# The Lyα emission from high-z galaxies hosting strong damped Lyα systems

Ravi Joshi, Raghunathan Srianand, Pasquier Noterdaeme, Patrick Petitjean

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Ravi Joshi, Raghunathan Srianand, Pasquier Noterdaeme, Patrick Petitjean. The Lyα emission from high-z galaxies hosting strong damped Lyα systems. Monthly Notices of the Royal Astronomical Society, 2017, 465 (1), pp.701-721. 10.1093/mnras/stw2733. hal-01582733

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# The Lyα emission from high-z galaxies hosting strong damped Lyα systems

Ravi Joshi,<fnref n="1" />★ Raghunathan Srianand,<fnref n="1" /> Pasquier Noterdaeme<fnref n="2" /> and Patrick Petitjean<fnref n="2" />

Accepted 2016 October 20. Received 2016 October 20; in original form 2016 June 27

## ABSTRACT

We study the average Lyα emission associated with high-z strong (log N(H1) ≥ 21) damped Lyα systems (DLAs). We report Lyα luminosities ($L_{Lyα}$) for the full as well as various sub-samples based on N(H1), z, (r − i) colours of QSOs and rest equivalent width of Si ii λ 1526 line (i.e. W1526). For the full sample, we find $L_{Lyα} < 10^{41}$ (3σ) erg s⁻<fnref n="1" /> with a 2.8σ level detection of Lyα emission in the red part of the DLA trough. The $L_{Lyα}$ is found to be higher for systems with higher W1526 with its peak, detected at ≥3σ, redshifted by about 300–400 km s⁻<fnref n="1" /> with respect to the systemic absorption redshift, as seen in Lyman break galaxies (LBGs) and Lyα emitters. A clear signature of a double-hump Lyα profile is seen when we consider W1526 ≥0.4 Å and (r − i) < 0.05. Based on the known correlation between metallicity and W1526, we interpret our results in terms of star formation rate (SFR) being higher in high-metallicity (mass) galaxies with high velocity fields that facilitates easy Lyα escape. The measured Lyα surface brightness requires local ionizing radiation that is 4–10 times stronger than the metagalactic UV background at these redshifts. The relationship between the SFR and surface mass density of atomic gas seen in DLAs is similar to that of local dwarf and metal-poor galaxies. We show that the low-luminosity galaxies will contribute appreciably to the stacked spectrum if the size–luminosity relation seen for H1 at low z is also present at high-z. Alternatively, large Lyα haloes seen around LBGs could also explain our measurements.

**Key words:** galaxies: high-redshift – galaxies: ISM – quasars: absorption lines – galaxies: star formation.

## 1 INTRODUCTION

Damped Lyα systems (DLAs) are the highest H1 column density absorbers seen in QSO spectra, with N(H1) ≥ 2 × 10²⁰ cm⁻<fnref n="2" />. These absorbers trace the bulk of the neutral hydrogen at 2 ≤ z ≤ 3 (Prochaska & Wolfe 2009; Noterdaeme et al. 2009, 2012b) and have long been considered to arise from the high-redshift precursors of present-day galaxies (for a review see, Wolfe, Gawiser & Prochaska 2005). Presence of enriched elements (Pettini et al. 1994), measured excitation of C ii fine-structure levels (Wolfe, Prochaska & Gawiser 2003a; Srianand et al. 2005), existence of a correlation between metallicity and the velocity spread of low-ion absorption lines (Ledoux et al. 2006) akin to the mass–metallicity relation seen in galaxies and rotational excitation of high J level of H₂ detected in a small fraction of DLAs (Ge & Bechtold 1997; Ledoux, Petitjean & Srianand 2003; Noterdaeme et al. 2008) etc. suggest that DLAs are associated in some way with star-forming regions. Even if one associates a moderate star formation rate (SFR) to DLAs, they will

contribute appreciably to the global inventory of the cosmic SFR density at high-z (Wolfe, Gawiser & Prochaska 2003b; Srianand et al. 2005).

A straightforward way to establish the link between the H1 gas and stellar components in DLAs is to detect the DLA host galaxies in line or continuum emission. A wide variety of galaxies hosting DLAs are found in imaging studies at low z (i.e. z ≤ 1; Chen & Lanzetta 2003; Rao et al. 2003). However, till date only a handful of detections of galaxy counterparts hosting DLAs at z > 2 are confirmed using spectroscopy (Möller, Fynbo & Fall 2004; Fynbo et al. 2010; Péroux et al. 2011; Bouché et al. 2012; Krogager et al. 2012, 2013; Noterdaeme et al. 2012a; Jorgenson & Wolfe 2014; Kashikawa et al. 2014; Hartoog et al. 2015; Rubin et al. 2015; Srianand et al. 2016). Based on all these efforts, the DLAs with emission line detections are found to trace a galaxy population with SFR of 0.4–25 M☉ yr⁻<fnref n="1" /> and impact parameters ranging from 0.4 to 182 kpc, with a mean value of ~25 kpc (Péroux et al. 2011; Krogager et al. 2012). Christensen et al. (2014) measured stellar masses of these DLA host galaxies. These are found to be consistent with the expected values based on mass–metallicity relations of high-z galaxies. Note, the number of such DLAs are

<footnote n="1">¹Inter-University Centre for Astronomy and Astrophysics, Post Bag 4, Ganeshkhind, Pune 411007, India</footnote>

<footnote n="2">²UPMC-CNRS, UMR7095, Institut d'Astrophysique de Paris, F-75014 Paris, France</footnote>

<footnote>* E-mail: rvjoshirv@gmail.com</footnote>
much smaller than the known number of DLAs or galaxies with spectroscopic redshifts. This is mainly because it is found to be challenging as the glare of the bright background quasar makes it difficult to detect the faint galaxy producing the DLA absorption at small impact parameters (Lowenthal et al. 1995; Bunker et al. 1999; Kulkarni et al. 2000, and several unreported attempts).

Issues related to detecting faint galaxies against the glare of a bright quasar can be easily addressed if one can find sightlines with two or more optically thick H₁ absorbers (Steidel & Hamilton 1992; O'Meara, Chen & Kaplan 2006; Christensen et al. 2009; Fumagalli et al. 2010). In these cases, the continuum emission from the foreground DLA can be detected at wavelengths less than 912 Å in the rest frame of the higher redshift optically thick H₁ absorber that removes the quasar light completely. Using this so-called Double-DLA technique, Fumagalli et al. (2015) have placed stringent constraints on the *in situ* star formation rates, SFR <0.09–0.27 M$_{\odot}$ yr$^{-1}$ at the position of the absorbing gas. By looking at the galaxies detected around the QSO sightlines, they concluded that most of the DLAs do not originate from highly star-forming galaxies that are coincident with the absorbers.

The effective optical radius of high-z galaxies are small (i.e. ~1 kpc, Shibuya, Ouchi & Harikane 2015). If we conjecture that the DLAs are associated with such galaxies, then the observed number of DLAs per unit redshift requires H₁ gas to be extended appreciably beyond the stellar regions. Image stacking analysis of high-z galaxies suggests that the stellar distribution is also extended (Hathi et al. 2008; Rafelski, Wolfe & Chen 2011). By comparing the covering factor of diffuse emission from Lyman break galaxies (LBGs) and H₁ gas in DLAs (with an assumed star formation efficiency), Rafelski et al. (2011) argued that the star formation in the outskirts of galaxies (i.e. ~6 kpc from the LBG's core) may be dominated by atomic gas probed by DLAs with log N(H₁) ≥ 21 but with at least a factor of 10 less efficiency compared to that seen in local stellar discs. If the framework presented by Rafelski et al. (2011) is true then in a DLA sample of log N(H₁) ≥ 21 one should see evidence for low surface star formation along the line of sight and a faint galaxy within ~6 kpc. Similarly narrow-band image stacking analysis of LBGs at the rest-frame Lyα wavelength has shown diffuse Lyα emission extending up to several tens of kpcs (Steidel et al. 2011). In addition, Rauch et al. (2008) have detected a population of faint Lyα emitters (LAEs) with spatially extended Lyα emission that have a total cross-section consistent with that of DLAs. While the Lyα scattering requires only N(H₁) ≥ 2 × 10¹⁷ cm⁻² to generate large diffuse Lyα emission around LBGs, associating these scattering regions in the LBG's outskirts to DLAs will also mean that the low impact parameter regions [that will contribute to the high N(H₁) systems (see, Krogager et al. 2012)] can have associated Lyα emission even if there is no *in situ* star formation being present. All this suggests that the spectral stacking analysis will be very useful in probing DLA-galaxy connection.

Interestingly, in spectroscopic surveys using fibres (e.g. 3 and 2 arcsec diameter fibres employed in SDSS-DR7 and SDSS-DR12 (BOSS), respectively) one integrates, in addition to the light from a distant quasar, the light from all the foreground galaxies that happen to fall within the fibre along our line of sight. This allows detection of nebular emission lines from such galaxies on top of the QSO spectrum (Wild, Hewett & Pettini 2007; Noterdaeme, Srianand & Mohan 2010; Ménard et al. 2011; Straka et al. 2015). In this regard, the Lyα emission holds great potential in determining the nature of high-z DLA hosts. Since the direct detections of Lyα emission from high-z galaxies hosting DLAs are very rare, several studies have attempted to detect the Lyα emission in the composite spectra

(Rahmani et al. 2010; Noterdaeme et al. 2014). For instance, in a stacking analysis of 341 DLAs of mean redshift z ~2.86, and log N(H₁) ≥ 20.62 seen in the SDSS quasar spectra, Rahmani et al. (2010) have found a 3σ upper limit on SFR of ≤1.2 M$_{\odot}$ yr$^{-1}$. Using ~95 extremely strong DLAs (ESDLAs) with log N(H₁) ≥ 21.7 in the SDSS-DR12 spectra, Noterdaeme et al. (2014) have detected the Lyα emission with an average luminosity of 6 × 10$^{41}$ erg s$^{-1}$ in the DLA core. These contemporary studies motivate us to unveil the average SFR of DLAs in the stacked spectra by exploiting the unprecedented number of DLAs found in the recent SDSS-BOSS survey (Noterdaeme et al. 2009, 2012b). Note that, this compilation contains ~14 972 DLAs with log N(H₁) ≥ 20.3 which is a factor of 2.2 larger than the DLA compilation based on SDSS-DR9 by Noterdaeme et al. (2012b). The available information from emission and absorption stacks will not only allow us to detect the Lyα emission but also to probe its dependence on various quantities, such as H₁ column density, rest equivalent width of metal absorption lines, QSO colours and absorption redshift.

This paper is organized as follows. Section 2 describes our sample selection criteria. In Section 3, we present details of our analysis and average Lyα emission for our full sample and various sub-samples. Section 4 presents discussion based on absorption lines detected in different stacked spectra. In Section 5, we discuss our results in the framework of different possible scenarios. A detailed summary of our study is presented in Section 6. Throughout, we have assumed the flat cosmology with $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_m = 0.3$ and $\Omega_\Lambda = 0.7$.

## 2 SAMPLE

For our analysis, we have constructed a sample of systems from the BOSS-DR12 DLA catalogue, based on automatic search for DLAs in the quasar spectra (see, Noterdaeme et al. 2012b) in SDSS-III BOSS Data Release.<fnref n="1" /> We use the following stringent criteria. (1) We consider only systems with log N(H₁) ≥ 21.0. In this case, the Lyα absorption has a dark core (with $\tau \ge 10$) spread at least over 7 times the average full width at half-maximum, FWHM (~160 km s$^{-1}$) of the instrumental profile of the BOSS spectrograph. (2) We consider DLAs detected in spectra with a median continuum-to-noise ratio (CNR) ≥5. This cut is used to ensure an accurate determination of the H₁ column density. The CNR is defined over a redshift range with the minimum redshift set redward of any possible Lyman break present in the spectrum and the maximum redshift defined at 5000 km s$^{-1}$ blueward of QSO emission redshift (see also Noterdaeme et al. 2009). (3) We consider only DLAs which are in the redshift range 2.3 ≤ z$_{abs}$ ≤ 3.44. This redshift range is chosen to avoid the poor signal-to-noise regions in the blue spectrum and to exclude wavelength ranges affected by the residuals from the poorly subtracted sky emission lines in the red. (4) We restrict ourselves to the observed wavelength range redward of the Ly β + O VI emission line of the quasar to avoid any broad associated O VI absorption being misidentified as a DLA. (5) We avoid DLAs located in the vicinity of the quasar (i.e. the so-called proximate DLAs) by considering the systems with velocity offset of ≥5000 km s$^{-1}$ with respect to the quasar emission redshift. After applying these conditions and avoiding sightlines having broad absorption lines from QSO outflows, we have a sample of 830 DLAs satisfying the above criterion.

<footnote n="1">¹ https://www.sdss3.org/</footnote>
<table><caption>**Table 1.** Details of our DLA sample.</caption><thead><tr><th>Sample</th><th>Criteria</th><th>DLAs</th></tr></thead><tbody><tr><td>Total DLAs</td><td>log N(H1) ≥ 20.3</td><td>14 515</td></tr><tr><td>Primary sample</td><td>2.3 ≤ z<sub>abs</sub> ≤ 3.44<br>CNR ≥ 5<br>log N(H1) ≥ 21.0</td><td>830</td></tr><tr><td></td><td>β ≥ 5000 km s<sup>-1</sup><br>non-BAL<br>z<sub>abs</sub> > z(Ly β)</td><td></td></tr><tr><td>New measurements</td><td>log N(H1) < 21.0</td><td>62</td></tr><tr><td></td><td>Noisy spectra</td><td>8</td></tr><tr><td></td><td>Uncertain z</td><td>52</td></tr><tr><td></td><td>Uncertain N(H1)</td><td>4</td></tr><tr><td>Final sample</td><td></td><td>704</td></tr></tbody></table>

To avoid any uncertainty in the absorption redshift and/or H1 column density measurements, we further visually inspected the entire sample and refitted the DLA absorption with a Voigt profile. As accurate redshift measurement is critical for detecting the signal in the co-added spectrum, we measure the absorption redshift by cross-correlating the low-ionization metal absorption lines (e.g. C II λ1334, Si II λ1526, Al II λ1670, Fe II λλ1608, 2344, 2374, 2382, 2586, 2600, Mg II λλ2796, 2803) redward to Lyα emission. When no metal absorption line is detected, we keep the redshift obtained from the H1 Lyman series transitions. This resulted in a final sample of 704 DLAs with good z<sub>abs</sub> and N(H1) measurements. Among them, ~80 per cent of the DLAs have redshift determined from metal lines. In the case of 62 DLAs, the re-measured N(H1) turns out to be less than our cutoff value 10<sup>21</sup> cm<sup>-2</sup>, thereby they go out of our primary sample. Similarly for 64 cases visual inspections, independently by two of the co-authors of this paper, suggest that the DLA trough is noisy and lack of higher H1 Lyman series lines and associated metal lines prevent accurate measurements of N(H1) and z. These 64 systems are also removed from our primary sample. Details of our sample are summarized in Table 1.

While most of the results presented here are based on this sample, sometimes to test our results based on smaller sub-samples, in the appendix we also present results based on samples where we have released the limiting CNR that defines our sample. The reduced one-dimensional spectra for our entire sample of 704 quasars were downloaded from the SDSS-BOSS Data Archive Server.<fnref n="2" /> Details about the BOSS spectral information can be found in Dawson et al. (2013). Briefly, the BOSS spectra were obtained with a fibre having a diameter of 2 arcsec (a maximum impact parameter of ~8 kpc for the above assumed cosmology) and cover a spectral range from 3650 to 10 400 Å at a resolution (λ/Δλ) of about 1560–2270 in the blue (i.e. λ ≤ 6350 Å) and 1850–2659 in the red channels.

# 3 ANALYSIS

## 3.1 The composite spectra

We have generated several composite spectra with an aim to study both the average emission and absorption properties of high column density DLAs. For stacking, an individual spectrum is shifted to the rest frame of the DLA, while conserving the flux and rebinning on to a common grid, keeping the same pixel size (constant in velocity

space) as the original data (Bolton et al. 2012). We produce the median, continuum normalized stacked spectrum for studying the absorption lines of DLAs. Each individual spectrum is normalized using the best-fitting principal component analysis (PCA) continuum model of Bolton et al. (2012). We apply a median-smoothing filter to remove additional imperfections on intermediate scales in the PCA continuum normalized spectra. For this, we mask out pixels possibly containing narrow absorption lines, by only keeping fluctuations within 1.5σ of the continuum. We mainly consider regions redward of the Lyα emission, to avoid the blending by the Lyα forest. The normalized flux error in each pixel of the stacked spectrum is the absolute deviation measured around the mean flux. These spectra are used to study any possible correlation between average absorption properties and the Lyα emission.

For studying the Lyα emission, we convert each spectrum into luminosity per unit wavelength using the luminosity distance at the redshift of the DLAs for the above-mentioned cosmological parameters. To remove contribution from any possible outliers, we have used the weighted mean and the 3σ clipped weighted mean statistics with 1/σ<fnref n="2" /> weighting with σ being the error in the luminosity. However, we note that using 1/σ<fnref n="2" /> weighting will give more weightage to the low-redshift systems, when we use the rest-frame spectra in luminosity. Therefore, we have constructed the median composite spectrum also. As the DLAs are characterized by a wide flat absorption trough with zero transmitted flux (i.e. dark core), we do not re-scale or normalize the spectrum before the co-addition. The 1σ uncertainty over each pixel in the stacked spectrum is estimated from the central interval encompassing 68 per cent of the flux distribution of the corresponding pixel.

To check and correct the non-zero flux offsets (e.g. see also Rahmani et al. 2010; Cai et al. 2014; Noterdaeme et al. 2014) in the core of DLAs due to poor background subtraction, we follow two methods. (i) The average residual flux is calculated by stacking the Lyman limit region (i.e. λ ≤ 912 × (1 + z) Å) of DLAs in the observer’s frame, where the average flux is expected to be negligibly small due to the large expected H1 optical depth. For this, we have used 3259 QSO spectra having DLAs in the redshift range 2.95 ≤ z ≤ 4.6, whose Lyman limit falls in the observed wavelength range 3600–5100 Å. The stacked spectrum based on three different statistics discussed above is shown in three different panels in Fig. 1 with a blue solid histogram. It is clear that there is a non-zero residual flux present over the redshift range 2.0 ≤ z ≤ 3.2 (i.e. 3650 ≤ λ ≤ 5100 Å) for the Lyα absorption. This residual flux could arise from leakage of far-UV radiation from the quasar host galaxy through the DLA (Cai et al. 2014) or due to imperfect fibre sky subtraction (Bolton et al. 2012) and should be corrected for before we do the stacking. However, for DLAs in our sample, at 3.2 ≤ z ≤ 3.4, we do not have the better residual measurements from the Lyman limit systems. For this z range we use the second method. (ii) We construct the residual flux template (background) by stacking the DLA core pixels with τ ≥ 10 that falls within this observed wavelength range. To be specific, we have used the core pixels of all the available DLAs in addition to those that are part of our sample. We generated a template spectrum by considering only the DLA core pixels with optical depth τ ≥ 10 for a given N(H1) value. In Fig. 1, we also show the average flux in the stacked spectrum of the DLA core regions, constructed using ~7072 DLAs with log N(H1) ≥ 20.62 and average CNR ≥2 (grey curve). We construct the background for all the three stacking statistics and found that in the overlapping observed wavelength range, the residual flux in the Lyman limit region and core of the DLA absorption profile matches well. Besides, we note that the residual flux template does

<footnote n="2">² http://dr12.sdss3.org/bulkSpectra</footnote>
<figure><img src="image_2.png" /><figcaption>**Figure 1.** The average residual flux in the Lyman limit region of DLAs (blue solid curve) and in the DLA core pixels with $\tau \ge 10$ (grey curve) are shown as a function of the observed wavelength. The dashed line represents the running average of residual flux in the DLA core pixels, which is used as a background template for correcting the residuals in the DLA core. The shaded region in the top panel shows the wavelength range of interest for our stacking analysis.</figcaption></figure>

not depend on the brightness and colour of the background quasar. Next, in order to account for this average residual flux, we compute its running average (dashed curve) over 300 pixels and construct the background template for each statistics. We find that the background template obtained by varying the smoothing window in the range of 100–300 pixels gives similar results. Finally, before generating the stacked spectrum in the DLA rest frame, we correct each spectrum for the non-zero flux offset by subtracting the flux given by the template for the corresponding statistics.

In addition, to the non-zero average flux in the bottom, we also notice that in several DLAs there are large pixel-to-pixel variations in the core pixels (see few examples in Fig. A1 in appendix). While we could associate this with the residual from telluric line subtraction for few cases, the origin of such fluctuations in most cases is unclear. By and large these pixel-to-pixel variations do not affect our results, but may have strong influence when we are dealing with sub-samples having much smaller number of systems. As stated earlier, in such cases we relax the CNR constraint which defines our main sample to have additional systems to test the results we get from our main sample.

## 3.2 The average Lyα emission of DLAs

In Fig 2, we show our stacked spectrum of 704 DLAs obtained using the weighted mean (bottom panel), median (middle panel) and $3\sigma$ clipped weighted mean (top panel) statistics. Our stacked spectra show an asymmetric flux distribution in the bottom of the DLA absorption profile, with a possible enhanced flux in the red half (i.e. $\lambda_r > 1215.67$ Å) within the core region. Such asymmetric Lyα line profiles are naturally produced by the radiative transport of Lyα

<figure><img src="image_3.png" /><figcaption>**Figure 2.** The stacked spectrum corresponds to the $3\sigma$ clipped weighted mean (top), median (middle) and weighted mean (bottom), for DLAs in the redshift range of $2.30 \le z \le 3.44$. The blue segment shows the DLA core with $\tau \ge 10$ for log N(HI) = 21.0. The dashed curves show the synthetic profiles for lower (i.e. log N(HI) = 21) and median (i.e. log N(HI) = 21.23) column density of DLAs used to get the stacked spectrum.</figcaption></figure>

photons in an outflowing medium (Neufeld 1990, 1991; Dijkstra, Haiman & Spaans 2006; Hansen & Oh 2006; Verhamme, Schaerer & Maselli 2006; Dijkstra 2014).

It is not a straightforward exercise to quantify the Lyα line luminosity (or its upper limit), as the resonant scattering of Lyα photons along with the gas kinematics introduces a wide range of velocity offsets for the escaping Lyα photons. In addition, by stacking, we average over a large number of systems with varied Lyα optical depths and velocity fields. Therefore, we estimate three Lyα luminosities from each stacked spectrum: (i) integrated luminosity over the entire DLA core region, (ii) integrated luminosity ($L_{\lambda}^{b}$) in the blue part (i.e. core regions having $\lambda_r < 1215$ Å) and (iii) the same ($L_{\lambda}^{r}$) in the red part (i.e. $\lambda_r > 1215$ Å). We define the core region of the DLA absorption profile to be the wavelength range where the Lyα optical depth, $\tau \ge 10$, for the minimum column density of log [N(HI) /cm$^{-2}$] ~21 in our sample (see above in Section 2). This is indicated in each panel as a horizontal line. The integrated Lyα luminosity over the DLA core regions is summarized in Table 2. For the full sample, we do not detect any statistically significant signal above the $2\sigma$ level. All three statistics give nearly the same integrated luminosities and errors suggesting that the stacked spectrum is not severely influenced by a few outliers.

Furthermore, to account for the uncertainties related to under/oversubtraction of non-zero flux (i.e. background correction) and to test if the stack is influenced by few outliers, we perform a bootstrap analysis. For this, we make stacked spectra of 1000 sub-samples by randomly selecting 90 and 80 per cent of the sample. Note that as we have performed sub-sampling bootstrap with 90 per cent of the sample, the measured statistical error from
<table><caption>**Table 2.** The Lyα line luminosity measured from the stacked spectra.</caption><thead><tr><th rowspan="2">Criteria</th><th rowspan="2">Number of systems</th><th rowspan="2">log N(H<sub>I</sub>)</th><th rowspan="2">⟨z<sub>abs</sub>⟩</th><th colspan="3">Lyα luminosity (×10<sup>40</sup> erg s<sup>-1</sup>)</th></tr><tr><th>Weighted mean</th><th>3σ clipped weighted mean</th><th>Median</th></tr></thead><tbody><tr><td>Final sample</td><td>704</td><td>21.23</td><td>2.71</td><td>4.54 ± 3.12</td><td>4.70 ± 3.19</td><td>5.24 ± 3.31</td></tr><tr><td>N(H<sub>I</sub>) ≥ 10<sup>21.23</sup> cm<sup>-2</sup></td><td>354</td><td>21.40</td><td>2.72</td><td>8.83 ± 5.05</td><td>9.15 ± 5.13</td><td>6.75 ± 5.32</td></tr><tr><td>10<sup>21</sup> ≤ N(H<sub>I</sub>) &lt; 10<sup>21.23</sup> cm<sup>-2</sup></td><td>350</td><td>21.10</td><td>2.69</td><td>3.06 ± 4.32</td><td>3.27 ± 4.36</td><td>6.29 ± 4.65</td></tr><tr><td>z<sub>abs</sub> ≥ 2.7</td><td>352</td><td>21.23</td><td>2.97</td><td>6.21 ± 4.80</td><td>6.18 ± 4.78</td><td>7.96 ± 5.00</td></tr><tr><td>z<sub>abs</sub> &lt; 2.7</td><td>352</td><td>21.22</td><td>2.48</td><td>3.41 ± 4.12</td><td>4.60 ± 4.35</td><td>3.84 ± 4.39</td></tr><tr><td>W<sub>1526</sub> ≥ 0.8 Å<sup>a</sup></td><td>288</td><td>21.25</td><td>2.68</td><td>10.51 ± 4.95</td><td>13.62 ± 5.02</td><td>15.82 ± 5.18</td></tr><tr><td>W<sub>1526</sub> &lt; 0.8 Å</td><td>406</td><td>21.26</td><td>2.75</td><td>4.07 ± 4.67</td><td>4.63 ± 4.05</td><td>4.71 ± 5.02</td></tr><tr><td>(r - i) ≥ 0.1</td><td>352</td><td>21.24</td><td>2.82</td><td>6.22 ± 4.49</td><td>5.43 ± 4.56</td><td>10.93 ± 4.70</td></tr><tr><td>(r - i) &lt; 0.1</td><td>352</td><td>21.22</td><td>2.59</td><td>3.19 ± 4.34</td><td>5.66 ± 4.47</td><td>4.02 ± 4.66</td></tr></tbody></table>

Note. <sup>a</sup>Only systems with Si II λ1526 absorption line detection are considered.

<table><caption>**Table 3.** Luminosity of Lyα line in bootstrap analysis for both the median and 3σ clipped weighted mean stacked spectra.</caption><thead><tr><th rowspan="2">Sample</th><th colspan="4">Lyα luminosity (×10<sup>40</sup> erg s<sup>-1</sup>)</th><th colspan="3">Lyα luminosity (×10<sup>40</sup> erg s<sup>-1</sup>)</th></tr><tr><th>DLA bottom</th><th colspan="2">Median stack</th><th>DLA bottom</th><th colspan="2">3σ clipped weighted mean</th></tr><tr><th></th><th></th><th>L<sub>λ</sub><sup>b</sup></th><th>L<sub>λ</sub><sup>r</sup></th><th></th><th>L<sub>λ</sub><sup>b</sup></th><th>L<sub>λ</sub><sup>r</sup></th></tr></thead><tbody><tr><td>Full with 90 per cent bootstrap</td><td>5.82 ± 3.48<sup>a</sup></td><td>-1.06 ± 2.44</td><td>6.88 ± 2.49</td><td>3.87 ± 3.20</td><td>-2.27 ± 2.21</td><td>6.14 ± 2.32</td></tr><tr><td>Full with 80 per cent bootstrap</td><td>5.27 ± 3.69</td><td>-1.32 ± 2.58</td><td>6.58 ± 2.65</td><td>3.86 ± 3.40</td><td>-2.22 ± 2.35</td><td>6.08 ± 2.46</td></tr><tr><td>N(H<sub>I</sub>) ≥ 10<sup>21.23</sup> cm<sup>-2</sup></td><td>6.24 ± 4.95</td><td>-2.33 ± 3.46</td><td>8.57 ± 3.53</td><td>5.58 ± 4.62</td><td>-2.51 ± 3.20</td><td>8.09 ± 3.33</td></tr><tr><td>10<sup>21</sup> ≤ N(H<sub>I</sub>) &lt; 10<sup>21.23</sup> cm<sup>-2</sup></td><td>3.55 ± 4.82</td><td>-1.00 ± 3.37</td><td>4.56 ± 3.45</td><td>2.97 ± 4.45</td><td>-1.72 ± 3.06</td><td>4.69 ± 3.23</td></tr><tr><td>z<sub>abs</sub> ≥ 2.7</td><td>7.34 ± 5.28</td><td>1.72 ± 3.68</td><td>5.61 ± 3.79</td><td>6.03 ± 4.92</td><td>1.17 ± 3.39</td><td>4.86 ± 3.57</td></tr><tr><td>z<sub>abs</sub> &lt; 2.7</td><td>4.02 ± 4.60</td><td>-3.11 ± 3.22</td><td>7.13 ± 3.29</td><td>2.66 ± 4.27</td><td>-4.64 ± 2.96</td><td>7.30 ± 3.07</td></tr><tr><td>W<sub>1526</sub> ≥ 0.8 Å<sup>b</sup></td><td>14.55 ± 5.50</td><td>2.49 ± 3.87</td><td>12.06 ± 3.91</td><td>11.33 ± 5.06</td><td>2.95 ± 3.51</td><td>8.39 ± 3.64</td></tr><tr><td>W<sub>1526</sub> &lt; 0.8 Å</td><td>0.64 ± 4.55</td><td>-3.66 ± 3.15</td><td>4.31 ± 3.21</td><td>0.80 ± 4.43</td><td>-5.20 ± 3.04</td><td>5.28 ± 3.22</td></tr><tr><td>W<sub>1526</sub> ≥ 0.4 Å<sup>b</sup></td><td>13.49 ± 4.05</td><td>1.16 ± 2.84</td><td>12.34 ± 2.90</td><td>11.34 ± 3.74</td><td>-0.79 ± 2.60</td><td>12.13 ± 2.69</td></tr><tr><td>(r - i) ≥ 0.1</td><td>9.32 ± 4.97</td><td>-0.45 ± 3.56</td><td>9.78 ± 3.48</td><td>4.84 ± 4.59</td><td>-2.57 ± 3.24</td><td>7.41 ± 3.24</td></tr><tr><td>(r - i) &lt; 0.1</td><td>3.45 ± 4.93</td><td>-0.75 ± 3.36</td><td>4.19 ± 3.61</td><td>3.42 ± 4.48</td><td>-1.59 ± 3.04</td><td>5.02 ± 3.30</td></tr></tbody></table>

Notes. <sup>a</sup>The individual errors are the maximum error seen in the 1000 realizations during the bootstrap analysis and not the statistical error of the bootstrap sample.

<sup>b</sup>Only systems with Si II λ1526 absorption line detection are considered.

bootstrap sample is smaller than the uncertainties in each realization (given by the size of the interval including 68 per cent of the pixel value). Therefore, the uncertainty over each pixel is taken to be the maximum error value over 1000 realizations.

In this exercise, we also test the presence of any systematic differences between the integrated flux in the blue and red parts of the DLA core in the stacked spectra. Results of this bootstrap analysis are summarized in Table 3. Here, we consider only the median stacked and 3σ clipped weighted mean spectrum. It is clear from this table that the total integrated luminosity over the core region of the DLA absorption trough is $(5.8 \pm 3.5) \times 10^{40}$ erg s$^{-1}$ and $(5.3 \pm 3.7) \times 10^{40}$ erg s$^{-1}$, respectively, when we consider sub-samples with 90 per cent and 80 per cent of the total sample and median stacked spectrum. In Table 3, we also give the Lyα luminosity integrated over the blue and red parts alone. It is clear that, while the luminosity is consistent with zero in the blue part [i.e. $L_{\lambda}^{b} = (-1.0 \pm 2.4) \times 10^{40}$ erg s$^{-1}$], there is a 2.8σ excess seen in the red part [i.e. $L_{\lambda}^{r} = (6.9 \pm 2.5) \times 10^{40}$ erg s$^{-1}$]. The distribution of the measured luminosities in the blue and red parts in individual median stacked spectrum considering 90 per cent of the total sample in each sub-sample is shown with stars in Fig. 3. First of all, we note that there is no correlation seen between the Lyα luminosity in red and blue parts, with a Spearman's rank correlation coefficient ($r_s$) of 0.16. It is also clear from the figure that the average Lyα luminosity of the blue part is consistent with zero, whereas, the red part always shows a positive luminosity. From Table 3, we also see this trend

being present even when we do bootstrap analysis for 3σ clipped weighted mean spectra albeit with slightly lower significance level. The main conclusion from this section is that the average Lyα luminosity of DLAs with $log N(H_1) \ge 21$ within 8 kpc to the QSO sightline, is less than one hundredth of the Lyα luminosity of an L<sub>λ</sub> galaxy (i.e. $L_{\lambda}(Ly\alpha) = 5.8 \times 10^{42}$ erg s$^{-1}$) at these redshifts (Ouchi et al. 2008). There is a hint of peak emission in the Lyα profile being redshifted with respect to the DLA systemic redshift.

In what follows, we will explore the dependence of Lyα luminosity on different measurable parameters of the DLA.

## 3.3 Dependence on N(H<sub>I</sub>)

The H<sub>I</sub> column density of DLAs is found to be strongly anticorrelated with the impact parameter of the quasar sightline, where a larger N(H<sub>I</sub>) typically originates from a sightline with smaller impact parameter (Zwaan et al. 2005; Péroux et al. 2011; Rao et al. 2011; Krogager et al. 2012). In this scenario, we expect the galaxies associated with high H<sub>I</sub> column density systems to preferably come inside the fibre, whereas those associated with low column density will contribute less to the Lyα emission in the fibre spectrum. Recall, that the fibre radius of 1 arcsec corresponds to a physical size of ~8 kpc at a median z of 2.7. From fig. 8 of Zwaan et al. (2005), we see that a line of sight having $log N(H_1) \ge 21$ will occur, at an impact parameter $\le 8$ kpc, 75 per cent of the times. Whereas, the same for $log N(H_1) \ge 21.7$ will occur within an impact
<figure><img src="image_4.png" /><figcaption>**Figure 3.** Distribution of the measured mean Lyα luminosities on the red ($L_\lambda^r$) and blue ($L_\lambda^b$) part of the DLA core region in the bootstrapped stacked spectra for our full sample (blue star) and for the sub-sample with $W_{1526} \ge 0.8$ Å (cyan open circle), using ~1000 random sub-samples generated by considering only 90 per cent size of the sample. The vertical and horizontal (dashed and dotted) lines represent the mean bootstrapped luminosity in the red and blue parts. The bar shows $1\sigma$ deviation of the bootstrap sample around the mean. No correlation is seen between the $L_\lambda^r$ and $L_\lambda^b$, with a Spearman's rank correlation coefficient ($r_s$) of 0.16 and 0.19, respectively, for both the sub-samples. We also show the histogram distribution of $L_\lambda^r$ and $L_\lambda^b$ in the figure.</figcaption></figure>

parameter of 8 kpc in 90 per cent of the cases. However, results for high-z DLAs suggest that a given $N(H_1)$ could come from a larger impact parameter (see, Péroux et al. 2011; Rao et al. 2011) than the one discussed by Zwaan et al. (2005) based on Hα emission from low-z galaxies.

On the other hand, Toribio et al. (2011) have found that the size of the Hα emitting region in low-z galaxies is proportional to their optical luminosity. Because of this, it is possible that a given Hα column density will originate from a large impact parameter for a high-luminosity galaxy compared to a low-luminosity one. It is also clear from fig. 18 in Zwaan et al. (2005) that when we have $\log N(H_1) \ge 21$, at an impact parameter of 10 kpc, the associated galaxy is most likely to be a sub-L* galaxy. Patra, Chengalur & Begum (2013), using Hα maps of low-z dwarf galaxies found roughly 10 per cent of the cross-section above $\log N(H_1) = 20.3$ at $z=0$ is provided by dwarf galaxies. However, this fraction sharply falls to $\le 1$ per cent by $\log N(H_1) \sim 21.5$ as the cross-section of such gas in dwarf galaxies are very small. Typically these high column density gas are located close to star-forming regions even when there is no one to one correspondence between them (Roychowdhury et al. 2014). In this scenario, for our sample with $\log N(H_1) \ge 21$, it is most likely that the stellar light that will directly go through the fibre will come from less luminous galaxies. The brighter galaxies contributing to the absorption likely fall outside the BOSS fibre (López & Chen 2012).

All this motivates us to explore the Lyα emission as a function of the Hα column density. In order to understand the dependence of Lyα luminosity on $N(H_1)$, we divided the sample into two around log $N(H_1) = 21.23$, the median log $N(H_1)$ of our sample. The median log

$N(H_1)$ of the two sub-samples are 21.1 and 21.4, respectively. The Lyα luminosities measured by integrating over the DLA core region are given in Table 2. For the weighted mean and $3\sigma$ clipped weighted mean, the Lyα luminosity measured in the high $N(H_1)$ sub-sample is slightly higher than those measured from the low column density sub-sample. However, the difference is not statistically significant considering the associated errors in both measurements. We also note that the measurements based on the median stacked spectrum are similar for both sub-samples. We then consider the bootstrapping results presented in Table 3. As in the case of the whole sample, the Lyα luminosities measured for the red part are always higher than those measured for the blue part. However, even for the red part the difference between two sub-samples based on $N(H_1)$ (i.e. $(4.0 \pm 4.9) \times 10^{40}$ erg s$^{-1}$ for the median stacked spectrum) is not statistically significant. However, we wish to note here that the range in $N(H_1)$ is very small in our sample to reveal any weak trend in the Lyα emission with $N(H_1)$.

## 3.4 Dependence on z

Next, we consider two sub-samples based on $z_{abs}$ by dividing the full sample at $z_{abs} = 2.7$. The sub-samples have median redshift of 2.48 and 2.97. From Table 2, we note that the Lyα luminosity measured over the core region appears to be slightly higher for the high-z sample but the difference is not statistically significant. Table 3, also suggests that there is no statistically significant difference in the average Lyα luminosity as a function of z. As the projected size of the fibre in the whole z range, over which our DLAs are detected, is nearly the same, it does not influence the z evolution. However, in all sub-samples the Lyα luminosity inferred is higher for the red part compared to that of the blue part, albeit with less statistical significance, compared to what we find for the full sample. We note that the smaller z range in our sample provides a limited leverage to reveal any weak trend of Lyα emission with z.

## 3.5 Dependence on $W_\alpha$ (Si II λ1526)

Prochaska et al. (2008) found a strong correlation between the rest equivalent width of Si II λ1526 line (denoted as $W_{1526}$) and metallicity (see also Jorgenson, Murphy & Thompson 2013; Neeleman et al. 2013). They found this correlation to be stronger than the correlation found between low-ion velocity widths ($\Delta v_{90}$) and metallicity (Ledoux et al. 2006). The correlation between $W_{1526}$ and metallicity has also been confirmed by Kaplan et al. (2010) in the case of metal strong DLAs. The measured correlation implies an average metallicity of $Z = -1.48 Z_\odot$ (respectively, $-1.06 Z_\odot$) for $W_{1526} = 0.4$ Å (respectively, 0.8 Å) with a scatter of 0.25 dex in the metallicity. The average metallicity will be 0.15 dex higher if we use the coefficients found by Jorgenson et al. (2013). As Si II λ1526 is a strong transition, the absorption line profile will be highly saturated for the high equivalent width systems. Therefore, the measured equivalent width usually samples the velocity field probed by the optically thin gas that is either part of an outflow or in the halo. Thus, this correlation was interpreted as a reflection of the underlying mass-metallicity relation in galaxies (Ledoux et al. 2006; Christensen et al. 2014).

We find the median $W_{1526} = 0.8$ Å, if we consider DLAs with a clear detection of Si II λ1526 absorption line. Therefore, we consider two sub-samples of DLAs with $W_{1526} < 0.8$ Å and $\ge 0.8$ Å. The median $W_{1526}$ of these two sub-samples are 0.53 Å and 1.18 Å. The DLAs without Si II λ1526 absorption line detections are also included in the $W_{1526} < 0.8$ Å sub-sample. From Table 2 (and Fig. 4), we can see that when we consider only systems with
<figure><img src="image_5.png" /><figcaption>**Figure 4.** Same as Fig. 2, for the sub-sample with $W_{1526} \ge 0.8$ Å. The dashed curve is the Lyα emission in the composite spectrum of LBGs (Shapley et al. 2003) with its amplitude scaled to match the luminosity seen in our stacked spectrum.</figcaption></figure>

$W_{1526} \ge 0.8$ Å, the integrated Lyα luminosity over the full DLA core itself is significantly more from zero by more than $3\sigma$ level (see Table 2). We find $L = (15.8 \pm 5.2) \times 10^{40}$ erg s$^{-1}$, a detection at $3\sigma$ level, in the median stacked spectrum. The weighted mean and $3\sigma$ clipped weighted mean stacked spectra also show consistent results albeit with slightly less significance. The measured Lyα luminosity in this case corresponds to 2–2.5 per cent of the Lyα luminosity of an $L_\alpha$ galaxy at these redshifts. However, for the sub-sample of systems having $W_{1526} < 0.8$ Å, we find $L (\simeq 4.7 \pm 5.0) \times 10^{40}$ erg s$^{-1}$ in the median stacked spectrum. Consistently lower luminosities are also seen in the case of weighted mean and $3\sigma$ clipped weighted mean stacked spectra. The difference in the Lyα luminosities between high and low equivalent width sub-samples is $(11.1 \pm 7.2) \times 10^{40}$ erg s$^{-1}$ for the median stacked spectrum. Therefore, we have a marginal evidence for the Lyα emission being high for high Si II equivalent width systems.

Next, we consider $L_\lambda^b$ and $L_\lambda^r$ for the two sub-samples. For the $W_{1526} \ge 0.8$ Å sub-sample, we find the $L_\lambda^r \sim (12.1 \pm 3.9) \times 10^{40}$ erg s$^{-1}$ is higher than $L_\lambda^b \sim (2.5 \pm 3.9) \times 10^{40}$ erg s$^{-1}$ in the median stacked spectrum. The difference, $(9.6 \pm 5.5) \times 10^{40}$ erg s$^{-1}$, is significant at a $1.7\sigma$ level. The result of bootstrap analysis retaining 90 per cent of the data is also shown in Fig. 3. First of all we notice that, unlike the full sample, the scatter is much higher in the $W_{1526}$ based sub-sample. However, it is clear that while there is no statistically significant detection of emission in the blue part, in the red part non-zero fluxes are detected at more than $3\sigma$ level. Interestingly, even in the $W_{1526} < 0.8$ Å sub-sample $L_\lambda^r = (4.3 \pm 3.2) \times 10^{40}$ erg s$^{-1}$ is stronger than $L_\lambda^b = (-3.7 \pm 3.2) \times 10^{40}$ erg s$^{-1}$ in the median stacked spectrum. We notice that $L_\lambda^r$ in the case of $W_{1526} \ge 0.8$ Å is higher than $L_\lambda^r$ measured for $W_{1526} < 0.8$ Å sub-

sample with the difference being $(5.1 \pm 5.5) \times 10^{40}$ erg s$^{-1}$. As can be seen from Table 3, these trends are also seen when we do the bootstrapping analysis using $3\sigma$ clipped weighted mean statistics. However, the difference in the total Lyα luminosity and $L_\lambda^r$ between two sub-samples are weaker when we use $3\sigma$ clipped weighted mean stacked spectrum instead of median stacked spectrum. Similar trends are also seen if we consider the results for CNR $\ge 4$ cutoff (see Table A1 in the appendix). Furthermore, we have tested the dependence of the strength of the Lyα emission on $W_{1526}$ cutoff by decreasing it to 0.4 Å and 0.6 Å and find no strong dependence on this cutoff (see Table A1, Fig. A2).

Results presented here are consistent with Lyα emission being stronger among systems having large $W_{1526}$. We also note that the distribution of $N(H\alpha)$ and $z$ are similar (see the quoted median values in Table 2) between different sub-samples discussed here purely based on $W_{1526}$. We discuss the implications of these results in detail in Section 5.

## 3.6 Dependence of QSO colours

The general population of DLAs tend to have very little dust, so they usually do not cause a strong reddening to the spectra of background QSOs (see, Vladilo, Prochaska & Wolfe 2008; Frank & Péroux 2010; Khare et al. 2012; Fukugita & Ménard 2015; Murphy & Bernet 2016). However, due to the resonant nature of the Lyα transition, presence of even a small amount of dust can modify the emerging Lyα line intensity and profile shape. In addition, the spatial distribution of dust can also change the profile of the emerging Lyα emission.

In this section, we study the possible connection between the QSO colour and the stacked Lyα emission. For a given dust content in DLAs, the observed colour of a QSO at a given redshift will also depend on the absorption redshift. We consider two sub-samples based on $(r-i)$ colour of the QSO.

From Table 2, we see that when we divide our whole sample into two based on $(r-i)$ colours, there is no statistically significant difference in the Lyα luminosity between the two sub-samples even though we see slightly higher value of the Lyα luminosity (the difference is only at $1.1\sigma$ level) for the redder QSOs when we consider the median stacked spectrum. From Table 3, we also notice that the red part shows a high luminosity of $L_\lambda^r = (9.8 \pm 3.5) \times 10^{40}$ erg s$^{-1}$ (at $\sim 2.8\sigma$) for the high $(r-i)$ sample. Whereas, the low $(r-i)$ sample shows an $L_\lambda^r = (4.2 \pm 3.6) \times 10^{40}$ erg s$^{-1}$ with an insignificant (only at $\sim 1.2\sigma$) luminosity difference between the two sub-samples. One needs to be careful in interpreting these results, as the difference in $(r-i)$ between the two sub-samples is not much and also there is a possible correlation between $(r-i)$ and $W_{1526}$. Just to explore the role played by the possible presence of dust, we restrict our future discussions to systems with $W_{1526} \ge 0.4$ Å. This is because we already have a clear detection of Lyα emission in this case. Thus, further dividing the sub-sample based on $(r-i)$ is more relevant for probing the connection between reddening and Lyα emission profile.

We compute the $\Delta(r-i)$ colours of QSOs having DLAs in our sample with $W_{1526} \ge 0.4$ Å by finding the deviation of the $(r-i)$ colour of the QSO with respect to the median $(r-i)$ colour of SDSS-DR12 QSOs at the same redshift. The $\Delta(r-i)$ is found to be strongly correlated with $(r-i)$ with an $r_s$ of 0.88. In Fig. 5, we show the histogram of $\Delta(r-i)$. The vertical dashed and dotted lines show the median $\Delta(r-i)$ and associated $1\sigma$ spread for the general QSOs with the same redshift range of quasars used in this work.
<figure><img src="image_6.png" /><figcaption>**Figure 5.** $\Delta(r - i)$ values of the background quasars for the sub-sample with $W_{1526} \ge 0.4$ Å, with a median value is shown with a blue arrow. The $\Delta(r - i)$ distribution for the lower and upper 30, 40, 50 per cent of the sub-samples from $(r - i)$ colour distribution is shown with blue, cyan, green, yellow, orange and red colours, respectively. The dashed line represents the median $\Delta(r - i)$ colour for the SDSS-DR12 QSOs with a similar redshift range of the quasars used in this work (i.e. $2.38 \le z_{em} \le 3.95$). The dotted lines show the one standard deviation away from that median.</figcaption></figure>

The QSOs in the DLA sample are slightly redder than the general population of QSOs. The presence of QSOs with negative $\Delta(r - i)$ indicates the possible intrinsic spread in the QSO continuum shape. We divide the $W_{1526} \ge 0.4$ Å sub-sample into six different groups by selecting the lower 30, 40, 50 and upper 50, 40 and 30 per cent DLAs from the cumulative distribution of measured $(r - i)$ colour. Note that, the average $W_{1526}$ for the lower and upper 30 per cent sub-samples are found to be $0.83 \pm 0.13$ Å and $1.04 \pm 0.16$ Å, respectively. In Table 4, we summarize our results. First six columns summarize the sample, range in $(r - i)$ considered, mean $\Delta(r - i)$ of the group, the total Ly$\alpha$ luminosity, $L_{\lambda}^{b}$ and $L_{\lambda}^{r}$. More information can be obtained from the visual inspection of the profiles given in Fig. 6.

It is clear from Table 4 that when we consider the stacked spectrum of lower 30 per cent of the quasars from $(r - i)$ colour distribution, the integrated Ly$\alpha$ luminosity over the full profile is found to be $(22.9 \pm 7.2) \times 10^{40}$ erg s$^{-1}$. Unlike all the other cases discussed till now, we do see a clear double hump profile for the Ly$\alpha$ emission in this case (see the bottom panels of Fig. 6). Double hump

feature for the lower 30 per cent of the quasars from $(r - i)$ colour distribution is also clearly seen even in the $3\sigma$ clipped mean stacked spectrum (see right lowermost panel in Fig. 6). It is also consistent with the finding of Noterdaeme et al. (2014), where excess Ly$\alpha$ emission is seen in ESDLAs (i.e. $\log N(H_1) \ge 21.7$) with lower $E(B - V)$ (i.e. $<0.025$). Results of double Gaussian fit to the observed Ly$\alpha$ profile for the median stacked spectra are also summarized in Table 4 (see columns 7–9). From Table 4, we notice that in the two low $(r - i)$ bins within the measurement errors, $L_{\lambda}^{r}$ is twice that of $L_{\lambda}^{b}$ and the blue peak has a slightly larger velocity shift with respect to the DLA redshift compared to the red peak (see columns 7 and 8 in Table 4). Table 4 also gives the deconvolved FWHM of the red peak obtained through Gaussian fitting (column 9). In Section 5.1, we use these values to compare DLAs with high-redshift LBGs and LAEs.

It is clear from the bottom three panels in Fig. 6 and top three columns in Table 4 that as we keep adding DLAs along the QSOs with high $(r - i)$ colours, the measured Ly$\alpha$ luminosity over the DLA core as well as $L_{\lambda}^{b}$ and $L_{\lambda}^{r}$ decrease. In addition, the significance of blue as well as red peaks decreases. This trend is in line with our expectation that the presence of dust in DLAs will destroy the Ly$\alpha$ photons. Based on this trend, naively we would expect the double hump to disappear when we consider sub-sample of DLAs towards QSOs having $(r - i)$ in the upper 50 per cent. While this is the case (see the second panel from top in Fig. 6), we notice that the total integrated Ly$\alpha$ luminosity over the core region is similar to what we find for the lower 50 per cent case. This is mainly because of the detection of non-zero luminosity (detected at $\sim 3.4\sigma$ level) in the red part. Indeed, when we consider only the upper 30 per cent of sightlines based on $(r - i)$ colours, the Ly$\alpha$ luminosity we measure is close to what we measure for the lower 30 per cent case albeit without a clear detection of the double hump. While the disappearance of the double hump (or decreasing flux in the blue part) with increasing $(r - i)$ is consistent with the effect of dust, the lack of systematic change in the integrated Ly$\alpha$ luminosity suggests that there are other effects at play while we change $(r - i)$.

It is to be remembered that when we consider the upper 30 per cent, we probably have a mixed population of QSOs, i.e. QSOs with bluer intrinsic colour reddened by the dust in the DLAs and QSOs that are intrinsically redder compared to the median QSO spectrum. Even if Ly$\alpha$ is present in all these cases, the profile shapes of the two populations are expected to be different that can lead to an incoherent addition and lack of a well-defined shape for Ly$\alpha$ emission. In addition, from fig. 3 of Krogager et al. (2012) one can see that the high H I column density systems may originate from smaller impact parameter when the metallicities are low. If this is the case with the low $(r - i)$ sub-sample then the fibre filling factor will also have a role to play. We also notice a possible correlation between $\Delta(r - i)$ and $W_{1526}$. As seen in the previous

<table><caption>**Table 4.** Ly$\alpha$ luminosity for the sub-sample with $W_{1526} \ge 0.4$ Å and various $(r - i)$ colour cut.</caption><thead><tr><th rowspan="2">Sample%</th><th rowspan="2">(r - i)</th><th rowspan="2">&lang;&Delta;(r - i)&rang;</th><th colspan="3">Ly&alpha; luminosity (x10<sup>40</sup> erg s<sup>-1</sup>)</th><th colspan="2">Velocity offset</th><th rowspan="2">FWHM (km s<sup>-1</sup>)</th></tr><tr><th>DLA bottom</th><th>L<sub>&lambda;</sub><sup>b</sup></th><th>L<sub>&lambda;</sub><sup>r</sup></th><th>&Delta;&nu;<sup>r</sup></th><th>&Delta;&nu;<sup>b</sup></th></tr></thead><tbody><tr><td>Lower 30%</td><td>&lt;0.05</td><td>-0.06</td><td>22.89 &pm; 7.21</td><td>6.02 &pm; 4.79</td><td>16.87 &pm; 5.38</td><td>298 &pm; 30</td><td>-368 &pm; 15</td><td>196 &pm; 70</td></tr><tr><td>Lower 40%</td><td>&lt;0.08</td><td>-0.04</td><td>21.29 &pm; 6.38</td><td>7.05 &pm; 4.28</td><td>14.23 &pm; 4.74</td><td>336 &pm; 26</td><td>-363 &pm; 16</td><td>144 &pm; 62</td></tr><tr><td>Lower 50%</td><td>&lt;0.11</td><td>-0.03</td><td>14.00 &pm; 5.82</td><td>2.71 &pm; 3.92</td><td>11.29 &pm; 4.30</td><td>272 &pm; 46</td><td>-337 &pm; 26</td><td>214 &pm; 98</td></tr><tr><td>Upper 50%</td><td>&gt;0.11</td><td>0.10</td><td>13.54 &pm; 5.61</td><td>0.34 &pm; 4.04</td><td>13.20 &pm; 3.89</td><td>-</td><td>-</td><td>-</td></tr><tr><td>Upper 40%</td><td>&gt;0.13</td><td>0.10</td><td>20.76 &pm; 6.24</td><td>2.96 &pm; 4.55</td><td>17.80 &pm; 4.28</td><td>-</td><td>-</td><td>-</td></tr><tr><td>Upper 30%</td><td>&gt;0.16</td><td>0.16</td><td>23.15 &pm; 7.18</td><td>6.08 &pm; 5.37</td><td>17.06 &pm; 4.77</td><td>-</td><td>-</td><td>-</td></tr></tbody></table>
<figure><img src="image_7.png" /><figcaption>**Figure 6.** Left-hand panel (bottom to top): the median stacked spectra for various sub-samples with $W_{1526} \ge 0.4$ Å with colour selection criteria of ($r - i$) <0.05, <0.08, <0.11, >0.11, >0.13 and >0.16. The blue segment shows the DLA core with $\tau \ge 10$ for $\log N(H_i) = 21.0$. The dashed curves show the synthetic profiles for lower (i.e. $\log N(H_i) = 21$) and median (i.e. $\log N(H_i) = 21.23$) column density of DLAs used to get the stacked spectrum. The solid red curve shows the Gaussian fit to the non-zero flux seen in the red and blue parts of the DLA core region. Right-hand panel: same as left-hand panel for $3\sigma$ clipped weighted mean stacks. See Fig. A3 in the appendix for results when we consider systems with $\text{CNR} \ge 4$. In each panel, the number of DLAs contributing to the plot are also indicated.</figcaption></figure>

section, systems with high $W_{1526}$ tend to have high Ly$\alpha$ luminosity. Therefore, in the sub-samples based on ($r - i$), it is not necessary that the only quantity that varies is the amount of reddening. Therefore interpretation of these results is not straightforward.

## 4 ABSORPTION LINES IN THE STACKED SPECTRUM

We also generated sets of median stacked spectra to study absorption lines, as described in section 3, for various sub-samples discussed in the previous section. We detect transitions of the low-ionization species (e.g. Fe II, Si II, Zn II, Cr II, Mg I), high-ionization species (Si IV, C IV) and weak transition lines (Fe II $\lambda\lambda$2249, 2260) as found earlier in the stacked SDSS spectra by York et al. (2006), Khare et al. (2012) and Noterdaeme et al. (2014). In addition, we clearly detect Fe II $\lambda$1611, Ni II $\lambda$1454 and Ti II $\lambda$1910 absorption which were marginally detected in the stacked spectrum of extreme-DLAs studied by Noterdaeme et al. (2014). The rest equivalent width ($W_r$) measured for various metal lines in the stacked spectrum of our entire sample as well as different sub-samples, with a single Gaussian fit, are listed in Table A3. A comparison of profiles of different absorption lines seen between different sub-samples is also shown in Fig. A4. We note that most lines are saturated even if the apparent optical depths of the strong absorption lines are small and appear unsaturated due to insufficient BOSS spectral resolution ($R \sim 2000$) and smoothing, resulting from redshift uncertainties when co-adding the spectra (York et al. 2006; Noterdaeme et al. 2014; Sardane, Turnshek & Rao 2015).

Next, using different absorption lines of Fe II and Si II, that originate from transitions with a broad range of oscillator strengths detected in our stacked spectra, we compute the column density and the effective Doppler parameter ($b_{\text{eff}}$) by constructing a single cloud curve growth (see left-hand panel of Fig. 7 for the sub-samples based on $W_{1526}$). It is clear that $b_{\text{eff}}$ is a function of $W_{1526}$. Jenkins (1986) has shown that a single cloud curve of growth (COG) gives nearly correct column density when using the $W_r$ of several lines with the line distribution function not having markedly irregular characteristic, even when different lines have large variation in central optical depth and inter velocity dispersion. While spectral resolution of BOSS spectra will not allow us to measure $N$ and $b$ parameters accurately, the purpose of using the single cloud COG is to find some tracers of column density and velocity field (i.e. $b_{\text{eff}}$).

In Fig. 8, we show the integrated Ly$\alpha$ luminosity in the DLA core region (lower panel) and in the red part of the DLA trough (upper panel) versus $b_{\text{eff}}$. In this figure, we mainly focus on the four sub-samples defined only based on $W_{1526}$. The $b_{\text{eff}}$ values for various sub-samples are given in Table 5 (column 2). These values are much higher than what one measures in individual components seen in high-resolution echelle spectra. While comparing the $\Delta v_{90}$ and $b_{\text{eff}}$ measured from high-resolution spectrum, Noterdaeme et al. (2014) have obtained an empirical relation of $\Delta v = 2.2 \ b_{\text{eff}} + 0.02 \ b_{\text{eff}}^2$ which at larger $b_{\text{eff}}$ departs from the linear theoretical relation in the Gaussian regime, $\Delta v = 2.33 \ b_{\text{eff}}$ (see their fig. 7). Therefore, $b_{\text{eff}}$ we measure is an indication of overall velocity field and not related to the thermal or microturbulent motion.

Interestingly, we notice nearly the same $b_{\text{eff}}$ (i.e. between 35 and $38 \text{ km s}^{-1}$) for the sub-samples divided based on $N(H_i)$ and
<figure><img src="image_8.png" /><figcaption>**Figure 7.** Left-hand panel: single component COG for absorption lines detected in the median stacked spectrum for the sub-samples with $W_{1526} \ge 0.8$ Å (open squares and solid line) and <0.8 Å (filled squares and dotted line). The best-fitted $b_{\text{eff}}$ obtained are also quoted in the figure. Right-hand panel: same for the stacked spectrum of $W_{1526} \ge 0.4$ Å and considering upper and lower 30 per cent samples from $(r-i)$ colour distribution.</figcaption></figure>

<figure><img src="image_9.png" /><figcaption>**Figure 8.** Lower panel: the dependence of $\text{Ly}\alpha$ luminosity (integrated over the entire DLA core region) on the Doppler parameter ($b_{\text{eff}}$) for the sub-samples with $W_{1526} \ge 0.6$, 0.8 and <0.6, 0.8 Å. Upper panel: same for the luminosity seen in the red part of DLA bottom. Note that the significant detection (i.e. >3$\sigma$ level) is only for three cases with $W_{1526} \ge 0.6$ and 0.8 Å.</figcaption></figure>

$z$ (not shown in this figure). Remember, these sub-samples do not show significant difference in their $\text{Ly}\alpha$ luminosities. Also, all of them tend to have less $\text{Ly}\alpha$ luminosity compared to the high $W_{1526}$ sub-sample. Note, Noterdaeme et al. (2014) have found a $b_{\text{eff}}$ =

40 km s⁻¹ for their ESDLAs with similar analysis. An increasing trend of $\text{Ly}\alpha$ luminosity with $b_{\text{eff}}$ is clearly seen while considering the total luminosity in the DLA core as well as the $L_{\text{\textit{\lambda}}}^r$ (see Fig. 8). These trends are consistent with high $W_{1526}$ systems originating from gas with high velocity fields and also tend to show higher $\text{Ly}\alpha$ luminosity.

Similar dependences are also seen between [O II] luminosity in the stacked spectrum and the equivalent width of Ca II and Mg II absorption lines in the low-$z$ absorbers (Wild et al. 2007; Noterdaeme et al. 2010; Ménard et al. 2011). This could be due to some physical connection between the SFR and velocity width (like for example, large-scale outflows as suggested by Ménard et al. 2011) or an artefact of known correction between the line equivalent width and impact parameter coupled with fibre filling factor decreasing with the impact parameter and hence the rest equivalent width (López & Chen 2012). Both these explanations will work for the $\text{Ly}\alpha$ emission also. In addition, we also have the possibility that the $\text{Ly}\alpha$ escape fraction being higher at higher $W_{\text{f}}$ (i.e. velocity field). Note that, from the observed $W_{1526}$ and metallicity correlation, the $\text{Ly}\alpha$ emission will also depend on the metallicity. These trends are consistent with the suggestion of Möller et al. (2004) based on DLA galaxies detected in their NICMOS sample. However, no such relation is seen with the $z$ and $N(H\text{i})$.

Next we consider two sub-samples based on $(r-i)$ colours of the background QSOs. As discussed in the previous section, the total $\text{Ly}\alpha$ luminosity is nearly the same between the sub-samples consisting of the lower and upper 30 per cent sources. This was unexpected based on a simple conjecture that systems towards QSOs with high $(r-i)$ colours and higher dust content should have less $\text{Ly}\alpha$ luminosity. The stacked spectra show $b_{\text{eff}} = 39$ km s⁻¹ for the lower 30 per cent sub-sample and $b_{\text{eff}} = 54$ km s⁻¹ for the upper 30 per cent sub-sample (see right-hand panel in Fig. 7). If one follows the idea of higher the $b_{\text{eff}}$ higher will be the intrinsic $\text{Ly}\alpha$ luminosity then the higher $\text{Ly}\alpha$ luminosity in the upper 30 per cent can be understood as the consequence of systems having more intrinsic $\text{Ly}\alpha$ luminosity.
<table><caption>**Table 5.** Parameters estimated using absorption stacks.</caption><thead><tr><th>Sample criteria</th><th>b<sub>eff</sub> (km s<sup>-1</sup>)</th><th>&Delta;v<sub>90</sub> (km s<sup>-1</sup>)</th><th>W<sub>1526</sub> (&Aring;)</th><th>FWHM(C<sub>IV</sub>) (km s<sup>-1</sup>)</th><th>DR<sup>a</sup></th><th>FWHM(Si II) (km s<sup>-1</sup>)</th></tr></thead><tbody><tr><td>W<sub>1526</sub> &ge; 0.8 &Aring;</td><td>59.3 &plusmn; 1.3</td><td>201</td><td>1.21 &plusmn; 0.01</td><td>252.8 &plusmn; 4.2</td><td>1.75</td><td>240.1 &plusmn; 2.4</td></tr><tr><td>W<sub>1526</sub> &lt; 0.8 &Aring;</td><td>25.1 &plusmn; 0.8</td><td>76</td><td>0.43 &plusmn; 0.01</td><td>220.1 &plusmn; 6.2</td><td>1.63</td><td>125.4 &plusmn; 2.2</td></tr><tr><td>W<sub>1526</sub> &ge; 0.6 &Aring;</td><td>50.8 &plusmn; 1.0</td><td>164</td><td>1.03 &plusmn; 0.01</td><td>251.4 &plusmn; 4.1</td><td>1.72</td><td>213.6 &plusmn; 2.3</td></tr><tr><td>W<sub>1526</sub> &lt; 0.6 &Aring;</td><td>21.9 &plusmn; 0.6</td><td>61</td><td>0.37 &plusmn; 0.01</td><td>197.7 &plusmn; 7.5</td><td>1.63</td><td>114.0 &plusmn; 2.6</td></tr><tr><td>W<sub>1526</sub> &ge; 0.4 &Aring;</td><td>45.3 &plusmn; 1.0</td><td>141</td><td>0.85 &plusmn; 0.01</td><td>250.3 &plusmn; 4.0</td><td>1.74</td><td>192.7 &plusmn; 2.2</td></tr><tr><td>N(HI) &ge; 10<sup>21.23</sup> cm<sup>-2</sup></td><td>34.4 &plusmn; 1.0</td><td>100</td><td>0.75 &plusmn; 0.01</td><td>245.0 &plusmn; 5.0</td><td>1.61</td><td>180.0 &plusmn; 3.1</td></tr><tr><td>N(HI) &lt; 10<sup>21.23</sup> cm<sup>-2</sup></td><td>36.0 &plusmn; 1.3</td><td>106</td><td>0.61 &plusmn; 0.01</td><td>237.7 &plusmn; 6.4</td><td>1.71</td><td>167.4 &plusmn; 3.6</td></tr><tr><td>z<sub>abs</sub> &ge; 0.7</td><td>37.2 &plusmn; 1.8</td><td>110</td><td>0.65 &plusmn; 0.11</td><td>228.3 &plusmn; 6.3</td><td>1.70</td><td>171.3 &plusmn; 3.6</td></tr><tr><td>z<sub>abs</sub> &lt; 0.7</td><td>37.8 &plusmn; 1.0</td><td>112</td><td>0.69 &plusmn; 0.10</td><td>246.6 &plusmn; 5.1</td><td>1.66</td><td>175.0 &plusmn; 3.1</td></tr><tr><td>(r &minus; i) upper 30 per cent</td><td>54.2 &plusmn; 1.7</td><td>179</td><td>1.05 &plusmn; 0.02</td><td>240.5 &plusmn; 6.5</td><td>1.83</td><td>222.4 &plusmn; 4.2</td></tr><tr><td>(r &minus; i) lower 30 per cent</td><td>38.9 &plusmn; 1.1</td><td>116</td><td>0.83 &plusmn; 0.01</td><td>270.4 &plusmn; 7.3</td><td>1.68</td><td>184.6 &plusmn; 3.5</td></tr></tbody></table>

Note. <sup>a</sup>C IV doublet ratio, given as $DR = W(C_{IV1548})/W(C_{IV1550})$.

<figure><img src="image_10.png" /><figcaption>**Figure 9.** The dependence of Lyα luminosity (integrated over the entire DLA core region) on the equivalent width, FWHM of C IV absorption and the doublet ratio (i.e. $W_{1548}/W_{1550}$) for the sub-samples with $W_{1526} \ge 0.6$, 0.8 and $<0.6$, 0.8 Å.</figcaption></figure>

In Table 5 (see column 3), we present the velocity width ($\Delta v_{90}$) estimated from the relationship between $\Delta v_{90}$ and $b_{eff}$ obtained by Noterdaeme et al. (2014, see above). There is a hint of excess Lyα emission with the large values of $\Delta v_{90}$. It is clear that the inferred $\Delta v_{90}$ is more than 100 km s$^{-1}$ in cases where we have significant Lyα detection. Interestingly, we note that in case of direct detection of Lyα emission from DLA host galaxies, the $\Delta v_{90}$ is found to be greater than ~100 km s$^{-1}$ (see, Fynbo et al. 2010; Noterdaeme et al. 2012a; Hartoog et al. 2015; Srianand et al. 2016). In cosmological hydrodynamic simulations, Bird et al. (2015) have shown that the velocity width closely tracks the virial velocity of the DLA host halo, and thus provides an analogue to the halo mass. In their models, $\Delta v_{90} \sim 100$ km s$^{-1}$ will correspond to typically halo mass of $10^{10}-10^{11}M_{\odot}$. In such haloes, higher gas accretion rate could lead to high SFR, hence higher intrinsic Lyα and UV emission, albeit with a very small Lyα escape fraction (Garel et al. 2015). Alternatively, the observed low Lyα luminosity could be related to the diffuse Lyα emission from the extended region of high-mass galaxies. As we noted above, the average luminosity is much less (i.e. 2–2.5 per cent) than that of an $L_{*}$ galaxy. We discuss these possibilities in the following section.

We now investigate the dependence of Lyα emission on the high-ionization C IV line which probes a wider region of velocity space than the neutral lines (Ledoux et al. 1998; Wolfe & Prochaska 2000; Fox et al. 2007). Fox et al. (2007) have shown that high-velocity C IV clouds are unbound to the central potential well and trace the outflows. In Fig. 9, we plot the Lyα luminosity, integrated over the

entire DLA bottom, as a function of C IV equivalent width, FWHM and the C IV doublet ratio (i.e. $DR = W(C_{IV1548})/W(C_{IV1550})$) for the sub-samples based on $W_{1526}$. The measured line widths (FWHM), corrected for instrumental resolution, DR and width of the Si II line are also summarized in Table 5. As noted by Fox et al. (2007), C IV lines are wider than Si II lines. A trend of increasing Lyα luminosity with increasing $W_{f}$ and FWHM of C IV is clearly noticeable. In addition, the larger DR (i.e. optically thin gas) and higher equivalent width suggest that the winds from the galaxy are responsible for the absorption (Bouché et al. 2006; Christensen et al. 2009). From Table 5, it is clear that there is very little difference in C IV (FWHM) between the sub-samples defined based on ($r - i$) colours. The correlation between $L_{\lambda}$ and $b_{eff}$ is not present when we consider the two sub-samples based on ($r - i$) colours as discussed before. This once again reiterates the importance of dust and difficulties related to interpreting Lyα profile in the stacked spectrum.

In summary, the Lyα emission from DLAs may be originating from systems with large velocity widths in both high and low ions. There are also indications that the metallicity and dust may have an important role to play in deciding the total Lyα luminosity and the profile shape.

## 5 DISCUSSION

In this section, we discuss implications of our Lyα measurements for different stacked spectra and draw some broad conclusions on the nature of high-z DLAs.
## 5.1 Lyα emission line profile

The Lyα emission line profile is a powerful probe of the star formation and associated feedback processes (i.e. infall and/or outflows) in young galaxies. In an extremely opaque static medium Lyα escapes through successive resonance scattering leading to a double-humped profile, with the position of the peaks determined by column density, temperature and kinematics of the medium (Neufeld 1990, 1991; Dijkstra et al. 2006; Hansen & Oh 2006; Verhamme et al. 2006; Dijkstra 2014). In addition, presence of bulk motions (e.g. outflow/inflow) and scattering by dust in the H₁ gas can further modify the emergent Lyα profile. For instance, the scattering through an outflowing (respectively, inflowing) medium results in an overall redshift (respectively, blueshift) of the Lyα spectral line with enhanced red (respectively, blue) peak and suppressed blue (respectively, red) peak (Dijkstra et al. 2006; Barnes & Haehnelt 2010).

We detect the Lyα emission predominantly in the red part of the Lyα trough. This implies the presence of outflows in high-z DLAs. We detect flux in the red part at $2.8\sigma$ when we consider the full sample and with more than $3\sigma$ level when we consider the sub-samples with $W_{1526} \ge 0.4$ Å, $0.6$ Å and $0.8$ Å (see Tables 2 and A1). In these cases, the average Lyα luminosity is found to be consistent within the ~3σ upper limit of ~21.8 × $10^{40}$ erg s⁻¹ obtained by co-adding 341 low column density ($\log N(H_1) \ge 20.62$) DLAs by Rahmani et al. (2010). However, less than the Lyα luminosity of ($60.0 \pm 20.0$) × $10^{40}$ erg s⁻¹, seen in the composite spectrum of ESDLAs ($\log N(H_1) \ge 21.7$) by Noterdaeme et al. (2014). We find the Lyα luminosity being higher for systems with higher $W_{1526}$ which indicates that high-metallicity DLAs could be associated with high star-forming galaxies. Recently, the detection rate of DLA host galaxies is found to be higher towards high-metallicity DLAs (Fynbo et al. 2010, 2011, 2013; Krogager et al. 2012, 2013). These galaxies show either no Lyα emission or a suppressed blue peak.

For the sub-samples with $W_{1526} \ge 0.4$ Å, $0.6$ Å and $0.8$ Å with clear Lyα detection, we model the red peak with a Gaussian profile and measure the centroid shift with respect to the redshift defined by the absorption lines detected in the quasar spectrum. As evident from Fig. 4, the final profile shape of the Lyα emission in the stacked spectrum depends on the statistics used. When we use median combined spectrum, we measure $\Delta v_{Ly\alpha}^r = 385 \pm 28, 331 \pm 28$ and $337 \pm 32$ km s⁻¹ for the sub-samples with $W_{1526} \ge 0.4$ Å, $0.6$ Å and $0.8$ Å, respectively. When we use the weighted mean spectrum, we get $\Delta v_{Ly\alpha}^r = 321 \pm 32, 395 \pm 34$ and $295 \pm 56$ km s⁻¹ for the three cases. From Table 2, we can see that the mean $N(H_1)$ in all these cases are nearly the same, i.e. $\log N(H_1) \sim 21.25$. In a pure static medium, we expect $\Delta v_{Ly\alpha}^r$ in the range 613–417 km s⁻¹ if we assume H₁ gas temperature to be in the range, $10^3–10^4$ K (Dijkstra 2014, see their equation 21).

Interestingly, the Lyα emission line profiles associated with a handful of confirmed DLA host galaxies are found to be predominantly redshifted with typical velocity offsets ranging from 10 to 400 km s⁻¹ (Møller & Warren 1993; Møller et al. 2002; Møller et al. 2004; Fynbo et al. 2010; Krogager et al. 2012; Hartoog et al. 2015; Srianand et al. 2016). Thus, what we measure in the stacked spectra are consistent with these individual measurements. In addition, Noterdaeme et al. (2014) have shown that the Lyα emission in the composite spectrum of ESDLAs also shows a redshifted profile relative to the systemic redshift. Till now only one DLA with Lyα emission ($z_{abs} = 2.207$ towards the quasar SDSS J113520.39–001053.56) shows a classical double hump profile

(Kulkarni et al. 2012; Noterdaeme et al. 2012a). Even in this case, detailed modelling of the Lyα radiative transport favours the emission being scattered from an outflowing gas (Noterdaeme et al. 2012a).

In the case of high-z LBGs and LAEs whenever systematic redshifts can be determined using optical nebular emission lines, the Lyα emission line profile is found to be asymmetric with the peak emission being redshifted in the majority of the cases. The typical peak shift measured is ~200 km s⁻¹ in the case of LAEs (Erb et al. 2014; Shibuya et al. 2014; Song et al. 2014; Trainor et al. 2015) and ~400–600 km s⁻¹ in the case of LBGs (see, Steidel et al. 2003, 2010; Kulas et al. 2012). It has also been suggested that the shifts seen in the case of LAEs and LBGs are statistically different. This could be attributed either to the differences in $N(H_1)$ or to the H₁ covering factor. Taking it at face value, the $\Delta v_{Ly\alpha}^r$ we measure for DLAs is higher than the mean value typically seen in the case of LAEs but less than that of LBGs.

As discussed in Section 3.6, a clear double hump profile is detected when we use DLAs detected towards bluer quasars. The two peaks are separated ~650 km s⁻¹, $L_{\lambda}^r/L_{\lambda}^b \sim 2$ and red part has a FWHM of ~200 km s⁻¹. About 30 per cent of LBGs (Kulas et al. 2012) and 50 per cent of LAEs (Yamada et al. 2012; Hashimoto et al. 2015) show multiple peaks in their Lyα emission with a typical peak separations of ~500 and 800 km s⁻¹, respectively, for LAEs and LBGs (Kulas et al. 2012; Hashimoto et al. 2015; Trainor et al. 2015). In addition, the mean FWHM of the red peak of Lyα line in LAEs and LBGs are found to be ~260 and 364 km s⁻¹, respectively (Trainor et al. 2015). The average property we measure for the sub-samples with double hump Lyα profile is intermediate between LAEs and LBGs. However, it is interesting to understand the connection between the appearance of a double hump feature and low reddening of the quasars. A double hump feature is naturally produced when Lyα transport occurs with nearly static H₁ gas with very little dust and weak bulk motions. Another possible explanation is that systems towards blue quasars having large velocity infall. However, a clear answer to our finding can only come from a detailed analysis of individual systems with double hump detections in front of blue quasars.

Some caution has to be exercised while comparing the DLA measurements with those of LBGs and LAEs. Noterdaeme et al. (2010) have detected nebular emission lines from $z \sim 0.6$ Mg II systems in the SDSS fibre spectra of distant QSOs. From their fig. 7, we can see that the relative velocity between the absorption redshift and the redshift from the nebular line can be between ±100 km s⁻¹. If similar shifts exist between the line of sight where we detect the DLA absorption and the star-forming regions from where the Lyα emission originates, then the intrinsic Lyα emission profile will be smeared by these additional random velocities. While this effect prevents us from establishing a link between DLAs and either LAEs or LBGs, the analysis presented here suggests that the average Lyα profile we find follows the trends seen in the case of LBGs and LAEs.

## 5.2 Fluorescent Lyα emission induced by UV background

We now explore the Lyα emission being induced by metagalactic UV background impinging on the optically thick gas as a possible mechanism for the Lyα emission from DLAs. For this, we use the recent computation of the background ionizing radiation by Khaire & Srianand (2015a,b), accounting for the contributions from both quasars and galaxies (with typical escape fraction, $f_{esc}$, of 2 and 4 per cent for galaxies). For the $f_{esc}$ of 0, 2 and 4 per cent,
<table><caption>**Table 6.** The SFR and surface brightness corresponding to the Lyα luminosity seen in the red and blue parts of DLA profile (median stack) in bootstrap analysis.</caption><thead><tr><th rowspan="2">Sample</th><th colspan="3">SFR (M<sub>☉</sub> yr<sup>-1</sup>)<sup>a</sup></th><th rowspan="2">Surface brightness<sup>b</sup><br>(×10<sup>-19</sup> erg s<sup>-1</sup> cm<sup>-2</sup> arcsec<sup>-2</sup>)</th></tr><tr><th>DLA bottom</th><th>Blue part</th><th>Red part</th></tr></thead><tbody><tr><td>Full with 90 per cent bootstarp</td><td>≤0.55</td><td>≤0.39</td><td>≤0.39</td><td>≤3.93</td></tr><tr><td>Full with 80 per cent bootstarp</td><td>≤0.58</td><td>≤0.41</td><td>≤0.42</td><td>≤4.17</td></tr><tr><td>W<sub>1526</sub> ≥ 0.4 Å</td><td>0.71 ± 0.21</td><td>≤0.45</td><td>0.65 ± 0.13</td><td>6.66 ± 1.56</td></tr><tr><td>W<sub>1526</sub> ≥ 0.8 Å</td><td>≤0.87</td><td>≤0.61</td><td>0.64 ± 0.17</td><td>6.51 ± 2.11</td></tr><tr><td>W<sub>1526</sub> < 0.8 Å</td><td>≤0.85</td><td>≤0.58</td><td>≤0.61</td><td>≤6.58</td></tr><tr><td>N(H1) ≥ 10<sup>21.23</sup> cm<sup>-2</sup></td><td>≤0.78</td><td>≤0.55</td><td>≤0.56</td><td>≤5.51</td></tr><tr><td>10<sup>21</sup> ≤ N(H1) < 10<sup>21.23</sup> cm<sup>-2</sup></td><td>≤0.76</td><td>≤0.53</td><td>≤0.55</td><td>≤5.53</td></tr><tr><td>z<sub>abs</sub> ≥ 2.7</td><td>≤0.84</td><td>≤0.58</td><td>≤0.60</td><td>≤4.79</td></tr><tr><td>z<sub>abs</sub> < 2.7</td><td>≤0.73</td><td>≤0.51</td><td>≤0.52</td><td>≤6.43</td></tr><tr><td>(r - i) ≥ 0.1</td><td>≤0.79</td><td>≤0.56</td><td>≤0.55</td><td>≤4.98</td></tr><tr><td>(r - i) < 0.1</td><td>≤0.78</td><td>≤0.53</td><td>≤0.57</td><td>≤5.26</td></tr></tbody></table>

Notes. <sup>a</sup> A 3σ upper limit is given in case of non-detection. Here, the errors are purely statistical and depend on the value of assumed $f_{esc}$ and IMF. See Section 5.3 for details of assumption involved in this estimate.

<sup>b</sup> Surface brightness corresponding to the Lyα luminosity seen in the red part of DLA bottom.

the integrated unidirectional flux of the H1 ionizing photons (i.e. $\phi_0 = \int_{v_0}^{\infty} \pi I_v d_v / h v$) at $z \sim 2.7$ are found to be $\sim 9.1 \times 10^4$, $1.3 \times 10^5$ and $1.7 \times 10^5$ photons s$^{-1}$ cm$^{-2}$, respectively (see, Shull, Danforth & Tilton 2014). If we assume an arcsec$^2$ optically thick H1 gas cloud in photoionization equilibrium with this background, it will produce a surface brightness of $\sim 3.1 \times 10^{-20}$, $4.6 \times 10^{-20}$ and $6.2 \times 10^{-20}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$. These values will be reduced if the clouds are observed at some inclination angle.

Note that, in the presence of a local ionizing source the observed surface brightness will increase and is related to surface brightness induced by UV background as, $SB = (1 + b)SB_{bgr}$. Here b is the boost factor defined as the ratio of H1 photoionization rate ($\Gamma_{H1}$) from the local source to that of the UV background. Table 6 lists surface brightness corresponding to the Lyα luminosity measured in the red part or 3σ upper limits. It is clear that, in cases where luminosity is measured at $>3\sigma$ level, the measured Lyα surface brightness requires a boost by a factor of more than 6 even if we consider metagalactic UV background computed with 4 per cent escape fraction for UV photons from galaxies. In particular, for systems with $W_{1526} \ge 0.8$ Å, we need a boost in the Lyα continuum flux by at least a factor of $\sim 10$ from local sources of excitation probably from *in situ* star formation or from star-forming regions close to the DLAs. In what follows, we discuss the implication of a local radiation field in some detail.

## 5.3 *In situ* star formation

Assuming that the Lyα photons mainly originate from HII regions around massive stars embedded in the DLAs and case B-recombination (Osterbrock & Ferland 2006), we relate the Lyα luminosity to the SFR ($\dot{M}_{SF}$) as

$$L_{Ly\alpha} = 0.68 h v_{\alpha} (1 - f_{esc}) N_{\gamma} \dot{M}_{SF}. \quad (1)$$

Here, $h v_{\alpha} = 10.2$ eV, is the energy of the Lyα photons. At the redshift of interest in this work $f_{esc} \le 0.04$ based on the $\Gamma_{H1}$ measurement using Lyα absorption by the intergalactic medium (Khaire et al. 2016). Vasei et al. (2016) have found an $f_{esc} \le 0.08$ using deep HST imaging of a gravitationally lensed galaxy at $z \sim 2.38$. In our calculations, we use $f_{esc} = 0.04$. The $N_{\gamma}$ represents the number of ionizing photons released per baryons of star formation. We use $N_{\gamma} = 9870$, corresponding to the average metallicity, i.e. $Z/Z_{\odot} =$

$-1.5$, of high-redshift DLA absorbers and a Salpeter initial mass function with $\alpha = 2.35$, given in Rahmani et al. (2010, references therein). Note that the observed Lyα luminosity depends on the escape fraction ($f_{esc}^{Ly\alpha}$) of Lyα photons and related to the emitted Lyα luminosity ($L_{Lyα}$) as $L_{Lyα}^{obs} = f_{esc}^{Ly\alpha} L_{Lyα}$. We use the $f_{esc}^{Ly\alpha}$ to be 5 per cent, as is estimated for the high-redshift ($z \sim 2.2$) star-forming galaxies by Hayes et al. (2010). Based on the Lyα to the Hα luminosity ratio, Trainor et al. (2015) have recently estimated the $f_{esc}^{Ly\alpha}$ to be 30 per cent for the LAE. For our purpose, we use 5 per cent escape in our analysis.

The SFR corresponding to the Lyα luminosity in the entire DLA bottom as well as in the red and blue parts for various sub-samples, measured in our bootstrap analysis, are given in Table 6. Note, if $f_{esc}^{Ly\alpha}$ is 30 per cent, as seen in the case of LAEs, then the inferred SFR will be less than what is given in Table 6 by a factor of 6. The average SFR in the stacked spectrum of our full sample is found to be $\le 0.4$ M$_{\odot}$ yr$^{-1}$ in the red and blue parts and also in the entire DLA trough. However, the sub-samples with $W_{1526} \ge 0.4$ Å and $W_{1526} \ge 0.8$ Å show a slightly higher SFR at $\sim 0.7$ M$_{\odot}$ yr$^{-1}$ for the total and the red part. In addition, for the sub-samples based on $N(H1)$ and $z_{abs}$, the 3σ upper limit on SFR is found to be $\sim 0.6$ M$_{\odot}$ yr$^{-1}$ (see also Table 6).

Next, we explore the Kennicutt-Schmidt (KS) relation between the gas surface mass density and the SFR per unit surface area in DLA galaxies. Assuming the typical size of the DLA galaxies to be $R = 8$ kpc (the fibre filling factor of 1), we compute the average SFR per unit area to be $\log \Sigma_{SFR} = -2.7$ and $-2.5$ M$_{\odot}$ yr$^{-1}$ kpc$^{-2}$ in the red part of DLA bottom for our entire sample and the sub-sample with $W_{1526} \ge 0.8$ Å, respectively. As the fibre filling factor may be smaller than one, the measured $\Sigma_{SFR}$ values should be a lower limits. For the mean log $N(H1)=21.23$, probed in this work, the expected average SFR is about 0.01 M$_{\odot}$ yr$^{-1}$ kpc$^{-2}$, if one considers the KS relation for low-redshift star-forming galaxies (Kennicutt 1998a,b). The surface SFR densities in DLA galaxies are a factor of $\sim 6$ lower than what one would predict from the local star formation law (see also, Chelouche & Bowen 2010; Rahmani et al. 2010; Rafelski et al. 2011).

In Fig. 10, we compare the gas surface density ($\Sigma_{gas}$) with the surface SFR ($\Sigma_{SFR}$) in different kinds of galaxies with what we measure for DLAs. It is clear from the figure that the DLAs occupy the region that is generally populated by low-metallicity dwarf galaxies and
<figure><img src="image_11.png" /><figcaption>Figure 10. The comparison of surface density of star formation and total (solid symbols) and H₁(open symbols) gas surface density in DLAs with different classes of star-forming galaxies. Orange points (open pentagon) represent low surface brightness galaxies (Wyder et al. 2009), black points (small hexagon) spatially resolved observations of Small Magellanic Cloud (Bolatto et al. 2011), purple points (circle) extremely metal-poor galaxies (Filho et al. 2016), green points (triangle) dwarf galaxies (Cormier et al. 2014), brown points (square) resolved observations of extremely metal-poor galaxies SextanA and (circle) ESO 146-G14 (Shi et al. 2014). The dashed line corresponds to Kennicutt–Schmidt law (Kennicutt & Evans 2012). The hatched region shows the surface brightness seen in the outskirts (i.e. 3–6 kpc) of LBGs (Rafelski et al. 2011).</figcaption></figure>

extremely metal-poor (i.e. $Z < 0.1 Z_{\odot}$) galaxies (Wyder et al. 2009; Bolatto et al. 2011; Cormier et al. 2014; Shi et al. 2014; Filho et al. 2016). In this figure, we also compare the surface brightness seen in the outer region of LBGs in the stacked images by Rafelski et al. (2011). Our measurements and the limits are consistent with this range. It is interesting to note that in the case of DLAs despite log $N(H_1) \ge 21$ we do not detect H₂ molecules in our stacked spectrum (see, Noterdaeme, Petitjean & Srianand 2015, for H₂-H₁ transition in DLAs using individual H₂ measurements) so the surface mass density is mainly atomic. However, even in the case of dwarf galaxies a considerable contribution to the measured surface mass density comes from molecular gas. While H₁ emission is widespread in the dwarf galaxies, regions with log $N(H_1) \ge 21$ project smaller cross-section and usually localized to the star-forming regions even if there is no one to one correspondence between SFR and $N(H_1)$ (Begum et al. 2006; Roychowdhury et al. 2014). Therefore, in this scenario the continuum and Ly$\alpha$ emission associated with the H₁ gas can be well within the SDSS fibre.

The low star formation efficiency in low-metallicity dwarf galaxies and extremely metal-poor galaxies is mainly attributed to the (i) low cooling rate in metal-poor environments (Filho et al. 2013, 2016); (ii) galactic winds which can suppress the SFR by throwing more than 80 per cent of gas back to the circumgalactic medium (Sánchez Almeida et al. 2014) and/or (iii) related to the star formation in compact regions which are fed by the accretion of metal-poor H₁ gas by cold cosmological accretion (Ekta & Chenzalur 2010; Sánchez Almeida et al. 2014). The same may be the

reason for low SFR seen in DLA galaxies as well, as the DLAs are also found to have a low average metallicity, i.e. ~1/10th of solar, at $z \sim 2$ (Rafelski et al. 2012; Noterdaeme et al. 2014) and signatures of outflows as an off-centred Ly$\alpha$ emission in our stack.

The low surface brightness galaxies or the low surface brightness outer regions of luminous galaxies contributing to the high-z DLA population is a good possibility, even though at low-z such galaxies do not contribute appreciably to the high $N(H_1)$ absorbers studied here (Patra et al. 2013). While we expect the number density of low-mass galaxies to be much higher at higher-z, Ly$\alpha$ forest–DLA cross-correlation analysis suggests that most of the high-z DLAs may originate from a biased regions with haloes of $10^{12} M_{\odot}$ (Font-Ribera et al. 2012).

Next, we consider the possibility of DLAs originating from passive H₁ gas at the external regions of massive star-forming galaxies. In this scenario, the Ly$\alpha$ emission will be detected either when the galaxy light falls inside the fibre (i.e. impact parameter less than 8 kpc) or when the Lyman continuum photons from the host galaxy inducing the Ly$\alpha$ fluorescence emission from the extended gas (i.e. the so-called Ly$\alpha$ halo).

## 5.4 Star-forming galaxies at small impact parameter

The number of DLAs at $z \sim 2.7$ having log $N(H_1) \ge 21.0$ per unit absorption distance is estimated to be $dN/dX \approx 0.028$, by integrating over the column density distribution function $f_{H_1}(N, X)$ from Noterdaeme et al. (2009). This can be used to derive the comoving incidence of star-forming galaxies at high redshift (LBGs) as

$$ \frac{dN_{\text{LBGs}}(L > L_{\text{min}})}{dl_c} = (1+z)^2 f \int_{L_{\text{min}}}^{\infty} \pi R(L)^2 \phi(L) dL. \quad (2) $$

Here, $R(L)$ represents size (or impact factor) of the H₁ gas below which the required $N(H_1)$ will be produced in a galaxy with luminosity $L$. We use a power-law relation, i.e. $R(L) \propto L^{0.4}$, obtained from H₁ size–optical luminosity relation from field galaxies in Arecibo Legacy Fast Arecibo L-Band Feed Array survey which is determined based on the galaxy R-band luminosity (Toribio et al. 2011). The f-factor is the gas filling factor which is considered to be one. We use the UV luminosity function of LBGs by Reddy & Steidel (2009) at redshifts $1.9 \le z \le 2.7$, with the best-fitting Schechter function parameter of $\phi^* = 2.75 \times 10^{-3} \text{Mpc}^{-3}$, $M_{\text{AB}}^{*}(1700 \text{Å}) = -20.70$ and $\alpha = -1.73$. Using the above scaling relation, $R(L) = R_{*}(L/L_{*})^{0.4}$, equation(2) becomes

$$ \frac{dN_{\text{LBGs}}}{dl_c} = \pi R_*^2 (1+z)^2 f \int_{L_{\text{min}}}^{\infty} \left(\frac{\phi^*}{L^*}\right) \left(\frac{L}{L^*}\right)^{(\alpha+0.8)} \times \exp\left(-\frac{L}{L^*}\right) dL. \quad (3) $$

Owing to the large magnifications by a lensing cluster, Alavi et al. (2014) have probed the LBGs down to very faint magnitudes, $M_{1500} < -13$ mag, about 100 times fainter than previous studies at the same redshift (Reddy & Steidel 2009). They have shown that the UV luminosity function of LBGs shows no turnover up to $M_{1500} < -13$ mag and have a faint-end slope, $\alpha = 1.74$, similar to Reddy & Steidel (2009). We compute the impact factor, $R \sim 13$ kpc, below which required $N(H_1)$ will be produced by an $L^*$ galaxy by integrating the above equation down to $L_{\text{min}} = 0.001 L^*$. It implies a high probability for the $L^*$ galaxy to lie outside the small region of 8 kpc probed by the BOSS fibre. Alternatively, taking the maximum value of $R$ to be ~8 kpc (corresponding to the fibre size) in above scaling relation $R = R_{*}(L/L^*)^{\alpha}$, we find the maximum luminosity
of a galaxy which will come inside the fibre to be 0.3 L*. In other words, the majority of galaxies contributing directly to the average luminosity seen in stacked spectrum will be below ~0.3 L*.

The average luminosity of LBGs with luminosity ranging from 0.001 to 0.3 L* is given as

$$ \langle L_{\text{LBG}} \rangle = \frac{\int_{L_{\text{min}}}^{L_{\text{max}}} L \phi(L) dL}{\int_{L_{\text{min}}}^{\infty} \phi(L) dL}. \qquad (4) $$

We find the average UV luminosity of LBGs to be 1.34 × 10^42 erg s⁻¹, corresponding to an SFR of 0.12 M☉ yr⁻¹. Interestingly, this is comparable to the average SFR of 0.09 < SFR < 0.27 M☉ yr⁻¹ (2σ limit) seen in direct imaging of DLA host galaxies within 2 and 12 kpc aperture sizes by Fumagalli et al. (2015), albeit for DLAs with lower N(HI) than what we consider here. Note that if we assume the gas filling factor to be less than 1, then R* will be > 13 kpc and Lmax < 0.3 L*, and the average UV luminosity will be even lower than the above quoted value. The expected SFR within 12 kpc aperture, using the above formalism, is found to be 0.14 M☉ yr⁻¹. Therefore, allowing for the impact parameter for a given N(HI) to scale with optical luminosity (as seen in the local Universe) for high-z DLAs gives the mean SFR in the correct ballpark value. Converting the average UV luminosity estimates into Lyα luminosity is not straightforward as the fraction of Lyα emitting galaxies may be a function of optical luminosity. However, one can get a simple estimate by substituting our constraint on SFR in equation (1). If we use fesc = 0.04 then in order to explain the Lyα luminosity we find for high W1526 sub-samples we need fLyα esc = 0.22. This is in between the values found for LBGs (i.e. 5 per cent) and LAEs (i.e. 30 per cent).

In the above discussion, we found that most of the Lyα emission from bright galaxies associated with a DLA will be outside the SDSS fibre. However, even in this case Lyα fluorescence from the large extended regions can still contribute to the Lyα emission inside the fibre. Extended diffuse Lyα emitting regions (up to 80 kpc radius) are detected around high-redshift LAEs and LBGs in recent observations by stacking analysis (Hayashino et al. 2004; Steidel et al. 2011; Matsuda et al. 2012; Feldmeier et al. 2013; Momose et al. 2014). The presence of these diffuse Lyα emitting haloes (LAHs) is also predicted by numerical simulations (Laursen & Sommer-Larsen 2007; Zheng et al. 2011; Dijkstra & Kramer 2012; Jeeson-Daniel et al. 2012; Verhamme et al. 2012).

Momose et al. (2016) have found prominent LAHs around LAEs with faint Lyα luminosities, bright UV luminosities and small Lyα equivalent width. The extent of the LAHs is found to be dependent on the environment, where LAHs in low-density environment are found to be smaller than the higher density environments having a typical scalelength of ~9.1 and 20.4 kpc, respectively (Matsuda et al. 2012; Momose et al. 2014). In addition, in continuum-subtracted Lyα emission images of LBGs, Steidel et al. (2011) have found similar diffuse Lyα haloes for all the galaxies, irrespective of whether the Lyα is seen in emission or absorption, up to radii of ~80 kpc. In Fig. 11, we compare the average Lyα surface brightness profile found by Steidel et al. (2011) around the LBGs with our measurements. Our discussions in the previous section suggest that an L* galaxy will have a typical impact parameter of ~13 kpc. It is clear from the figure that surface brightness contributions at these impact parameters will be much higher than the average we find for our samples. Inclusion of this contribution will increase the Lyα luminosity inside the fibre from stars that do not contribute to the continuum stacked image. Therefore, the fLyα esc can be smaller than 0.22 which we inferred to match our Lyα luminosities with the SFR constrains from the stacking experiment.

<figure><img src="image_12.png" /><figcaption>Figure 11. The average surface brightness profile of the Lyα emission in stacked images of 92 LBGs (solid black) and in a sub-sample with LBGs seen in Lyα emission (orange dashed) and in Lyα absorption (green dot-dashed) by Steidel et al. (2011). The grey lines represent the average surface brightness profile of continuum light (at 1220 Å) for all LBGs (dash-dotted) as well as the sub-samples with only Lyα emission (dotted) and absorption (dashed). The hatched region shows the surface brightness seen in our DLA stacks for entire sample (blue) and for a sub-sample (red) with W1526 ≥ 0.4 Å, respectively. The arrow represents the expected R* for unit covering factor of HI gas around galaxies from discussion in Section 5.4.</figcaption></figure>

# 6 SUMMARY

In order to probe the star formation in neutral gas clouds at high redshift, we perform a stacking analysis of 704 DLAs with a large H1 column density (log N(HI) ≥ 21) and a median redshift of ~2.7. We generate both emission and absorption line stacked spectra for various sub-samples based on N(HI), zabs, W1526 and (r - i) colours of QSOs which led us to the following conclusions.

(1) For the full sample, we measure the Lyα luminosity of (5.2 ± 3.3) × 10^40 erg s⁻¹ when we integrate the luminosity over the full core regions of the DLA in the median spectrum. Similar values are obtained when we considered the stacked spectra obtained using weighted mean and 3σ clipped weighted mean. This luminosity is ≤0.1 per cent of the Lyα luminosity of L* galaxies at these redshifts. In the bootstrap analysis, the measured luminosities in the red and blue parts of the core regions of the DLA troughs in the median stacked spectrum are (−1.1 ± 2.4) × 10^40 erg s⁻¹ and (6.9 ± 2.5) × 10^40 erg s⁻¹, respectively, with an ~2.8 σ excess Lyα emission in the red part of the DLA trough.

(2) In the sub-samples based on W1526, we find the Lyα luminosities of (13.5 ± 4.1) × 10^40 erg s⁻¹ and (14.6 ± 5.5) × 10^40 erg s⁻¹ when we integrate the luminosity over the core of Lyα trough in the median spectrum for W1526 ≥ 0.4 and 0.8 Å, respectively. These luminosities are mainly contributed from the red part with respective luminosities of Ls* = (12.3 ± 2.9) × 10^40 erg s⁻¹ and (12.1 ± 3.9) × 10^40 erg s⁻¹. Blue parts have luminosities consistent with zero. For systems with Si II detections having W1526 < 0.8 Å, we find the Lyα luminosity to be (0.6 ± 4.5) × 10^40 erg s⁻¹. Even in this case most of the observed luminosity comes from the
red part albeit with a significance of 1.3σ level. As the median $N(H_1)$ in all these sub-samples are the same, the difference in the $Lyα$ luminosity is not likely related to the differences in the $H_1$ optical depth. Prochaska et al. (2008) found a strong correlation between $W_{1526}$ and metallicity using echelle spectroscopic data. This correlation has been interpreted as a mass-metallicity relation as high $W_{1526}$ absorption tends to trace low optical depth clouds in the halo or outflowing gas. Thus, the enhanced $Lyα$ luminosity in the high $W_{1526}$ could be an indication of high star formation in high-metallicity systems together with an easier escape of $Lyα$ photos enabled by the large outflowing gas. We measure large $b_{eff}$ for systems with high $W_{1526}$, which supports this hypothesis.

(3) The sub-samples based on $N(H_1)$ do not show any detectable difference in the measured $Lyα$ luminosities either in the full core region or in the red part alone. The same is the case when we divided the sample into two redshift bins. We note that all the four sub-samples have nearly similar mean $W_{1526}$ (see Table A3).

(4) The sub-samples based of $(r-i)$ colours and $W_{1526} > 0.4$ Å show a double hump profile for the low $(r-i)$ sub-sample. The double hump disappears when we add more and more red QSOs sightlines to this sample. However, total $Lyα$ luminosity does not show any monotonous trend with $(r-i)$. In addition, the $L_{\lambda}^{r}/L_{\lambda}^{b}$, FWHM and peak separation are intermediate between what is seen in LBGs and LAEs. However, to uniquely identify the reason behind the appearance of double-humped $Lyα$ line towards blue QSOs, we need detailed analysis of such individual systems using radiative transfer models. Establishing any trend between the dust indicators and $Lyα$ profile will help us to discriminate between different models, such as (i) static medium (Neufeld 1990; Zheng & Miralda-Escudé 2002; Dijkstra et al. 2006), (ii) expanding/inflowing shell (Ahn, Lee & Lee 2003; Verhamme et al. 2006; Scherer et al. 2011) or (iii) multiphase media (Gronke & Dijkstra 2016), of the $Lyα$ radiative transport in DLAs. As these models predict different shape as a function of dust content.

(5) In our full sample as well as in all the sub-samples, we detected the $Lyα$ emission predominantly in the red part of the $Lyα$ trough albeit with varied significance. The measured shift in the peak location of $Lyα$ emission with respect to the absorption redshift is in the range 300–400 km s⁻¹. The redshifted profiles are typical of what is seen in the case of LBGs and LAEs at the same redshifts. In these cases, the observed redshifted $Lyα$ profiles are considered as a signature of predominantly outflowing gas in these galaxies. Thus, the $Lyα$ profile we measure in the case of high-z DLAs is consistent with the presence of outflowing gas. The measured shift in DLAs is higher than what is typically seen in LAEs but less than those of LBGs. However, as our reference redshift is the absorbing gas along the QSO sightline and not the non-resonant nebular emission from the galaxy itself, we need to exercise caution in interpreting this result.

(6) Using the updated metagalactic UV background radiation contributed by QSOs and galaxies, we find the expected $Lyα$ fluorescence is 4–10 times less than what we measure. This means most of the DLAs are not passive clouds in ionization equilibrium with the metagalactic UV background. Local excess ionizing radiation either from a nearby star-forming region or from *in situ* star formation is needed to produce the observed $Lyα$ luminosities. We discuss different scenarios such as *in situ* star formation in low-luminosity galaxies or scattered $Lyα$ emission from extended $Lyα$ haloes around high-luminosity galaxies as a possible alternative to explain the observed $Lyα$ luminosity. If the bias estimated for the DLA population based on clustering analysis as well as our mass

estimation from the stacked absorption lines are true then the second alternative will be favoured.

## ACKNOWLEDGEMENTS

We thank the anonymous referee for constructive comments. We are grateful to Vikram Khaire for providing the metagalactic UV background. RS, PN and PPJ acknowledge the support from Indo-French Centre for the Promotion of Advance Research (IFCPAR) under project number 5504-2.

Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the US Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society and the Higher Education Funding Council for England. The SDSS website is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory and the University of Washington.

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APPENDIX A

A1 Dependence on QSOs colour

Here, we present the dependence of Lyα emission on background
QSOs colour by considering the CNR ≥ 4 systems. The results

are shown in Fig. A3 and measured luminosities are summarized
in Table A2 in the appendix. The double hump profile is clearly
visible even in this case that has ~30 per cent more systems. How-
ever, the total luminosity measured in the median stacked spec-
trum, (19.1 ± 6.6) × 10^40 erg s^−1, is slightly less but consistent
within 0.4σ level to what we measure for CNR ≥ 5. The dou-
ble Gaussian fits are shown in Fig. A3. It is also clear from Ta-
ble A2 that derived luminosities using Gaussian fits to the blue
and red humps are significant at more than 3.3 σ level. Here also
we find L_λ^r being higher than L_λ^b. However, the excess is less than
what we have seen in the CNR ≥ 5 sample. The measured ve-
locity shifts of Gaussian peaks with respect to the central wave-
length systemic redshift are consistent with that measured for the
CNR ≥ 5 case. This enlarged sample also shows: (i) the pres-
ence of enhanced emission in the blue part; (ii) the blue peak hav-
ing more velocity shift with respect to the systemic redshift com-
pared to the red and (iii) L_λ^r being higher than L_λ^b for systems with
(r − i) < 0.05.

In addition, the trend of total luminosity being nearly same, and
the double hump present only in low (r − i) sub-samples, as seen
in Fig. 6, is also present when we consider systems with CNR ≥ 4.
However, the average luminosities are slightly less from what we
find for CNR ≥ 5 sample.

<figure><img src="image_13.png" /><figcaption>Figure A1. Examples of DLAs from our parent sample with large pixel-to-pixel variation in the absorption trough (see also, Section 3.1).</figcaption></figure>
<figure><img src="image_14.png" /><figcaption>**Figure A2.** Left to right: comparison of median stacked spectrum for a sub-sample with CNR ≥ 5 and $W_{1526}$ greater than (top panel) and less than (bottom panels) 0.4, 0.6 and 0.8 Å, respectively. The blue segment shows the DLA core with $\tau \ge 10$ for $\log N(H_i) = 21.0$. The dashed curves show the synthetic profiles for lower (i.e. $\log N(H_i) = 21$) and median (i.e. $\log N(H_i) = 21.23$) column density of DLAs used to get the stacked spectrum.</figcaption></figure>

<figure><img src="image_15.png" /><figcaption>**Figure A3.** Same as Fig. 6, for the sub-sample with relaxed CNR limit to CNR ≥ 4.</figcaption></figure>
<figure><img src="image_16.png" /><figcaption>**Figure A4.** Comparison of several absorption lines in the composite spectra of a number of sub-samples, defined in Table 2. Each row shows different lines in the composite spectra of the same two sub-samples (one in red and the other in black), which are detailed in the last column of each row. Each column contains lines of the same species, indicated in the top row of that column, the wavelength of which is indicated below each column.</figcaption></figure>

<table><caption>**Table A1.** Luminosity of Ly$\alpha$ line in bootstrap analysis of median stacked spectra for various sub-samples based on $W_{1526}$ and CNR limit of $\ge 5$ and $\ge 4$.</caption><thead><tr><th rowspan="2">Sample criteria</th><th rowspan="2">Number</th><th colspan="3">Ly&alpha; luminosity (&times;10<sup>40</sup> erg s<sup>-1</sup>)</th></tr><tr><th>DLA bottom</th><th>Median stack<br>L<sup>b</sup><sub>&lambda;</sub></th><th>L<sup>r</sup><sub>&lambda;</sub></th></tr></thead><tbody><tr><td>W<sub>1526</sub> &ge; 0.4 &Aring; and CNR &ge; 5</td><td>513</td><td>13.49 &plusmn; 4.05</td><td>1.16 &plusmn; 2.84</td><td>12.34 &plusmn; 2.90</td></tr><tr><td>W<sub>1526</sub> &lt; 0.4 &Aring; and CNR &ge; 5</td><td>112</td><td>-26.27 &plusmn; 8.33</td><td>-10.78 &plusmn; 5.92</td><td>-15.45 &plusmn; 5.85</td></tr><tr><td>W<sub>1526</sub> &ge; 0.6 &Aring; and CNR &ge; 5</td><td>391</td><td>13.76 &plusmn; 4.69</td><td>1.70 &plusmn; 3.00</td><td>12.05 &plusmn; 3.33</td></tr><tr><td>W<sub>1526</sub> &lt; 0.6 &Aring; and CNR &ge; 5</td><td>279</td><td>3.41 &plusmn; 5.44</td><td>-1.43 &plusmn; 3.76</td><td>4.85 &plusmn; 3.93</td></tr><tr><td>W<sub>1526</sub> &ge; 0.4 &Aring; and CNR &ge; 4</td><td>627</td><td>11.05 &plusmn; 3.67</td><td>2.02 &plusmn; 1.65</td><td>11.48 &plusmn; 1.75</td></tr><tr><td>W<sub>1526</sub> &lt; 0.4 &Aring; and CNR &ge; 4</td><td>119</td><td>-28.18 &plusmn; 8.09</td><td>-11.60 &plusmn; 5.78</td><td>-16.59 &plusmn; 5.66</td></tr><tr><td>W<sub>1526</sub> &ge; 0.6 &Aring; and CNR &ge; 4</td><td>496</td><td>9.92 &plusmn; 4.18</td><td>-0.43 &plusmn; 2.93</td><td>10.34 &plusmn; 2.98</td></tr><tr><td>W<sub>1526</sub> &lt; 0.6 &Aring; and CNR &ge; 4</td><td>320</td><td>1.87 &plusmn; 5.00</td><td>-2.90 &plusmn; 3.47</td><td>4.77 &plusmn; 3.60</td></tr><tr><td>W<sub>1526</sub> &ge; 0.8 &Aring; and CNR &ge; 4</td><td>370</td><td>8.72 &plusmn; 4.78</td><td>-0.80 &plusmn; 3.36</td><td>9.52 &plusmn; 3.40</td></tr><tr><td>W<sub>1526</sub> &lt; 0.8 &Aring; and CNR &ge; 4</td><td>511</td><td>-0.04 &plusmn; 4.03</td><td>-4.07 &plusmn; 2.80</td><td>4.03 &plusmn; 2.91</td></tr></tbody></table>

<table><caption>**Table A2.** Same as Table 4, after relaxing the CNR limit to CNR $\ge 4$.</caption><thead><tr><th rowspan="2">Sample%</th><th rowspan="2">(r - i)</th><th rowspan="2">&lang;&Delta;(r - i)&rang;</th><th rowspan="2">DLA-bottom</th><th rowspan="2">Ly&alpha; luminosity<br>(&times;10<sup>40</sup> erg s<sup>-1</sup>)<br>L<sup>b</sup><sub>&lambda;</sub></th><th colspan="3">Velocity offset</th><th rowspan="2">FWHM<br>(km s<sup>-1</sup>)</th></tr><tr><th>L<sup>r</sup><sub>&lambda;</sub></th><th>&Delta;&nu;<sup>r</sup></th><th>&Delta;&nu;<sup>b</sup></th></tr></thead><tbody><tr><td>Lower 30%</td><td>&lt;0.04</td><td>-0.08</td><td>19.07 &plusmn; 6.61</td><td>7.65 &plusmn; 4.42</td><td>11.41 &plusmn; 4.92</td><td>283 &plusmn; 24</td><td>-401 &plusmn; 19</td><td>-</td></tr><tr><td>Lower 40%</td><td>&lt;0.07</td><td>-0.06</td><td>13.07 &plusmn; 5.84</td><td>3.92 &plusmn; 3.94</td><td>9.15 &plusmn; 4.31</td><td>301 &plusmn; 22</td><td>-383 &plusmn; 24</td><td>-</td></tr><tr><td>Lower 50%</td><td>&lt;0.10</td><td>-0.04</td><td>10.47 &plusmn; 5.24</td><td>0.95 &plusmn; 3.51</td><td>9.52 &plusmn; 3.89</td><td>263 &plusmn; 48</td><td>-336 &plusmn; 54</td><td>205 &plusmn; 112</td></tr><tr><td>Upper 50%</td><td>&gt;0.10</td><td>0.12</td><td>10.87 &plusmn; 5.10</td><td>-0.65 &plusmn; 3.66</td><td>11.52 &plusmn; 3.56</td><td>-</td><td>-</td><td>-</td></tr><tr><td>Upper 40%</td><td>&gt;0.13</td><td>0.13</td><td>13.27 &plusmn; 5.64</td><td>-1.10 &plusmn; 4.05</td><td>14.37 &plusmn; 3.93</td><td>-</td><td>-</td><td>-</td></tr><tr><td>Upper 30%</td><td>&gt;0.16</td><td>0.16</td><td>10.77 &plusmn; 6.69</td><td>-0.38 &plusmn; 4.93</td><td>11.15 &plusmn; 4.52</td><td>-</td><td>-</td><td>-</td></tr></tbody></table>
<table><caption>Table A3. Equivalent widths of metal absorption lines for various sub-samples.</caption><thead><tr><th rowspan="2">Transition</th><th colspan="2">All</th><th colspan="2">W<sub>1526</sub></th><th colspan="4">W<sub>r</sub>(Å)</th><th colspan="2">log N(H<sub>1</sub>)</th><th colspan="2">z<sub>abs</sub></th></tr><tr><th>(S1)</th><th>≥0.8 Å<br>(S2)</th><th>(S3)</th><th>≥0.4 Å<br>(S4)</th><th>≥21.23<br>(S5)</th><th>≥21.23<br>(S6)</th><th>≥2.7<br>(S7)</th><th>≥2.7<br>(S8)</th><th>(S1)</th><th>≥0.8 Å<br>(S2)</th><th>≥0.4 Å<br>(S4)</th><th>≥21.23<br>(S5)</th><th>≥2.7<br>(S7)</th></tr></thead><tbody><tr><td>Al IIλ1670</td><td>0.732 ± 0.008</td><td>1.218 ± 0.012</td><td>0.483 ± 0.008</td><td>0.864 ± 0.009</td><td>0.841 ± 0.012</td><td>0.589 ± 0.011</td><td>0.713 ± 0.011</td><td>0.735 ± 0.012</td><td>0.713 ± 0.008</td><td>0.589 ± 0.011</td><td>0.713 ± 0.008</td><td>0.735 ± 0.008</td></tr><tr><td>Al IIIλ1854</td><td>0.171 ± 0.006</td><td>0.332 ± 0.009</td><td>0.092 ± 0.007</td><td>0.219 ± 0.007</td><td>0.200 ± 0.008</td><td>0.153 ± 0.008</td><td>0.161 ± 0.009</td><td>0.175 ± 0.008</td><td>0.161 ± 0.008</td><td>0.153 ± 0.008</td><td>0.161 ± 0.008</td><td>0.175 ± 0.008</td></tr><tr><td>Al IIIλ1862</td><td>0.103 ± 0.006</td><td>0.192 ± 0.009</td><td>0.050 ± 0.007</td><td>0.135 ± 0.006</td><td>0.134 ± 0.008</td><td>0.079 ± 0.008</td><td>0.084 ± 0.009</td><td>0.105 ± 0.007</td><td>0.084 ± 0.008</td><td>0.079 ± 0.008</td><td>0.084 ± 0.008</td><td>0.105 ± 0.007</td></tr><tr><td>Co IIλ1466</td><td>0.009 ± 0.003</td><td>0.022 ± 0.004</td><td>0.002 ± 0.004</td><td>0.012 ± 0.003</td><td>0.007 ± 0.004</td><td>0.002 ± 0.004</td><td>0.011 ± 0.005</td><td>0.008 ± 0.003</td><td>0.011 ± 0.005</td><td>0.002 ± 0.004</td><td>0.011 ± 0.005</td><td>0.008 ± 0.003</td></tr><tr><td>Cr IIλ2056</td><td>0.085 ± 0.007</td><td>0.142 ± 0.010</td><td>0.060 ± 0.009</td><td>0.077 ± 0.007</td><td>0.087 ± 0.009</td><td>0.038 ± 0.010</td><td>0.056 ± 0.012</td><td>0.093 ± 0.007</td><td>0.056 ± 0.012</td><td>0.038 ± 0.010</td><td>0.056 ± 0.012</td><td>0.093 ± 0.007</td></tr><tr><td>Cr IIλ2066</td><td>0.043 ± 0.006</td><td>0.089 ± 0.009</td><td>0.048 ± 0.008</td><td>0.047 ± 0.007</td><td>0.032 ± 0.008</td><td>0.039 ± 0.009</td><td>0.043 ± 0.010</td><td>0.054 ± 0.007</td><td>0.043 ± 0.010</td><td>0.039 ± 0.009</td><td>0.043 ± 0.010</td><td>0.054 ± 0.007</td></tr><tr><td>Cu λ1334</td><td>0.934 ± 0.010</td><td>1.513 ± 0.015</td><td>0.644 ± 0.011</td><td>1.167 ± 0.012</td><td>1.094 ± 0.015</td><td>0.812 ± 0.014</td><td>0.866 ± 0.015</td><td>0.977 ± 0.014</td><td>0.866 ± 0.015</td><td>1.094 ± 0.014</td><td>0.866 ± 0.015</td><td>0.977 ± 0.014</td></tr><tr><td>Cu λ1548</td><td>0.362 ± 0.007</td><td>0.676 ± 0.012</td><td>0.253 ± 0.007</td><td>0.508 ± 0.008</td><td>0.399 ± 0.009</td><td>0.367 ± 0.010</td><td>0.321 ± 0.010</td><td>0.420 ± 0.009</td><td>0.321 ± 0.010</td><td>0.367 ± 0.010</td><td>0.321 ± 0.010</td><td>0.420 ± 0.009</td></tr><tr><td>Cu λ1550</td><td>0.228 ± 0.006</td><td>0.403 ± 0.011</td><td>0.171 ± 0.007</td><td>0.305 ± 0.007</td><td>0.247 ± 0.008</td><td>0.234 ± 0.009</td><td>0.175 ± 0.009</td><td>0.264 ± 0.008</td><td>0.175 ± 0.009</td><td>0.234 ± 0.009</td><td>0.175 ± 0.009</td><td>0.264 ± 0.008</td></tr><tr><td>Fe IIλ1608</td><td>0.456 ± 0.006</td><td>0.750 ± 0.010</td><td>0.293 ± 0.007</td><td>0.569 ± 0.007</td><td>0.534 ± 0.009</td><td>0.411 ± 0.008</td><td>0.432 ± 0.009</td><td>0.494 ± 0.009</td><td>0.432 ± 0.009</td><td>0.411 ± 0.008</td><td>0.432 ± 0.009</td><td>0.494 ± 0.009</td></tr><tr><td>Fe IIλ1611</td><td>0.048 ± 0.004</td><td>0.070 ± 0.005</td><td>0.028 ± 0.005</td><td>0.056 ± 0.004</td><td>0.078 ± 0.005</td><td>0.025 ± 0.005</td><td>0.040 ± 0.005</td><td>0.050 ± 0.005</td><td>0.040 ± 0.005</td><td>0.025 ± 0.005</td><td>0.040 ± 0.005</td><td>0.050 ± 0.005</td></tr><tr><td>Fe IIλ2249</td><td>0.072 ± 0.010</td><td>0.164 ± 0.015</td><td>0.062 ± 0.014</td><td>0.105 ± 0.011</td><td>0.158 ± 0.014</td><td>0.053 ± 0.015</td><td>0.094 ± 0.018</td><td>0.102 ± 0.011</td><td>0.053 ± 0.015</td><td>0.158 ± 0.014</td><td>0.094 ± 0.018</td><td>0.102 ± 0.011</td></tr><tr><td>Fe IIλ2260</td><td>0.089 ± 0.009</td><td>0.156 ± 0.014</td><td>0.078 ± 0.012</td><td>0.110 ± 0.010</td><td>0.132 ± 0.013</td><td>0.080 ± 0.013</td><td>0.076 ± 0.017</td><td>0.124 ± 0.010</td><td>0.076 ± 0.017</td><td>0.132 ± 0.013</td><td>0.076 ± 0.017</td><td>0.124 ± 0.010</td></tr><tr><td>Fe IIλ2344</td><td>0.894 ± 0.023</td><td>1.471 ± 0.029</td><td>0.578 ± 0.030</td><td>1.152 ± 0.022</td><td>1.083 ± 0.031</td><td>0.819 ± 0.030</td><td>0.886 ± 0.045</td><td>0.891 ± 0.019</td><td>0.886 ± 0.045</td><td>1.083 ± 0.031</td><td>0.886 ± 0.045</td><td>0.891 ± 0.019</td></tr><tr><td>Fe IIλ2374</td><td>0.605 ± 0.020</td><td>0.988 ± 0.028</td><td>0.427 ± 0.026</td><td>0.704 ± 0.022</td><td>0.744 ± 0.029</td><td>0.477 ± 0.027</td><td>0.565 ± 0.041</td><td>0.622 ± 0.018</td><td>0.565 ± 0.041</td><td>0.477 ± 0.027</td><td>0.565 ± 0.041</td><td>0.622 ± 0.018</td></tr><tr><td>Fe IIλ2382</td><td>1.221 ± 0.021</td><td>1.902 ± 0.030</td><td>0.894 ± 0.026</td><td>1.384 ± 0.024</td><td>1.329 ± 0.030</td><td>1.075 ± 0.030</td><td>1.142 ± 0.043</td><td>1.223 ± 0.020</td><td>1.142 ± 0.043</td><td>1.329 ± 0.030</td><td>1.142 ± 0.043</td><td>1.223 ± 0.020</td></tr><tr><td>Fe IIλ2586</td><td>0.791 ± 0.041</td><td>1.363 ± 0.061</td><td>0.532 ± 0.047</td><td>1.038 ± 0.043</td><td>0.913 ± 0.053</td><td>0.706 ± 0.058</td><td>0.587 ± 0.120</td><td>0.856 ± 0.030</td><td>0.587 ± 0.120</td><td>0.913 ± 0.053</td><td>0.587 ± 0.058</td><td>0.856 ± 0.030</td></tr><tr><td>Fe IIλ2600</td><td>1.186 ± 0.042</td><td>1.895 ± 0.059</td><td>0.778 ± 0.051</td><td>1.448 ± 0.042</td><td>1.353 ± 0.053</td><td>1.069 ± 0.062</td><td>0.998 ± 0.127</td><td>1.265 ± 0.032</td><td>0.998 ± 0.127</td><td>1.353 ± 0.053</td><td>0.998 ± 0.127</td><td>1.265 ± 0.032</td></tr><tr><td>Mg IIλ2796</td><td>1.901 ± 0.086</td><td>2.954 ± 0.111</td><td>1.267 ± 0.116</td><td>2.184 ± 0.084</td><td>2.188 ± 0.098</td><td>1.689 ± 0.134</td><td>– ± –</td><td>1.904 ± 0.084</td><td>– ± –</td><td>2.188 ± 0.098</td><td>1.689 ± 0.134</td><td>– ± –</td></tr><tr><td>Mg IIλ2803</td><td>1.813 ± 0.105</td><td>2.858 ± 0.140</td><td>1.123 ± 0.104</td><td>2.107 ± 0.094</td><td>2.010 ± 0.097</td><td>1.716 ± 0.150</td><td>– ± –</td><td>1.828 ± 0.105</td><td>– ± –</td><td>2.010 ± 0.097</td><td>1.716 ± 0.150</td><td>– ± –</td></tr><tr><td>Mg IIλ2852</td><td>0.418 ± 0.075</td><td>0.712 ± 0.100</td><td>0.254 ± 0.112</td><td>0.586 ± 0.083</td><td>0.560 ± 0.103</td><td>0.326 ± 0.105</td><td>– ± –</td><td>0.416 ± 0.076</td><td>– ± –</td><td>0.560 ± 0.103</td><td>0.326 ± 0.105</td><td>– ± –</td></tr><tr><td>Mn IIλ2576</td><td>0.142 ± 0.033</td><td>0.169 ± 0.054</td><td>0.121 ± 0.032</td><td>0.157 ± 0.037</td><td>0.176 ± 0.029</td><td>0.107 ± 0.062</td><td>0.064 ± 0.103</td><td>0.155 ± 0.018</td><td>0.064 ± 0.103</td><td>0.176 ± 0.029</td><td>0.107 ± 0.062</td><td>0.155 ± 0.018</td></tr><tr><td>Mn IIλ2594</td><td>0.057 ± 0.031</td><td>0.105 ± 0.052</td><td>0.049 ± 0.034</td><td>0.115 ± 0.035</td><td>0.108 ± 0.039</td><td>0.037 ± 0.042</td><td>0.010 ± 0.103</td><td>0.093 ± 0.018</td><td>0.010 ± 0.103</td><td>0.108 ± 0.039</td><td>0.037 ± 0.042</td><td>0.093 ± 0.018</td></tr><tr><td>Mn IIλ2606</td><td>0.032 ± 0.026</td><td>0.083 ± 0.039</td><td>0.010 ± 0.033</td><td>0.061 ± 0.028</td><td>0.096 ± 0.036</td><td>0.010 ± 0.036</td><td>0.021 ± 0.084</td><td>0.055 ± 0.018</td><td>0.021 ± 0.084</td><td>0.096 ± 0.036</td><td>0.010 ± 0.036</td><td>0.055 ± 0.018</td></tr><tr><td>Ni IIλ1317</td><td>0.131 ± 0.009</td><td>0.209 ± 0.015</td><td>0.096 ± 0.011</td><td>0.121 ± 0.011</td><td>0.123 ± 0.012</td><td>0.111 ± 0.013</td><td>0.109 ± 0.013</td><td>0.157 ± 0.012</td><td>0.109 ± 0.013</td><td>0.123 ± 0.012</td><td>0.111 ± 0.013</td><td>0.157 ± 0.012</td></tr><tr><td>Ni IIλ1370</td><td>0.057 ± 0.005</td><td>0.078 ± 0.008</td><td>0.034 ± 0.007</td><td>0.083 ± 0.006</td><td>0.105 ± 0.007</td><td>0.012 ± 0.007</td><td>0.056 ± 0.007</td><td>0.078 ± 0.008</td><td>0.056 ± 0.007</td><td>0.105 ± 0.007</td><td>0.012 ± 0.007</td><td>0.056 ± 0.007</td></tr><tr><td>Ni IIλ1454</td><td>0.029 ± 0.003</td><td>0.049 ± 0.005</td><td>0.027 ± 0.004</td><td>0.044 ± 0.004</td><td>0.039 ± 0.005</td><td>0.026 ± 0.005</td><td>0.034 ± 0.005</td><td>0.031 ± 0.004</td><td>0.034 ± 0.005</td><td>0.039 ± 0.005</td><td>0.026 ± 0.005</td><td>0.031 ± 0.004</td></tr><tr><td>Ni IIλ1467</td><td>0.024 ± 0.003</td><td>0.055 ± 0.005</td><td>0.009 ± 0.004</td><td>0.039 ± 0.004</td><td>0.025 ± 0.005</td><td>0.014 ± 0.005</td><td>0.031 ± 0.005</td><td>0.026 ± 0.004</td><td>0.031 ± 0.005</td><td>0.025 ± 0.005</td><td>0.014 ± 0.005</td><td>0.026 ± 0.004</td></tr><tr><td>Ni IIλ1709</td><td>0.059 ± 0.004</td><td>0.095 ± 0.007</td><td>0.036 ± 0.006</td><td>0.083 ± 0.005</td><td>0.087 ± 0.006</td><td>0.039 ± 0.006</td><td>0.038 ± 0.006</td><td>0.073 ± 0.006</td><td>0.038 ± 0.006</td><td>0.087 ± 0.006</td><td>0.039 ± 0.006</td><td>0.073 ± 0.006</td></tr><tr><td>Ni IIλ1741</td><td>0.079 ± 0.005</td><td>0.144 ± 0.008</td><td>0.054 ± 0.007</td><td>0.101 ± 0.006</td><td>0.111 ± 0.007</td><td>0.035 ± 0.007</td><td>0.070 ± 0.007</td><td>0.078 ± 0.007</td><td>0.070 ± 0.007</td><td>0.111 ± 0.007</td><td>0.035 ± 0.007</td><td>0.078 ± 0.007</td></tr><tr><td>Ni IIλ1751</td><td>0.063 ± 0.005</td><td>0.069 ± 0.007</td><td>0.031 ± 0.006</td><td>0.067 ± 0.005</td><td>0.072 ± 0.006</td><td>0.023 ± 0.007</td><td>0.059 ± 0.006</td><td>0.034 ± 0.007</td><td>0.059 ± 0.006</td><td>0.023 ± 0.007</td><td>0.059 ± 0.006</td><td>0.034 ± 0.007</td></tr><tr><td>O IIλ1302</td><td>0.732 ± 0.009</td><td>1.186 ± 0.013</td><td>0.457 ± 0.011</td><td>0.903 ± 0.010</td><td>0.794 ± 0.013</td><td>0.669 ± 0.013</td><td>0.715 ± 0.013</td><td>0.724 ± 0.013</td><td>0.715 ± 0.013</td><td>0.669 ± 0.013</td><td>0.715 ± 0.013</td><td>0.724 ± 0.013</td></tr><tr><td>Si IIλ1304</td><td>0.546 ± 0.009</td><td>0.930 ± 0.014</td><td>0.341 ± 0.012</td><td>0.649 ± 0.011</td><td>0.592 ± 0.014</td><td>0.520 ± 0.013</td><td>0.531 ± 0.014</td><td>0.540 ± 0.013</td><td>0.531 ± 0.014</td><td>0.592 ± 0.014</td><td>0.520 ± 0.013</td><td>0.540 ± 0.013</td></tr><tr><td>Si IIλ1526</td><td>0.674 ± 0.007</td><td>1.206 ± 0.010</td><td>0.427 ± 0.007</td><td>0.846 ± 0.008</td><td>0.752 ± 0.010</td><td>0.606 ± 0.010</td><td>0.649 ± 0.011</td><td>0.694 ± 0.010</td><td>0.649 ± 0.011</td><td>0.752 ± 0.010</td><td>0.606 ± 0.010</td><td>0.694 ± 0.010</td></tr><tr><td>Si IIλ1808</td><td>0.141 ± 0.006</td><td>0.262 ± 0.009</td><td>0.090 ± 0.007</td><td>0.207 ± 0.006</td><td>0.221 ± 0.008</td><td>0.100 ± 0.008</td><td>0.142 ± 0.008</td><td>0.159 ± 0.008</td><td>0.142 ± 0.008</td><td>0.221 ± 0.008</td><td>0.100 ± 0.008</td><td>0.159 ± 0.008</td></tr><tr><td>Si ivλ1393</td><td>0.321 ± 0.006</td><td>0.543 ± 0.010</td><td>0.208 ± 0.007</td><td>0.373 ± 0.007</td><td>0.342 ± 0.009</td><td>0.268 ± 0.008</td><td>0.312 ± 0.009</td><td>0.331 ± 0.008</td><td>0.312 ± 0.009</td><td>0.342 ± 0.009</td><td>0.268 ± 0.008</td><td>0.331 ± 0.008</td></tr><tr><td>Si ivλ1402</td><td>0.169 ± 0.005</td><td>0.356 ± 0.008</td><td>0.120 ± 0.006</td><td>0.231 ± 0.006</td><td>0.190 ± 0.007</td><td>0.158 ± 0.007</td><td>0.183 ± 0.007</td><td>0.204 ± 0.007</td><td>0.183 ± 0.007</td><td>0.190 ± 0.007</td><td>0.158 ± 0.007</td><td>0.204 ± 0.007</td></tr><tr><td>Ti IIλ1910</td><td>0.021 ± 0.005</td><td>0.039 ± 0.007</td><td>0.015 ± 0.006</td><td>0.037 ± 0.005</td><td>0.030 ± 0.007</td><td>0.011 ± 0.007</td><td>0.015 ± 0.008</td><td>0.017 ± 0.006</td><td>0.015 ± 0.008</td><td>0.030 ± 0.007</td><td>0.011 ± 0.007</td><td>0.017 ± 0.006</td></tr><tr><td>Zn II+Mg IIλ2026</td><td>0.111 ± 0.007</td><td>0.179 ± 0.010</td><td>0.057 ± 0.008</td><td>0.120 ± 0.008</td><td>0.137 ± 0.010</td><td>0.070 ± 0.009</td><td>0.082 ± 0.011</td><td>0.097 ± 0.008</td><td>0.082 ± 0.011</td><td>0.137 ± 0.010</td><td>0.070 ± 0.009</td><td>0.097 ± 0.008</td></tr><tr><td>Zn II+Cr IIλ2062</td><td>0.102 ± 0.006</td><td>0.189 ± 0.010</td><td>0.079 ± 0.008</td><td>0.123 ± 0.007</td><td>0.116 ± 0.009</td><td>0.061 ± 0.009</td><td>0.079 ± 0.011</td><td>0.129 ± 0.007</td><td>0.079 ± 0.011</td><td>0.116 ± 0.009</td><td>0.061 ± 0.009</td><td>0.129 ± 0.007</td></tr></tbody></table>

This paper has been typeset from a TeX/LATEX file prepared by the author.