spatial pattern of ionization through [O III]/[O II], [O II]/H$\alpha$, and [O III]/H$\alpha$. At the position of knot A, [O II]/H$\alpha$ is low and [O III]/H$\alpha$ high, suggesting a low optical depth along the line of sight, and they argue that knot A may therefore be the source of LyC leakage from Haro 11.

We deem it interesting here to compare to the velocity sliced [O III]/H$\alpha$ mapping based on MUSE data and presented in Menacho et al. (2019). While the ionization is indeed very high around knot A, many of these structures seem to extend in the plane toward S and W. Very high ionization ([O III]/H$\alpha \gtrsim 3)$ is prominent >1 kpc SW of knot A from -300 to -150 km s⁻¹ (and also at +250 km s⁻¹; Menacho et al. 2019, Figure 3) but the ionization at the actual center of knot A is moderate ([O III]/H$\alpha \gtrsim 1). Hence, this structure is likely outflowing, created by ionizing photons from knot A and optically thin to LyC photons, but pointing toward us at some angle, not directly into the line of sight.

However, between -100 and 0 km s⁻¹, there is a point-like source of high ionization in Menacho et al. (2019), ~0.5 kpc W of the center of knot A (hence within the knot), and it reappears as highly ionized at $v > 250$ to $350$ km s⁻¹. It can be identified with the westernmost star cluster of knot A (see Figure 2, where it is indicated by a blue circle), and in Figure 1 as the bright red (indicating strong $H\alpha$ emission) source near the western edge of the COS aperture. The source is also clearly seen in the work of Keenan et al. (2017): in Figure 1 as the bright rightmost source in the box labeled knot A and in the central panel of Figure 5 as the source in knot A with [O III]/[O II] ~ 10. This source can also be identified with one of the young massive star clusters (cluster #15) in Adamo et al. (2010) and has an estimated age of 3.5 Myr and a mass of $5 \times 10^5 M_{\odot}$. The $H\alpha$ velocity at this location is ~+80 km s⁻¹ (compared to the systemic velocity 6194 km s⁻¹ of Menacho et al. 2021, coincident with the dusty arm), while the average for knot A is $57$ km s⁻¹. In the kinematic study of Östlin et al. (2015), it was shown that the stellar and ionized gas velocities broadly agree in knot A but also that the ionized gas here has several distinct velocity components due to the projection onto the more redshifted dusty arm (Menacho et al. 2021; referred to as the “ear” in Östlin et al. 2015). Hence, it is probable that this cluster actually lies behind knot A, in the dusty arm. However, it also coincides with a local minimum in the nebular reddening ($E(B-V)_n \approx 0.16$), slightly lower than the rest of knot A and the dusty arm.

The ionization of this source is moderate at its systemic velocity but increasing for more redshifted velocities. The ionization drops at yet higher velocities, which may due to blending with the overall halo gas that has somewhat lower ionization due to exhaustion of more energetic photons (>35 eV) capable of doubly ionizing oxygen.

Cluster #15 is the strongest $H\alpha$ source within the knot A aperture and accounts for approximately 30% of the $H\alpha$ flux in the COS aperture, but only contributes ~10% of the far-UV flux. This means that the UV continuum and $H\alpha$ will have slightly different luminosity weighted mean velocities, and, as a consequence, we might expect a deficiency of ISM absorption for velocities $v > -20$ km s⁻¹, which may affect the Si II absorption in knot A close to 0 km s⁻¹.

<footnote n="15">¹⁵ The results differ from those of Rivera-Thorsen et al. (2017), where the integration over the width of each velocity bin, amounting to a factor of 17 corresponding to the width in km s⁻¹ per bin, was omitted.</footnote>
Another high-ionization structure, in the form of a cone, is seen extending north from knot B at −50 to +100 with an opening angle of ~5° and spatially extending to the boundary of the galaxy, where the ionization even increases. This is most likely another LyC escape channel directed roughly perpendicular to the line of sight (Menacho et al. 2019; Figure 4.)

Knot C shows low to moderate ionization at all velocities. Knot C is spatially resolved, hence not a bona fide star cluster, but more like a very UV-bright galaxy core or nuclear star cluster. Presumably, it has experienced an extended star formation history, perhaps for the past 40 Myr when the massive star cluster formation was triggered (Adamo et al. 2010). Despite this, the electron temperature is highest in knot C (see further discussion on this in Menacho et al. 2021.)

## 4.2. Dust Geometry

We return now to $\mathcal{R}$ the “extinction-corrected” Ly$\alpha$ escape fractions. In general, galaxies are observed to have $\mathcal{R}$ significantly lower than 1 (Giavalisco et al. 1996; Östlin et al. 2009; Hayes et al. 2011; Atek et al. 2014), indicating that part of Ly$\alpha$ has been scattered out of the line of sight and/or that scattering has increased the effective path length and hence the effective dust column. In the rare case where $\mathcal{R} > 1$, then either (i) Ly$\alpha$ has been scattered into the line of sight, or (ii) has experienced a lower dust column, or (iii) the dust geometry is not in the form of a screen, but mixed with the sources or in discrete clumps (Scarlata et al. 2009). Case (i) is not likely for spectra centered on bright UV continuum sources, such as knots A, B, and C, because the Ly$\alpha$ production outside the apertures is much lower than within them but is a reason why Ly$\alpha$/H$\alpha$ > 8.4 is observed in the halo of many galaxies including Haro 11 (Hayes et al. 2007; Östlin et al. 2009, 2014). Case (ii) can happen due to geometrical effects in the ISM, e.g., if dust is mainly contained in dense clouds and Ly$\alpha$ scatters on their surfaces but do not traverse them (Neufeld 1991; Laursen et al. 2013; Duval et al. 2014), or similarly, if there are ionized dust-poor channels where Ly$\alpha$ scatters on the interfaces. For case (iii), the wavelength dependence on the optical depth will be softened (Scarlata et al. 2009), effectively increasing Ly$\alpha$/H$\alpha$. It should be noted that cases (ii) and (iii) are not distinct: the dust is likely, and observed to be, inhomogeneously distributed but in a more complex fashion than having clouds of fixed opacity with no dust around them, and Ly$\alpha$ will always have more complicated radiative transfer since it is prone to scattering on H I.

From Table 1, we see that knot A has $\mathcal{R} \ll 1$, and hence, Ly$\alpha$ has been significantly attenuated by scattering, increasing the effective dust optical depth. For knot C, $\mathcal{R} > 1$ and Ly$\alpha$ must have been scattered through less dusty paths in a clumpy two-phase medium or through ionized channels. Those paths with low H I density would also be optically thin to LyC photons.

Interestingly, this scenario is also consistent with less neutral absorbers per line of sight for clump C as estimated in Section 3.6. In such a case, (ii) described above is more likely as then the majority of Ly$\alpha$ photons will escape via “random walk” between the clumps and, thus, are less susceptible to dust embedded within them. If, on the other hand, the number of absorbers increases, more Ly$\alpha$ photons will escape via

“frequency excursion” and cross the absorbing clouds and their dust (Gronke et al. 2017).

## 4.3. Stellar versus Nebular Reddening

The amount of reddening of the stellar component for young populations can be estimated from the UV continuum slope, $\beta$ ($f_{\lambda} \propto \lambda^{\beta}$) (Calzetti et al. 2000). To measure $\beta$, we used the derivation from Micheva et al. (2020) based on images in the ACS/SBC/F140LP, STIS/NUV/F25QTZ, and WFC3/UVIS/WFC3 filters and retrieved the value for each knot by integrating over the COS aperture after having applied the COS vignetting function (Goudfrooij et al. 2012), and correcting for foreground Milky Way extinction (as above), to obtain a $\beta$ representative for the observed COS spectrum of each knot. (see Table 1).

A value $\beta = -1.73$ indicates 0.85 mag of extinction at 1600 Å, and $E(B-V)_n \approx 0.19$ (using Equations (9) and (10) in Calzetti et al. 2000). Hence, for knot A, the values of $E(B-V)_n = 0.18$ and $\beta = -1.73$ are perfectly consistent with the Calzetti et al. (2000) prescription that $E(B-V)_s = 0.44E(B-V)_n$. For knot C, the reddening implied from $\beta$ ($A_{1600} \lesssim 0.5$) is a factor of 3.2 lower at 1600 Å than what $E(B-V)_n$ suggests ($A_{1600} \approx 1.6$). For knot B, the situation is the reverse, with $\beta$ implying 2.9 mag of extinction at 1600 Å, while $H\alpha/H\beta$ suggests $A_{1600} \approx 1.7$ mag. The dust geometries are quite complex in knots B and C, which is apparent from just looking at the HST multicolor image (Adamo et al. 2010), where dust lanes streak through the centers of knots B and C. Possibly, in knot C, the dusty nebular gas-emitting $H\alpha$ and $H\beta$ only covers a small part of the UV radiation (consistent with the low neutral gas-covering fraction), while in knot B, dust could preferentially be located in the UV luminous parts, surrounded by a bright emission-line nebula with lower dust content. Indeed, the $H\alpha/H\beta$ and $Br\gamma/H\beta$ measurements for knots B and C from XSHOOTER in a 0.9' wide slit (Guseva et al. 2012) differ significantly from our values, with the center of knot B having a higher extinction: $E(B-V)_n = 0.51$ and 0.60 from $H\alpha/H\beta$ and $Br\gamma/H\beta$, respectively (implying $A_{1600} = 2.3$, and 2.6, respectively), closer to the $\beta$ estimate. Knot C instead has a lower inferred $E(B-V)_n = 0.2$ (implying $A_{1600} = 0.9$). These differences corroborate the conclusion that the dust distribution in knots B and C are highly nonuniform (see also Menacho et al. 2021).

The above discussion used the Calzetti et al. (2000) law. If adopting instead the SMC law of Prevot et al. (1984), lower stellar extinctions (about half) would result, mainly because the selective extinction ($R_V$) is lower, and in this case, the far-UV extinctions inferred from $\beta$ and $H\alpha/H\beta$ in knot B would be largely consistent, while for knot C, the discrepancy would increase.

Anyway, an inhomogeneous dust distribution seems a like a natural explanation for both the small stellar reddening and high $\mathcal{R}$ of knot C and lends support to knot C as a probable source of LyC leakage. In knot B, geometry may have worked in the opposite direction to preferentially redden the UV continuum, and the fairly high $\mathcal{R}$ indicates that the average escaping Ly$\alpha$ photon has encountered less dust than the UV continuum. We finally note that the individual clusters in knot A (including #15) have bluer spectra ($\beta \lesssim -2$) than knot A integrated. The same applies for knot C, where the center has $\beta \approx -2.1$ (Micheva et al. 2020).
### 4.4. Ionizing Photon Budget

We now pose the question if any of the knots have the required UV flux to be a potential source of the LyC flux density observed with FUSE, which is $f_{\lambda}^{LyC} = 4 \pm 0.9 \times 10^{-15}$ erg s$^{-1}$ cm$^{-2}$ Å$^{-1}$ (Leitet et al. 2011).

The magnitude of the Lyman discontinuity depends on the age and metallicity of the stellar populations, which we are not addressing in detail in the current paper. We note, however, that knot B is the youngest with most clusters younger than 5 Myr, while knot A has a spread in age from a few to 10 Myr (Adamo et al. 2010), and knot C likely has had an extended star formation history for the past 40 Myr. Adopting a Starburst99 (Leitherer et al. 1999) model with a continuous SFR and metallicity $Z=0.008$ (the results are very similar for $Z=0.004$) with ages 10, 5, and 50 Myr leads to Lyman discontinuities of $f_{\lambda_{912}}^{+}/f_{\lambda_{912}}^{-}=2.9$, 2.3 and 3.8 for knots A, B, and C, respectively. Here, $f_{\lambda_{912}}^{+}/f_{\lambda_{912}}^{-}$ is actually defined as $f_{\lambda_{1080-1120}}/f_{\lambda_{870-900}}$ (Leitherer et al. 1999).

We use the observed continuum flux densities at 1265 Å and the $\beta$ values above to extrapolate to the expected flux densities at 1100 Å, and find $f_{\lambda_{912}}^{+}=2.2 \times 10^{-14}$, $5.2 \times 10^{-15}$, and $5.2 \times 10^{-14}$ erg s$^{-1}$ cm$^{-2}$ Å$^{-1}$ for knots A, B, and C, respectively. We can then make the rough estimates that knots A, B, and C would need to have escape fractions of ~53%, ~177%, and ~29%, respectively, to account for the measured LyC flux density. Hence, knot B seems quite unlikely as a major contributor to the observed LyC, and also knot A in regard to the high neutral ISM covering. Knot C requires the least intrinsic escape and in addition has the lowest inferred ISM covering (with an estimated LyC transmission of 2%–35%; see Section 3.6), making it plausible as a significant source of the escaping LyC.

Zooming in on the aforementioned cluster #15 (the westernmost cluster in knot A), could it account for a fraction of the observed LyC flux? For an age of 3.5 Myr, and a Starburst99 standard model (Leitherer et al. 1999) with instantaneous burst and $Z=0.004$, $f_{\lambda_{912}}^{+}/f_{\lambda_{912}}^{-}=2.9$, and even if 100% of its LyC photons would escape, it could not account for more than 25% of the observed LyC flux. Because this source has ~10% of the far-UV flux of knot A, a 100% local escape would require a neutral covering fraction in knot A of $f_c \lesssim 0.9$, while we derive $f_c > 0.95$ over a wide velocity range ($150 \text{ km s}^{-1}$). Moreover, the H$\alpha$ emission from this cluster is quite round and symmetric with a diameter of ~0.5 (as judged from the HST H$\alpha$ image) suggesting a quite high covering of the nebular gas. The equivalent width of H$\alpha$ at cluster #15 is indeed very high: EW(H$\alpha$) = 1490 Å integrated over the 0.5 diameter. This value is close to what Starburst99 predicts (1320 Å) for the age in question, and hence most LyC photons seem to result in ionization, implying $f_{\text{esc}}^{LyC} \ll 1$. We conclude that cluster #15 is not likely to contribute to the observed LyC flux. The eastern part of knot A has a more much more patchy and fractured H$\alpha$ emission, and while EW(H$\alpha$) is far from uniform, it is on average lower (450 Å). It does in this respect seem more likely than the other parts (than cluster #15) of knot A are leaking LyC photons to the SW parts of the galaxy and powering the high ionization seen there (Keenan et al. 2017; Menacho et al. 2019).

While knot B (based on the H$\alpha$ flux) produces more LyC photons than knots A and C together, it appears quite unlikely to be a major source of the detected LyC emission, because it

would require the LyC photons to have some exclusive escape paths not available to the UV continuum at $\lambda \sim 1200$ Å. Neither the ISM covering nor the extinction measurements suggest this to be the case. Hence, while there may also be other sources to the observed LyC flux (e.g., recombinations directly to the ground state from the ionized halo; Östlin et al. 2021, in preparation), and in principle, more than one knot could be leaking, at least we can rule out knot B as a significant source of the observed LyC in Haro 11. Knots A and C are both plausible sources, but the former requires a higher intrinsic escape fraction at odds with the ISM covering.

### 4.5. Implications from the Ly$\alpha$ Profiles and AOD Analysis

The connection between Ly$\alpha$ and LyC escape has been studied theoretically by Verhamme et al. (2015) and empirically by Verhamme et al. (2017). The 2015 work highlighted the presence of Ly$\alpha$ emission at line center as a smoking gun for LyC escape. We see a nonzero flux at $(1 + v_{\text{sys}}/c) \times \lambda_{\text{Ly}\alpha,0}$ for all knots, especially for knots A and B, where the red peak is at $\sim+150 \text{ km s}^{-1}$ and hence consistent with the instrumental broadening of FWHM $150 \text{ km s}^{-1}$, implying that the detection of Ly$\alpha$ at systemic velocity is not significant. Moreover, the clumpy model employed by Verhamme et al. (2015) had zero neutral hydrogen density in the interclump medium, which is not realistic in view of ionized nebulae having neutral hydrogen fractions of $10^{-5}$ to $\gtrsim 10^{-4}$ (Osterbrock & Ferland 2006; Gray et al. 2019), enough to make Ly$\alpha$ optically thick. In Verhamme et al. (2017) and Izotov et al. (2018a, 2018b), the Ly$\alpha$ properties of confirmed LyC leakers were studied, finding a positive correlation between Ly$\alpha$ and LyC escape and a negative correlation between LyC escape and Ly$\alpha$ peak separation. Peak separations of $\sim 400 \text{ km s}^{-1}$, akin to knot B (and possibly C), are found in some confirmed LyC leakers with escape fractions of $\lesssim 10\%$ (Izotov et al. 2018b). Based on the Ly$\alpha$ escape fractions, line widths, and peak separation, it appears from these empirical relations that knot A is an unlikely source of LyC. On the other hand, the radiative transport of Ly$\alpha$ and LyC is very different: Ly$\alpha$ bounces on the neutral clouds but is still optically thick in the ionized medium, and Jaskot et al. (2019) find no correlation between $f_{\text{esc}}^{Ly\alpha}$ and [O III]/[O II]. However, they find that the Ly$\alpha$ peak separations correlate with [O III]/[O II] and that $f_{\text{esc}}^{Ly\alpha}$ correlates with the LIS gas-covering fraction, and suggest the LIS lines trace the fraction of low-column-density channels and that the ionization level is related to the transparency of those. In this context, this could mean that knot C has many free channels but that they are relatively opaque. On the other hand, the inferred Si IV columns are not higher in knot C, which also has markedly lower (at least locally) electron densities (James et al. 2013; Menacho et al. 2021). The ISM absorption lines exclusively measure the gas properties along the line of sight (which is what matters for LyC escape), and the ionization level has limited power in deducing the line-of-sight LyC optical depth because it is not affected by H I and may therefore have greater diagnostic power on the sample level than for individual galaxies/star-forming regions.

Kakiichi & Gronke (2021) also study the relation between LyC escape and Ly$\alpha$ peak separation based on radiation hydrodynamics simulations and in addition also the Ly$\alpha$ red peak asymmetry (see their Figure 13, and Table 1). The properties of knots A and B are consistent with those expected for the radiation-bounded regime and low LyC escape. Knot C
has a markedly higher red peak asymmetry, which is consistent with the regime where LyC escapes through holes.

Chisholm et al. (2016) find that metal absorption lines of similar strengths but different ionization states (among them Si II 1260 and Si IV 1393 used here) show similar velocity and line width, suggesting they trace the same structures (see also Henry et al. 2015). We arrive at a similar conclusion, although we see a somewhat wider velocity distribution for the ionized gas than for the neutral component. Chisholm et al. (2018) showed that the Si II covering fractions are indeed correlated with the H I covering derived from the Lyman series (excluding Ly$\alpha$) absorption and that this correlates with the observed LyC escape fraction. Using the relation (which has a considerable scatter) between Si II and H I covering in Chisholm et al. (2018), we expect knots A and B to have H I covering fractions of unity, and for knot C, $f_{c,H1} \sim 0.8$, and thus a net LyC escape.

While all of these relations have considerable scatter, knot C stands out as the most likely source of LyC.

Returning to the shape of the absorption profiles and the choice of reference velocity in Figure 7, where in the right panel all profiles are compared relative to the $H\alpha$ velocity of knot C. Knots B and C differ by only ~30 km s⁻¹ in systemic $H\alpha$ velocity (and less, considering stellar velocities), and changing to individual reference velocities have little effect on the comparison. Knot A, however, differs by ~110 km s⁻¹ from knot C, and in the comparison in the right panel of Figure 7, the absorption profile of knot A is significantly blueshifted with respect to knots B and C. The similarity of the absorption profiles (and also $N_{Si II}$ and $f_c$) in the right panel of Figure 7 (in particular for A and B, agreeing both in shape and magnitude) is quite striking and suggest that they may be reflecting the same global outflow structure, which suggests that the absorption is perhaps not dominated by matter in the knots themselves, but farther out, along the line of sight toward us; the projected distance between knots A and B is just 1.4 kpc. The absorption profile of knot C is significantly shallower but covers the same range of velocity. Haro 11 is a merger (Östlin et al. 2015) and likely the division is such that knots A and B originate in the same late-type galaxy (with very active star formation, many star clusters, and dust lanes) while knot C is from the other, less gas-rich, galaxy participating in the merger. This division is also supported by kinematics as the emission-line profile halfway between knot C and A are markedly double peaked with two components separated by ~100 km s⁻¹ (Östlin et al. 2015).

The ionized halo of Haro 11 as seen in the MUSE $H\alpha$ image is quite round and extends to a radius of 13 kpc (Menacho et al. 2019). If we assume that the radius along the line of sight is similar and that this also equals the extent of the neutral outflows, which is observed to be $\gtrsim 500$ km s⁻¹, we find that on average each 17 km s⁻¹ velocity bin would correspond to ~400 pc. The individual neutral clouds probably have a radius on the order of $\lesssim 1$ pc (McCourt et al. 2018), in which case the area of the aperture (taken to have a diameter of $0.3''$) is $\sim 10^4$ times the cross section of an individual cloud. The probability of a photon hitting a single cloud is then $10^{-4}$, and for 30,000 randomly distributed clouds a covering fraction of $f_c = 0.95$ would be expected, while 7000 clouds would yield $f_c = 0.5$ and 1000 clouds $f_c = 0.1$. The resulting volume-filling fractions of neutral clouds per velocity bin would then be ~0.005 (for $f_c = 0.95$). These are of course very crude estimates: the velocity versus radius scaling is not likely to be linear, and the

cloud size is likely smaller at small radii but nevertheless illustrates that high covering fractions may be found for low volume-filling fractions.

## 5. Conclusions

We have presented new ultraviolet HST/COS spectroscopy of knots A and B in Haro 11 and investigated the $Ly\alpha$ line profiles and absorption lines from silicon arising in the neutral (Si II) and ionized (Si IV) ISM. We combine these with available HST/COS spectroscopy for knot C to make a comparative analysis of which knot, as inferred from the neutral ISM covering, might be the most likely source of the LyC escape detected in this galaxy. We have included optical data from the ESO/VLT/MUSE integral field spectrometer to derive the nebular extinction (from $H\alpha/H\beta$), ionization level, and $Ly\alpha$ escape fractions for knots A, B, and C.

All three knots show prominent $Ly\alpha$ line emission detected at high signal to noise. Knots A and B have double peaks in $Ly\alpha$ while knot C only shows a hint of a blue peak (or more likely a P Cyg profile). Broad H I absorption near $Ly\alpha$ is seen in knots A and B, but not in knot C. The Si II and Si IV lines for all knots reveal outflowing material with blueshifted velocities up to $\gtrsim 500$ km s⁻¹. From the Si II and Si IV lines, we have derived column densities and covering fractions for the intervening neutral and ionized gas, respectively, and as a function of velocity. Below we summarize our conclusions:

1. Knot A has the smallest nebular reddening and a high ionization, but yet has a $Ly\alpha$ escape fraction of only ~1%, significantly lower than its reddening suggests, indicating that $Ly\alpha$ has been strongly attenuated by dust through multiple resonant scatterings. The $Ly\alpha$ line is double peaked, but with a line width and peak separation that indicate a rather high H I column density. Broad absorption around $Ly\alpha$, analyzed by Voigt profile fitting rendered a H I column of $log(N_{H1}) = 20.7$ (cm⁻²). The ISM absorption lines reveal covering fractions of neutral gas close to unity over $\gtrsim 300$ km s⁻¹. While the high ionization in and around knot A led Keenan et al. (2017) to suggest it as the most likely source of leaking LyC radiation, the high neutral covering implies that significant LyC escape along the line of sight is unlikely.

2. Knot B has the highest observed $H\alpha$ flux and equivalent width, but also the highest reddening, hence this is where most LyC and $Ly\alpha$ photons are produced, but the $Ly\alpha$ escape fraction is <1%. The $Ly\alpha$ line width is similar to knot A, but the peak separation is smaller, with values comparable to some confirmed LyC leaking galaxies, but the Voigt profile fitting returns a H I column, $log(N_{H1}) = 21$, higher than for knot A. Moreover, the ISM absorption lines reveal a covering fraction of neutral gas very similar to knot A. Also, the UV continuum slope ($\beta$) is quite red, indicating even higher reddening than estimated from $H\alpha/H\beta$. In conclusion, significant LyC escape along the line of sight from knot B is quite unlikely.

3. Knot C has the highest $Ly\alpha$ flux and has an order-of-magnitude higher $Ly\alpha$ escape fraction than knots A and B, despite a high nebular reddening (similar to that of knot B). Curiously, the UV continuum slope is very blue, indicating a significantly lower reddening than the nebular estimate. In fact, a lot more $Ly\alpha$ photons escape knot C than the nebular reddening would suggest, indicating that $Ly\alpha$ escapes through less dusty paths than $H\alpha$. The $Ly\alpha$ profile is narrow and with no sign of broad absorption, all suggesting low H I column density. The Si II lines suggest a low ~50% covering fraction
of neutral gas, suggesting a significant transmission (2%–35%) of LyC along the line of sight. The O I 1302 Å line is however shallower than the Si II lines, suggesting that the former may be formed also in the phase where hydrogen is ionized and that the neutral covering may be even smaller, around 20%. Based on the ionization budget, knot C would need to have an LyC escape fraction of 30% to account for all the observed LyC flux detected in the FUSE observations, and our results suggest this is plausible.

In conclusion, the observations here presented and analyzed favor knot C as the most likely of the three knots to significantly contribute to the observed LyC leakage from Haro 11.

This research has made use of the NASA/IPAC Extra-
galactic Database (NED), which is operated by the Jet
Propulsion Laboratory, California Institute of Technology,
under contract with the National Aeronautics and Space
Administration. This work was supported by the Swedish
Research Council (Vetenskapsrådet). M.S.O. acknowledges
support from NASA, HST-GO-15352.002-A. D.K. is funded
by the Centre National d'Études Spatiales (CNES)/Centre
National de la Recherche Scientifique (CNRS); Convention No
131425. J.M.M.H. is funded by Spanish State Research
Agency grants PID2019-107061GB-C61 and MDM-2017-
0737 (Unidad de Excelencia María de Maeztu CAB). M.G.
was supported by NASA through the Hubble Fellowship grant
HST-HF2-51409. The Cosmic Dawn Center is funded by the
Danish National Research Foundation under grant No. 140. We
thank the anonymous referee for a thorough and constructive
review.

Facilities: HST (ACS, COS), ESO (VLT/MUSE).

ORCID iDs

Göran Östlin https://orcid.org/0000-0002-3005-1349

T. Emil Rivera-Thorsen https://orcid.org/0000-0002-9204-3256

Matthew Hayes https://orcid.org/0000-0001-8587-218X

Axel Runnholm https://orcid.org/0000-0002-1025-7569

Genoveva Micheva https://orcid.org/0000-0003-4376-2841

M. S. Oey https://orcid.org/0000-0002-5808-1320

Angela Adamo https://orcid.org/0000-0002-8192-8091

Arjan Bik https://orcid.org/0000-0001-8068-0891

John M. Cannon https://orcid.org/0000-0002-1821-7019

Max Gronke https://orcid.org/0000-0003-2491-060X

Daniel Kunth https://orcid.org/0000-0002-3635-0112

Peter Laursen https://orcid.org/0000-0003-4207-0245

J. Miguel Mas-Hesse https://orcid.org/0000-0002-8823-9723

Jens Melinder https://orcid.org/0000-0003-0470-8754

Matteo Messa https://orcid.org/0000-0003-1427-2456

Mattia Sirressi https://orcid.org/0000-0002-0923-8352

Linda Smith https://orcid.org/0000-0002-0806-168X

References

Adamo, A., Östlin, G., Zackrisson, E., et al. 2010, MNRAS, 407, 870

Alexandroff, R. M, Heckman, T. M, Borthakur, S., et al. 2015, ApJ, 810, 104

Atek, H., Kunth, D., Schaerer, D., et al. 2014, A&A, 561, A89

Bacon, R., Accardo, L., Adjali, M., et al. 2010, Proc. SPIE, 7735, 773508

Basu-Zych, A. R., Lehmer, B., Fragos, T., et al. 2016, ApJ, 818, 140

Bergvall, N., & Östlin, G. 2002, A&A, 390, 891

Bergvall, N., Zackrisson, E., Anderson, B.-G., et al. 2006, A&A, 448, 513

Brocklehurst, M. 1971, MNRAS, 153, 471

Calzetti, D., Armus, L., Bohlin, R. C., et al. 2000, ApJ, 533, 682

Cappellari, M. 2017, MNRAS, 466, 798

Cardamone, C., Schawinski, K., Sarzi, M., et al. 2009, MNRAS, 399, 1191

Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245

Chisholm, J., Gazagnes, S., Schaerer, D., et al. 2018, A&A, 616, A30

Chisholm, J., Tremonti, C. A., Leitherer, C., Chen, Y., & Wofford, A. 2016, MNRAS, 457, 3133

Churchill, C. W., Evans, J. L., Stemock, B., et al. 2020, ApJ, 904, 28

Cormier, D., Madden, S. C., Leboutellier, V., et al. 2014, A&A, 564, 121

Dittenber, B., Oey, M. S., Hodges-Kluck, E., et al. 2020, ApJL, 890, L12

Duval, F., Schaerer, D., Östlin, G., & Laursen, P. 2014, A&A, 562, 52

Fitzpatrick, E. L. 1999, PASP, 111, 63

Fox, A. 2018, Instrument Science Rep. COS 2018-07

Fox, A. 2018, Instrument Science Rep. COS 2018-08

Garnett, D. R., Dufour, R. J., & Peimbert, M. 1995, ApJL, 449, L77

Giavalisco, M., Koratkar, A., & Calzetti, D. 1996, ApJ, 466, 831

Goudfrooij, P., Burgh, E., Aloisi, A., Hartig, G., & Penton, S. 2012, Instrument
Science Rep. COS 2010-10(v2)

Grady, C. A., Bjorkman, K. S., & Snow, T. P. 1987, ApJ, 320, 376

Gray, W. J., Scannapieco, E., & Lehnert, M. 2019, ApJ, 875, 110

Green, J. C., Froning, C. S., & Osterman, S. 2012, ApJ, 744, 60

Grimes, J. P., Heckman, T., Strickland, D., et al. 2007, ApJ, 688, 891

Gronke, M., Dijkstra, M., McCourt, M., & Oh, S. P. 2017, A&A, 607, A71

Guseva, N. G., Izotov, Y. I., Fricke, K. J., et al. 2012, A&A, 541, 115

Hayes, M. 2019, SAAS, 46, 319

Hayes, M., Östlin, G., Atek, H., et al. 2007, MNRAS, 382, 1465

Hayes, M., Schaerer, D., Östlin, G., et al. 2011, ApJ, 730, 8

Heckman, T. M., Alexandroff, R. M., Borthakur, S., et al. 2015, ApJ, 809, 147

Henry, A., Scarlata, C., Martin, C. L., & Erb, D. 2015, ApJ, 809, 19

Hirschauer, et al. 2021, Cosmic Origins Spectrograph Instrument Handbook, v13.0 (Baltimore, MD: STScI)

Howarth, I. D., & Prinja, R. K. 1989, ApJS, 69, 527

Inoue, A. K., & Iwata, I. 2008, MNRAS, 387, 1681

Iwata, I., Inoue, A. K., Matsuda, Y., et al. 2009, ApJ, 692, 1287

Izotov, Y. I, Schaerer, D., Worseck, G., et al. 2018a, MNRAS, 478, 4851

Izotov, Y. I., Worseck, G., Schaerer, D., et al. 2018b, MNRAS, 474, 4514

James, B. L., Tsamis, Y. G., Walsh, J. R., Barlow, M. J., & Westmoquette, M. S. 2013, MNRAS, 430, 2097

Jaskot, A. E., Dowd, T., Oey, M. S., Scarlata, C., & McKinney, J. 2019, ApJ,
885, 96

Jaskot, A. E., & Oey, M. S. 2013, ApJ, 766, 91

Jones, T. A., Ellis, R. S., Schenker, M. A., et al. 2013, ApJ, 779, 52

Kakiichi, K., & Gronke, M. 2021, ApJ, 908, 30

Keenan, R. P., Oey, M. S., & Jaskot, A. E. 2017, ApJ, 848, 12

Laursen, P., Duval, F., & Östlin, G. 2013, ApJ, 766, 124

Leitet, E., Bergvall, N., Piskunov, N., & Andersson, B.-G. 2011, A&A,
532, 107

Eitherer, C., Hernandez, S., Lee, J., & Oey, M. S. 2016, ApJ, 823, 64

Leitherer, C., Schaerer, D., Goldader, J. D., et al. 1999, ApJS, 123, 3

MacHattie, J. A., Irwin, J. A., Madden, S. C., Cormier, D., & Remy-Ruyer, A. 2014, MNRAS, 438, L66

McCandliss, S. R., France, K., Osterman, S., Green, J. C., MacPhate, J. B., & Wilkinson, E. 2010, ApJL, 709, L813

McCourt, M., Peng, O. S., O'Leary, R., & Madigan, A.-M. 2018, MNRAS,
473, 5704

Menacho, V., Östlin, G., Bik, A., et al. 2019, MNRAS, 487, 3183

Menacho, V., Östlin, G., Bik, A., et al. 2021, MNRAS, submitted

Micheva, G., Iwata, I., Inoue, A. K., et al. 2017, MNRAS, 465, 316

Micheva, G., Östlin, G., Melinder, J., et al. 2020, ApJ, 903, 123

Mostardi, R. E., Shapley, A. E., Nestor, D. B., et al. 2013, ApJ, 779, 65

Mostardi, R. E, Shapley, A. E., Steidel, C. C., et al. 2015, ApJ, 810, 107

Neufeld, D. 1991, ApJL, 375, L85

Osterbrock, D. E., & Ferland, G. J. 2006, Astrophysics of Gaseous Nebulae and Active Galactic Nuclei (2nd ed.; Sausalito, CA: Univ. Science Books)

Östlin, G., Amram, P., Bergvall, N., et al. 2001, A&A, 374, 800

Östlin, G., Hayes, M., Duval, F., et al. 2014, ApJ, 797, 11

Östlin, G., Hayes, M., Kunth, D., et al. 2009, AJ, 138, 923

Östlin, G., Marquart, T., Cumming, R. J., et al. 2015, A&A, 583, 55

Pardy, S. A., Cannon, J. M., Östlin, G., et al. 2016, AJ, 152, 178

Pellegrini, E. W., Oey, M. S., Winkler, P. F., et al. 2012, ApJ, 755, 40

Petitjean, P. 1998, arXiv:astro-ph/9810418

Prestwich, A. H., Jackson, F., Kaaret, P., et al. 2015, ApJ, 812, 166

Prevot, M. L., Lequeux, J., Maurice, E., Prevot, L., & Rocca-Volmerange, B. 1984, A&A, 132, 389

Puschnig, J., Hayes, M., Östlin, G., et al. 2017, MNRAS, 469, 3252
Rivera-Thorsen, T. E., Dahle, H., Chisholm, J., et al. 2019, *Sci*, **366**, 738

Rivera-Thorsen, T. E., Hayes, M., Östlin, G., et al. 2015, *ApJ*, **805**, 14

Rivera-Thorsen, T. E., Östlin, G., Hayes, M., et al. 2017, *ApJ*, **837**, 27

Sánchez-Blázquez, P., Peletier, R. F., Jiménez-Vicente, J., et al. 2006, *MNRAS*, **371**, 703

Savage, B. D., & Sembach, K. R. 1991, *ApJ*, **371**, 245

Scarlata, C., Colbert, J., Teplitz, H. I., et al. 2009, *ApJL*, **704**, L198

Schlafly, E. F., & Finkbeiner, D. P. 2011, *ApJ*, **737**, 103

Shepard, K, Gies, D. R., Lester, K. V., et al. 2020, *ApJ*, **888**, 82

Siana, B., Shapley, A. E., Kulas, K. R., et al. 2015, *ApJ*, **804**, 17

Vader, P., Frogel, J. A., Tendrup, D. M., Heisler, C. A., et al. 1993, *AJ*, **106**, 1743

Vanzella, E., Guo, Y., Giavalisco, M. tt, et al. 2012, *ApJ*, **751**, 70

Vazdekis, A., Sánchez-Blázquez, P., Falcón-Barroso, J., et al. 2010, *MNRAS*, **404**, 1639

Verhamme, A., Orlitova, I., Schaerer, D., et al. 2017, *A&A*, **597**, A13

Verhamme, A., Orlitova, I., Schaerer, D., & Hayes, M 2015, *A&A*, **578**, A7

Verhamme, A., Schaerer, D., & Maselli, A. 2006, *A&A*, **460**, 397