Probing the Nature of High-redshift Weak Emission Line Quasars: A Young Quasar with a Starburst Host Galaxy

We use the Pan-STARRS1 catalog internal release version (PS1 version 3.4; Chambers et al. 2016) as the main data for the initial quasar candidate selection. The 5σ limiting magnitudes of this stack catalog are g = 23.3, r = 23.2, i = 23.1, z = 22.3, y = 21.3. Due to intervening intergalactic medium (IGM), the z ≳ 6.2 quasar's flux on the blue side of Lyα is expected to be heavily absorbed, creating a strong break in the flux. This will also make them basically undetected, i.e., have very low signal-to-noise ratios (S/Ns), in all bands bluer than the z band at the PS1 limiting magnitude. For the selection, we require candidates to be detected in the PS1 y band and have very red colors. In summary, the criteria used are

$$\begin{eqnarray}&&{\rm{S}}/{\rm{N}}({g}_{\mathrm{PS}1},\,{r}_{\mathrm{PS}1},\,{i}_{\mathrm{PS}1})\lt 8\end{eqnarray} \tag{ 1 }$$

$$\begin{eqnarray}&&{\rm{S}}/{\rm{N}}({y}_{\mathrm{PS}1})\gt 7\end{eqnarray} \tag{ 2 }$$

$$\begin{eqnarray}&&{\rm{S}}/{\rm{N}}({z}_{\mathrm{PS}1})\geqslant 3\,\mathrm{and}\,{z}_{\mathrm{PS}1}-{y}_{\mathrm{PS}1}\gt 1.2\end{eqnarray} \tag{ 3 }$$

$$\begin{eqnarray}&&\begin{array}{r}\mathrm{or}\\ {\rm{S}}/{\rm{N}}({z}_{\mathrm{PS}1})\lt 3\,\mathrm{and}\,{z}_{\mathrm{PS}1,\mathrm{lim}}-{y}_{\mathrm{PS}1}\gt 1.2\\ {y}_{\mathrm{PS}1}\gt 15\end{array}\end{eqnarray} \tag{ 4 }$$

where all measured fluxes and magnitudes are based on point-spread function (PSF) photometry, unless stated otherwise. The magnitude limit criterion in Equation (4) is used to exclude unusually bright objects or spurious sources.

However, quasars of extreme brightness could still be detected in some PS1 "dropout" bands. This might be because of their intrinsically high luminosity (e.g., Wu et al. 2015) or their apparent fluxes are boosted by gravitational lensing events (e.g., Fan et al. 2019; Fujimoto et al. 2020). Moreover, most strongly lensed quasars would be removed as candidates due to Equation (1) because the massive lensing galaxy at intermediate redshift would contribute flux in those bands. Hence, additional criterion is applied if Equation (1) is not fulfilled to actually find this population:

$$\begin{eqnarray}&&({g}_{\mathrm{PS}1},\,{r}_{\mathrm{PS}1},\,{i}_{\mathrm{PS}1})-{y}_{\mathrm{PS}1}\gt 3.0.\end{eqnarray} \tag{ 5 }$$

There are some useful parameters in the PS1 catalog that can be utilized to remove most of the contaminants. The first one is to exclude objects showing extended morphology and only choose point like or very compact sources by requiring

$$\begin{eqnarray}&&| {y}_{\mathrm{PS}1,\mathrm{aper}}-{y}_{\mathrm{PS}1}| \lt 0.5,\end{eqnarray} \tag{ 6 }$$

where ${y}_{\mathrm{PS}1,\mathrm{aper}}$ and ${y}_{\mathrm{PS}1}$ are the PS1 catalog aperture and PSF magnitudes of stacked images, respectively. This cutoff threshold was chosen based on tests of spectroscopically confirmed galaxies and stars, following Bañados et al. (2016). Second, the measured PSF magnitudes also need to be consistent with each other:

$$\begin{eqnarray}&&| {y}_{\mathrm{PS}1,\mathrm{stk}}-{y}_{\mathrm{PS}1,\mathrm{wrp}}| \lt 0.5,\end{eqnarray} \tag{ 7 }$$

where "stk" and "wrp" denote photometry from stacked images and the mean photometry of single-epoch images, respectively. Lastly, the expected weighted PSF flux is required to be located in good pixels to a percentage of 85% or more, i.e., $\mathrm{PSF}\_\mathrm{QF}\gt 0.85$ in the catalog. By using the aforementioned criteria, we obtained ∼17 million candidates at the preliminary selection step (see Table 1).

Table 1.  Summary of Quasar Selections

Step	Selection	Candidates	Known
			Quasars
1	Initial query with:
	$\mathrm{PSF}\_\mathrm{QF}\gt 0.85$
	S/N $({y}_{\mathrm{PS}1})\geqslant 5$
	${y}_{\mathrm{PS}1}\gt 15$
	${z}_{\mathrm{PS}1}-{y}_{\mathrm{PS}1}\gt 1.0$
	Excluding Galactic plane
	Excluding M31	17170000	27
2	Detection in NIR or MIR	5631392	27
3	S/N(yPS1) ≥ 7	3627100	27
4	$| {y}_{\mathrm{PS}1,\mathrm{aper}}-{y}_{\mathrm{PS}1}| \lt 0.5$	1158750	24
5	$| {y}_{\mathrm{PS}1,\mathrm{stk}}-{y}_{\mathrm{PS}1,\mathrm{wrp}}| \lt 0.5$	1059373	24
6	${z}_{\mathrm{PS}1}-{y}_{\mathrm{PS}1}\gt 1.2$	367190	22
7	S/N $({g}_{\mathrm{PS}1},\,{r}_{\mathrm{PS}1},\,{i}_{\mathrm{PS}1})\lt 8$, or
	$({g}_{\mathrm{PS}1},\,{r}_{\mathrm{PS}1},\,{i}_{\mathrm{PS}1})-{y}_{\mathrm{PS}1}\gt 3$	74374	22
8	SED fitting	7808	20
9	Forced photometry	1263	20
10	Visual inspection	155	20

Note. Each selection step shows the number of candidates and recovered known quasars at z > 6. Our method is most sensitive to select quasars at 6.3 ≤ z ≤ 7.1 due to step 1 requirements.

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Some areas of PS1 are also covered by the Dark Energy Spectroscopic Instrument Legacy Imaging Surveys Data Release 8 (DELS; Dey et al. 2019) and the Dark Energy Survey Data Release 1 (DES; Abbott et al. 2018). The advantages of DELS and DES are that they reach magnitude levels ∼1 mag fainter compared to PS1. DELS¹² is conducted using imaging data from three different telescopes, covering ∼14,000 deg² of extragalactic sky visible from the northern hemisphere ($-18^\circ \lt \mathrm{decl}.\lt +84^\circ$) in three optical bands. These data reach 5σ depths of g = 24.0, r = 23.4, and z = 22.5 mag (Dey et al. 2019). On the southern hemisphere, DES¹³ is utilizing the Dark Energy Camera mounted on the Cerro Tololo Interamerican Observatory 4 m Blanco telescope to image ∼5000 deg² of the southern Galactic cap region in five broad bands. The median coadded catalog depth for a 195 diameter aperture at S/N = 10 is g = 24.33, r = 24.08, i = 23.44, z = 22.69, and Y = 21.44 mag (Abbott et al. 2018). Hence, the DELS and DES catalogs are cross-matched to our main PS1 catalog with 2'' radius and we will use their photometric data if available.

Selecting high-redshift ("high-z," z ≳ 6) quasar candidates through the color criteria becomes complicated due to contamination from other populations: late-M stars along with L and T dwarfs (MLT dwarfs) and elliptical galaxies at z = 1–2 (hereafter ellipticals). These populations have a higher surface density than the target high-z quasars themselves but have similar NIR colors. Hence, to reduce the number of MLT dwarfs contaminating the candidate sample, we exclude the sky region around M31 ($7^\circ \lt {\rm{R}}.{\rm{A}}.\,\lt 14^\circ ;\,37^\circ \lt \mathrm{decl}.\,\lt \,43^\circ$) and the Milky Way plane ($| b| \lt 20^\circ$). However, with the advantage of the Bayestar19 (Green et al. 2019) dust maps, we include some candidates in the Galactic plane region if they have reddening of $E(B-V)\lt 1$ in the selection, as a difference from previous studies that rigorously excluded all regions with $E(B-V)\gt 0.3$ (e.g., Bañados et al. 2016; Mazzucchelli et al. 2017).

To make a more refined and more robust selection, we need to take advantage of infrared photometry, because it will give us a parameter dimension to distinguish quasars from MLT dwarfs and ellipticals.

We take advantage of public surveys when their sky footprint overlapped with that of PS1. In this case, NIR photometry for the northern hemisphere was taken from the UKIRT Infrared Deep Sky Survey DR10 (UKIDSS; Lawrence et al. 2007) and the UKIRT Hemisphere Survey DR1 (UHS; Dye et al. 2018), while in the southern hemisphere, we used the Vista Hemisphere Survey DR6 (VHS; McMahon et al. 2013) with a 2'' cross-matching radius. The J- as well as H- and K-band magnitudes are very useful to discriminate between quasars and MLT dwarfs (Mazzucchelli et al. 2017; Yang et al. 2019; Wang et al. 2019). This is also an efficient method to remove spurious sources like cosmic rays that are usually only detected in one survey and not in the others.

In addition, mid-infrared (MIR) data were taken from the unWISE catalog (Schlafly et al. 2019), which contains roughly two billion objects, observed by the Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010) over the whole sky. The W1 (3.4 μm) and W2 (4.6 μm) of WISE photometric bands are useful to separate quasars from MLT dwarfs (e.g., the $W1-W2$ color). The advantage of unWISE compared to the original WISE catalog (AllWISE) is significantly deeper imaging data and improved source extraction in crowded regions (Schlafly et al. 2019). A 3'' cross-match radius to our main catalog was applied. By default, we defined that a source is detected in the NIR/MIR if the catalog entry of the corresponding flux measurements is not empty, i.e., there is a match within the cross-matching radius.

We first confirmed PSO J083.8371+11.8482 (hereafter PSO J083+11) as a quasar at z ∼ 6.3 via low-resolution NIR spectroscopic follow-up by using 6.5 m Magellan/FIRE (Simcoe et al. 2013) on 2018 December 31 with a total integration time on target of only 5 minutes. The observation was done using the high-throughput prism mode with a slit width of 06, resulting in a spectral resolution of R = 500. The observed wavelength range covered by this instrument is λ_obs ∼ 0.82–2.51 μm.

A second observing run to take a substantially deeper NIR spectrum of PSO J083+11 was done in January and February 2019, using the same telescope. The quasar was observed for 5 hr in the high-resolution echellette mode with the 06 slit. This in principle gives us R = 6000 spectral resolution or around ∼50 km s⁻¹ velocity resolution in the wavelength range of 0.82–2.51 μm. Unfortunately, this second observation run suffered from suboptimal condition, which degraded the S/N. Still, a quick look at the spectra showed that all key emission lines—e.g., Lyα λ 1216, C iv λ 1549, and Mg ii λ 2798—appear unexpectedly weak. We modeled the emission lines and continuum despite the noisy spectrum and found that its continuum power-law slope is consistent with a Type 1 quasar located at z = 6.34. We have to note that the region around Mg ii is heavily affected by telluric absorption, which made the accurate line measurement difficult. To mitigate this issue and to reduce instrument-specific effects, we carried out another spectroscopic observing run.

We obtained deep NIR spectroscopy with Gemini Near-InfraRed Spectrograph (GNIRS) at the 8.1 m Gemini North telescope (GN-2019A-FT-204, PI: M. Onoue) on 2019 March 20–22 with a total integration time on target of 8060 s. The runs were executed in the cross-dispersed mode to cover the observed wavelength range λ_obs ∼ 0.9–2.5 μm, corresponding to ${\lambda }_{\mathrm{rest}}\sim 1200-3400$ Å in the rest frame. We utilized the "short" camera with a resolution of 015 per pixel and a 31.7 l/mm grating. We used a slit with aperture of 0675, resulting in a spectral resolution of R ∼ 750. The exposure time for single frame was set to 155 s, and a standard ABBA offset pattern was applied between exposures to better remove the contribution from the sky emission. The airmass range of the observations was ∼1.1–1.7.

Data reduction was performed with PypeIt,¹⁶ an open-source spectroscopic data reduction pipeline (Prochaska et al. 2020). Each exposure was bias-subtracted and flat-fielded using standard procedures. The wavelength solution was obtained by comparing the spectrum of the sky with the prominent OH (Rousselot et al. 2000) and water lines.¹⁷ After removing contamination from cosmic rays, the pipeline optimally subtracts the background by modeling the sky emission with a b-spline function that follows the curvature of the spectrum on the detector. The 1D spectrum of the quasar was created for each exposure using optimal weighting (Kelson 2003). Relative flux calibration was performed with A-type stars observed before or after the target exposures. We used Molecfit¹⁸ (Kausch et al. 2015; Smette et al. 2015) to do telluric absorption correction, after which single-exposure one-dimensional spectra were coadded. All 1D spectra were coadded and scaled to the observed UHS J-band photometry (J_UHS = 20.09 ± 0.13, not corrected for Galactic extinction) for absolute flux calibration. The reddening due to Galactic extinction is then corrected by using the dust map from Green et al. (2019) and the extinction law from Gordon et al. (2016). Figure 2 shows the final spectrum in rest frame. From now on, we will use the GNIRS spectrum for the primary analysis instead of the FIRE spectrum, unless otherwise stated explicitly.

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To model the PSO J083+11 spectrum, we used a multicomponent fitting approach. The global continuum was modeled using combined power-law and UV Fe ii templates (Vestergaard & Wilkes 2001; Tsuzuki et al. 2006; Salviander et al. 2007). Free parameters are the scaling factor for each component and the power-law continuum slope. We implemented this approach by using a modified version of PyQSOFit,¹⁹ a code to fit typical quasar spectra (Guo et al. 2018). We improved the performance of PyQSOFit with respect to high-z quasars by optimizing the fit of the narrow and broad emission lines, substituting the continuum windows and other minor modifications. We define the continuum window for the fit by iteratively marking emission- and telluric-line-free regions. The following wavelength ranges were selected as continuum windows: λ_rest = 1285–1290 Å, 1315–1325 Å, 1350–1370 Å, 1445–1465 Å, 1580–1650 Å, 2140–2300 Å, 2340–2400 Å, 2420–2480 Å, 2630–2710 Å, 2745–2765 Å, and 2850–3000 Å.

The rest-frame 1450 Å absolute magnitude and the monochromatic luminosity at 3000 Å (${L}_{\lambda }(3000\,\mathring{\rm A} )$) were calculated from the best-fit power-law continuum. Then, the bolometric luminosity L_bol was derived using the empirical correction from Richards et al. (2006):

$$\begin{eqnarray}&&{L}_{\mathrm{bol}}=5.15\times \lambda {L}_{3000}.\end{eqnarray} \tag{ 12 }$$

After subtracting the best-fit continuum and scaled iron templates, each broad emission line was modeled with single Gaussian functions. We applied Monte Carlo simulations and created mock spectra to estimate the errors of the derived parameters. Random flux density errors are drawn assuming a normal distribution using the noise spectrum then applied to the raw spectrum with 1000 iterations. The measurement lower and upper-limit values are taken as the 16th and 84th percentiles of the distribution of these repeated measurements, respectively. Finally, we obtained the equivalent width (EW), central wavelength, FWHM, and velocity dispersion for each line.

However, as seen in Figure 2, the only strongly detected broad line is Mg ii, where we found FWHM (Mg ii) $={4140}_{-430}^{+880}$ km s⁻¹ and EW (Mg ii)_rest $=\,{8.71}_{-0.64}^{+0.67}$ Å. Moreover, we obtained the redshift of ${z}_{\mathrm{Mg}{\rm\small{II}}}=6.346\pm 0.001$, determined from the observed Mg ii central wavelength. The Lyα is weak, and we do not significantly detect the C iv line. The derivation of the Lyα and C iv EWs will be explained later in Section 6. Assuming that C iv will have an FWHM similar to or larger than that of Mg ii, we do not see any potential broad absorption-line signatures in the region where C iv is expected to be present (see the bottom panel of Figure 2). Hence, we conclude that the absence of C iv is not simply caused by the broad absorption-line phenomenon, but rather by the actual nature of this quasar (see Section 6).

The mass of the black hole is derived from our single-epoch NIR spectra. With the assumption that the virial theorem is valid for the BLR dynamics, we use the scaling relation for the Mg ii line from Vestergaard & Osmer (2009):

$$\begin{eqnarray}&&\displaystyle \frac{{M}_{\mathrm{BH}}}{{M}_{\odot }}={10}^{6.86}{\left(\displaystyle \frac{\mathrm{FWHM}(\mathrm{Mg}{\rm\small{II}})}{{10}^{3}\mathrm{km}{{\rm{s}}}^{-1}}\right)}^{2}{\left(\displaystyle \frac{\lambda {L}_{\lambda }(3000\mathring{\rm A} )}{{10}^{44}\mathrm{erg}{{\rm{s}}}^{-1}}\right)}^{0.5}\end{eqnarray} \tag{ 13 }$$

where $\lambda {L}_{\lambda }(3000\,\mathring{\rm A} )$ is the rest-frame luminosity at 3000 Å and FWHM(Mg ii) is the Mg ii FWHM. Then, we calculated the Eddington luminosity as

$$\begin{eqnarray}&&{L}_{\mathrm{Edd}}=1.3\times {10}^{38}\left(\displaystyle \frac{{M}_{\mathrm{BH}}}{{M}_{\odot }}\right)\,\mathrm{erg}\,{{\rm{s}}}^{-1}.\end{eqnarray} \tag{ 14 }$$

After that, we can derive the Eddington ratio as ${L}_{\mathrm{bol}}/{L}_{\mathrm{Edd}}$. We obtained a black hole mass of $\mathrm{log}({M}_{\mathrm{BH}}/{M}_{\odot })={9.30}_{-0.10}^{+0.16}$ and normalized accretion rate of ${L}_{\mathrm{bol}}/{L}_{\mathrm{Edd}}={0.51}_{-0.17}^{+0.13}$ for the GNIRS spectrum. As a comparison, we derived $\mathrm{log}({M}_{\mathrm{BH}}/{M}_{\odot })\,={9.06}_{-0.16}^{+0.24}$ and ${L}_{\mathrm{bol}}/{L}_{\mathrm{Edd}}={0.77}_{-0.33}^{+0.33}$ from the combined FIRE spectra. We report the error in measured virial black hole mass, which we calculated by propagating monochromatic luminosity errors and Mg ii line width. However, please note that we did not explicitly add systematic errors, which tend to be larger (≈0.5 dex) compared to random measurement errors (Vestergaard & Osmer 2009; Shen 2013).

In order to compare PSO J083+11 properties to other quasars at high redshifts, we compile Mg ii line-width and continuum luminosity measurements for ∼70 previously published z > 5.7 quasars (Jiang et al. 2007; Willott et al. 2010; De Rosa et al. 2011, 2014; Wu et al. 2015; Mazzucchelli et al. 2017; Eilers et al. 2018; Fan et al. 2019; Onoue et al. 2019; Shen et al. 2019; Tang et al. 2019). Then, we classify five quasars with EW $({\rm{C}}\,{\rm\small{IV}})\lt 10$ Å in Shen et al. (2019) as high-z WLQs. On the other hand, 261 WLQs at z ∼ 1.3 are taken from Meusinger & Balafkan (2014). We use their Mg ii line-width and continuum measurements to derive bolometric luminosities and virial masses with the same cosmology assumption and scaling relation (see Equations (12) and (13)) as those applied for PSO J083+11. Figure 3 shows the distribution of calculated parameters in the black hole mass–bolometric luminosity plane. Our object populates the same L_bol/L_Edd and M_BH parameter space as other quasars at z > 5.7 and those observed at z ∼ 2 from SDSS Data Release 7. Hence, we find that PSO J083+11 is powered by a typical mature and actively accreting SMBH, which has been reported in other quasars at similar luminosity ranges. A summary of calculated physical parameters is shown in Table 2.

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Table 2.  Derived Physical Parameters of PSO J083+11

Parameter	Value	Unit
αλ	$-{1.66}_{-0.04}^{+0.01}$
M1450	−26.67 ± 0.01	mag
MBH	$\left({2.00}_{-0.44}^{+0.74}\right)\times {10}^{9}$	M⊙
Lbol	$\left({1.33}_{-0.03}^{+0.01}\right)\times {10}^{47}$	erg s−1
Lbol/LEdd	${0.51}_{-0.17}^{+0.13}$
FWHM (Mg ii)	${4140}_{-430}^{+880}$	km s−1
${\rm{\Delta }}v(\mathrm{Mg}\,{\rm\small{II}}-[{\rm{C}}\,{\rm\small{II}}])$ a	237 ± 150	km s−1
EW ${(\mathrm{Mg}{\rm\small{II}})}_{\mathrm{rest}}$	${8.71}_{-0.64}^{+0.67}$	Å
EW ${(\mathrm{Ly}\alpha +{\rm{N}}{\rm\small{V}})}_{\mathrm{rest}}$	${5.65}_{-0.66}^{+0.72}$	Å
EW ${({\rm{C}}{\rm\small{IV}})}_{\mathrm{rest}}$ b	$\leqslant 5.83$	Å
Rp	1.17 ± 0.32	pMpc
tQ	${10}^{3.4\pm 0.7}$	yr
${z}_{[{\rm{C}}{\rm\small{II}}]}$	6.3401 ± 0.0004
FWHM ([C ii])	229 ± 5	km s−1
Flux ([C ii])	10.22 ± 0.35	Jy km s−1
${S}_{244\mathrm{GHz}}$	5.10 ± 0.15	mJy
${S}_{258\mathrm{GHz}}$	5.54 ± 0.16	mJy
${L}_{[{\rm{C}}{\rm\small{II}}]}$	$(1.04\pm 0.04)\times {10}^{10}$	$\,{L}_{\odot }$
${L}_{\mathrm{FIR}}$	$(1.22\pm 0.07)\times {10}^{13}$	$\,{L}_{\odot }$
${L}_{\mathrm{TIR}}$	$(1.72\pm 0.09)\times {10}^{13}$	$\,{L}_{\odot }$
SFR${}_{[{\rm{C}}{\rm\small{II}}]}$	800–4900	$\,{M}_{\odot }\,{\mathrm{yr}}^{-1}$
SFR${}_{\mathrm{TIR}}$	900–7600	$\,{M}_{\odot }\,{\mathrm{yr}}^{-1}$
${M}_{\mathrm{dust}}$	$(4.88\pm 0.14)\times {10}^{8}$	${M}_{\odot }$

Notes.

^aVelocity shift of Mg ii with respect to [C ii]. ^bThis is the 3σ upper-limit value; see Section 6.

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We obtained high-resolution imaging for PSO J083+11 by using HST (GO 15707, PI: K. Jahnke). Our goal is to test the quasar image being subject to gravitational lensing by utilizing two methods. The first one is by searching for multiple quasar images using the NIR F125W filter (λ_eff = 12365 Å) on the Wide Field Camera 3 (WFC3). The second one is directly searching for an intervening galaxy with the Advanced Camera for Surveys (ACS) ramp filter FR853N (λ_eff = 8528 Å) in the quasar Gunn–Peterson absorption trough just shortward of Lyα, where the quasar light will be nearly fully absorbed, maximizing visibility of any intervening lensing galaxy. Each band was exposed for two orbits, for a total integration time of ∼80 minutes per filter. For the WFC3/IR imaging, we rotated the field between the two orbits by ∼15°, to analyze not only the images directly, but also the difference image that reduced the interference of instrumental point-source spikes integral to the HST PSF. Data reduction was carried out with the HST pipeline for a final image pixel scale of 0128 pixel⁻¹. The 5σ surface brightness limit for a 1 arcsec² aperture is ∼26 mag arcsec⁻².

Our next step is to analyze the WFC3/F125W HST image and search for extended emission, which could be due to an intervening lensing galaxy, companion source, or the host galaxy of PSO J083+11 itself. Our approach is to remove the point-like emission from the central quasar, utilizing the Photutils²⁰ software (Bradley et al. 2019). We chose eight stars in the field to construct a PSF using median-averaging. The stars are chosen so that they are sufficiently away from the edge of the CCD or potential contaminants. They also have to be 1.5×–15× brighter than the quasar to get accurate PSF wings. Note that we did not take into account the effect of the spectral types of selected reference stars; hence, there might be color-dependent uncertainties due to systematic SED differences. Then, this model is fitted to the quasar's nuclear emission in the image, allowing the PSF centroid to move by less than a pixel. The observed image and PSF subtraction residual are shown in Figure 4.

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We identified emission in the WFC3/F125W images, which could be attributed to an intervening foreground galaxy having a potential gravitational lensing effect, located at 1'' to the southwest from the central quasar. However, we could not constrain a redshift of this emission by using HST data only. Its AB magnitude was estimated using aperture photometry, with an aperture of 4 pixel (∼05) radius with a value of ${\mathrm{mag}}_{{\rm{F}}125{\rm{W}}}=25.42\pm 0.07$. However, this emission does not show up in the ACS/FR853N image. We did the photometry with the same aperture size and obtained a 5σ upper-limit magnitude of ${\mathrm{mag}}_{\mathrm{FR}853{\rm{N}}}=23.25$.

We test whether the intervening galaxy could potentially boost the apparent quasar emission by modeling the galaxy's emission from the far-ultraviolet to the microwave regimes using BAGPIPES²¹ (Carnall et al. 2018). We used the measured flux from WFC3/F125W image aperture photometry, complemented with upper-limit fluxes from the ACS/FR853N and PS1 bands as input data. Here, we assume the foreground emission, which was found in the F125W images to be from a star-forming main-sequence galaxy. To derive the stellar mass and star formation rate (SFR), we need to scale the galaxy SED by choosing a particular star formation history (SFH) model. The simplest form of a parametric model uses up to three shape parameters and a normalization. The most simple and widely used parametric SFH model is exponentially declining (tau model, τ). For this work, we consider the delayed exponentially declining SFHs, which is a tau model multiplied by the time since star formation began (T₀). This would remove both a discontinuity in SFR at T₀ and models with rising SFHs if τ is large (see Carnall et al. 2018). The model will fit the total formed stellar mass and time since star formation began with uniform priors in the value ranges of ${M}_{\mathrm{formed}}={10}^{1-15}{M}_{\odot }$ and ${T}_{0}=0.5-0.8\,\mathrm{Gyr}$, respectively. We assume a fixed value for the SFR timescale (τ = 0.3 Gyr), the metallicity (Z = 0.02; equals to solar metallicity), and the nebular emission parameter ($\mathrm{log}(U)=-3$). The value of A_V = 1.0 mag was also applied following the Calzetti et al. (2000) extinction law.

For us, the central output calculated by BAGPIPES is the redshift-dependent galaxy stellar mass, which we will use to constrain the possible magnification for PSO J083+11 by applying the lensing equation (see Schneider 2006 for a review). Given the upper limit for the mass-to-light ratio (∼100) for disk and ellipticals, we can estimate the maximum total lens mass, including the dark matter contribution, from the galaxy stellar mass output of the model. Assuming a point-mass gravitational lens configuration, the possible combinations of Einstein angle (θ_E) and magnification (μ) can be calculated. Given the calculated lensing galaxy masses, we find limits of θ_E ≤ 05 and μ ≤ 1.07. Therefore, strong magnification of the quasar's emission by the foreground galaxy can be excluded. The calculated Einstein angles as a function of simulated galaxy masses and redshifts are shown in Figure 5.

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The ALMA band-6 observations were performed to spatially resolve the [C ii] 158 μm line, determine an accurate redshift, and investigate the host galaxy of PSO J083+11 (2019.1.01436.S, PI: I. T. Andika). We integrated for a total on-source time of 3145 s with observation carried out on 2019 October 9 using ALMA's C43-4 array configuration. The receivers were set to cover ∼258 GHz, which is the expected [C ii] frequency (${\nu }_{\mathrm{rest}}=1900.5369$ GHz) at z = 6.34 (see Section 4.2).

The data set was calibrated with a pipeline implemented in Common Astronomy Software Application²² (CASA; McMullin et al. 2007). Next, the TCLEAN task was applied using natural weighting to image the visibilities while maximizing the point-source sensitivity. We focused on a 5'' circular region around the central quasar and performed the cleaning down to 2σ. Then, we masked out the line-containing channels and used a simple median approximation to model the continuum. After that, we produced continuum-free data by subtracting the continuum model from visibilities. The final imaged cubes have a 30 MHz channel width, a synthesized beam of around 042 × 037, and an rms noise level of ∼0.24 mJy beam⁻¹.

The [C ii] spectrum is extracted with an aperture radius of 15 (equivalent to 8.3 kpc) to maximize the recoverable emission. We chose this value because there is no further [C ii] flux recovered outside this radius. Ill-defined units make the flux measurement in interferometric maps challenging (see Novak et al. 2019 for further details). Assume that we have the aperture fluxes measured inside the dirty image D, the clean component C only, and the residual image R. Then, the scaling factor for corrections is $\epsilon =C/(D-R)$. The most important factors that govern the scaling factor are the clean and dirty beam sizes. In our case,  ≈ 0.6. Then, we use the dirty image flux and multiply it with to get the final flux density in proper units. Note that if using the final image only to measure the flux, we would obtain ∼10% larger values. This happens because the final image is a superposition of clean Gaussian components plus residuals. The extracted spectrum is shown in Figure 9, and a Gaussian fit results in the integrated [C ii] line flux of 10.22 ± 0.35 Jy km s⁻¹, FWHM of 229 ± 5 km s⁻¹, and a redshift of 6.3401 ± 0.0004. With respect to [C ii] (host galaxy tracer), the Mg ii (quasar BLR tracer) is redshifted by 237 ± 150 km s⁻¹. As a comparison, Schindler et al. (in preparation) found that Mg ii in high-z quasars can be significantly shifted with respect to the [C ii] with a median velocity shift of $-{416}_{-398}^{+304}$ km s⁻¹. The Mg ii line, which arises from the BLR, may experience strong internal motions or winds, potentially displacing the emission line centers from the systemic redshift (Mazzucchelli et al. 2017).

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We covered two dust continuum spectral windows in our observation, one centered at 244 GHz and the other at 258 GHz. Pure continuum data are produced by selecting channels that are free from line emission, where we chose the effective bandwidth for each spectral window to be around 2500 km s⁻¹. Then, we collapsed each spectral window to create a moment-zero map of the dust continuum. The final 244 GHz and 258 GHz dust continuum maps have rms noise levels of 30.7 and 35.8 μJy beam⁻¹, respectively. Figure 10 shows the continuum map centered at 258 GHz. By using the same circular aperture size as for [C ii]—i.e., a radius of 15—we obtain flux densities of ${S}_{244\mathrm{GHz}}=5.10\pm 0.15$ mJy and ${S}_{258\mathrm{GHz}}=5.54\,\pm 0.16$ mJy. The moment-zero map of [C ii] is also created by collapsing the 700 km s⁻¹ cube width as shown in Figure 10. This width is equivalent to 3  × [C ii] FWHM.

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As is visible in Figure 10, both continuum and [C ii] emission are spatially resolved with approximate sizes of 2.3 kpc × 1.7 kpc and 8.2 kpc × 4.3 kpc, respectively. These sizes are defined as the major and minor axis FWHMs of the 2D Gaussian function, which is fitted to those emissions. In addition, the continuum emission peak overlaps with the central [C ii] emission. In Figure 11, we show the velocity field and dispersion maps of the continuum-subtracted [C ii] emission.²³ A mild velocity gradient is seen with the northern part redshifted and the southern part blueshifted. The morphology in Figure 10 is seen as a single resolved blob while the integrated spectrum (Figure 9) is represented well by a simple Gaussian profile.

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Finally, we overlaid the ALMA moment maps on the PSF-subtracted HST image as seen in Figure 12. Previously, we identified a potential lensing galaxy as shown in Figure 4. However, there is no apparent dust or [C ii] emission at the position of this potential lensing companion (see Figure 12), which makes it unlikely to be physically associated.

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To derive a star formation rate in the quasar host galaxy, we first calculate the [C ii] 158 μm line luminosity following Carilli & Walter (2013):

$$\begin{eqnarray}&&\displaystyle \frac{{L}_{\mathrm{line}}}{{L}_{\odot }}=1.04\times {10}^{-3}\displaystyle \frac{{S}_{\mathrm{line}}{\rm{\Delta }}v}{\mathrm{Jy}\,\mathrm{km}\,{{\rm{s}}}^{-1}}{\left(\displaystyle \frac{{D}_{L}}{\mathrm{Mpc}}\right)}^{2}\displaystyle \frac{{\nu }_{\mathrm{obs}}}{\,}\mathrm{GHz}.\end{eqnarray} \tag{ 15 }$$

We obtain a luminosity of ${L}_{[{\rm{C}}{\rm\small{II}}]}=(1.04\pm 0.04)\times {10}^{10}\,{L}_{\odot }$, which interestingly makes this one of the most luminous [C ii] line z > 6 quasars detected to date (Decarli et al. 2018; Venemans et al. 2018). The star formation rate (SFR) can be estimated by applying the De Looze et al. (2014) SFR–${L}_{[{\rm{C}}{\rm\small{II}}]}$ scaling relations for z > 5 galaxies:

$$\begin{eqnarray}&&\displaystyle \frac{{\mathrm{SFR}}_{[{\rm{C}}{\rm\small{II}}]}}{{M}_{\odot }{\mathrm{yr}}^{-1}}=3.0\times {10}^{-9}{\left(\displaystyle \frac{{L}_{[{\rm{C}}{\rm\small{II}}]}}{{L}_{\odot }}\right)}^{1.18}.\end{eqnarray} \tag{ 16 }$$

This will give us a value of the SFR${}_{[{\rm{C}}{\rm\small{II}}]}\sim 1990\,{M}_{\odot }\,{\mathrm{yr}}^{-1}$. However, one needs to take into account a factor of ∼2.5 for the systematic uncertainty from the scaling relation which makes the possible range of derived SFR for PSO J083+11 SFR${}_{[{\rm{C}}{\rm\small{II}}]}=800-4900\,{M}_{\odot }\,{\mathrm{yr}}^{-1}$.

Another important parameter that we can estimate from the ALMA data are the total infrared luminosity (L_TIR, rest-frame 3–1100 μm), far-infrared luminosity (L_FIR, rest-frame 42.5–122.5 μm), and the dust mass, from which we can again calculate an independent star formation rate (e.g., Helou et al. 1988; Kennicutt & Evans 2012; Carilli & Walter 2013). This can be done by assuming low dust optical depth in the Rayleigh–Jeans regime, thus modeling the SED with a modified blackbody will be sufficient (e.g., Beelen et al. 2006; Novak et al. 2019). The expression is

$$\begin{eqnarray}&&{S}_{{\nu }_{\mathrm{obs}}}={f}_{\mathrm{CMB}}(1+z){\left({D}_{L}\right)}^{-2}{\kappa }_{{\nu }_{\mathrm{rest}}}{M}_{\mathrm{dust}}{B}_{{\nu }_{\mathrm{rest}}}({T}_{\mathrm{dust},z}),\end{eqnarray} \tag{ 17 }$$

where ${B}_{{\nu }_{\mathrm{rest}}}$ is the blackbody radiation function, D_L is the luminosity distance, M_dust is the dust mass, while the observed and rest frequencies are related as ${\nu }_{\mathrm{rest}}=(1+z){\nu }_{\mathrm{obs}}$, and all values are in SI units. The adopted opacity coefficient following Dunne et al. (2003) and Novak et al. (2019) is

$$\begin{eqnarray}&&{\kappa }_{{\nu }_{\mathrm{rest}}}={\kappa }_{{\nu }_{0}}{\left(\displaystyle \frac{{\nu }_{\mathrm{rest}}}{{\nu }_{0}}\right)}^{\beta }=2.64{\left(\displaystyle \frac{{\nu }_{\mathrm{rest}}}{c/125\mu {\rm{m}}}\right)}^{\beta }\,{{\rm{m}}}^{2}\,{\mathrm{kg}}^{-1},\end{eqnarray} \tag{ 18 }$$

where c is the speed of light and β is the dust spectral emissivity index. Note that the estimated dust mass will have at least a factor of 2 for the systematic uncertainty due to the adopted opacity coefficient scaling relation.

Dust heating by the cosmic microwave background (CMB) plays an important role at z ≳ 6 and needs to be taken into consideration following da Cunha et al. (2013), namely,

$$\begin{eqnarray}&&{f}_{\mathrm{CMB}}=1-\displaystyle \frac{{B}_{{\nu }_{\mathrm{rest}}}({T}_{\mathrm{CMB},z})}{{B}_{{\nu }_{\mathrm{rest}}}({T}_{\mathrm{dust},z})},\end{eqnarray} \tag{ 19 }$$

$$\begin{eqnarray}&&{T}_{\mathrm{dust},z}={\left({T}_{\mathrm{dust}}^{\beta +4}+{T}_{\mathrm{CMB},{\rm{z}}=0}^{\beta +4}\left[{\left(1+z\right)}^{\beta +4}-1\right]\right)}^{\tfrac{1}{\beta +4}},\end{eqnarray} \tag{ 20 }$$

where T_dust is the intrinsic dust temperature of the source assuming it is located at redshift zero and T_CMB = 2.73(1+z) K is the temperature of the CMB at a given z. In our case, T_CMB = 20.04 K.

Note that we only have two data points of the continuum measurements—at 244 and 258 GHz—both of them located on the Rayleigh–Jeans tail and not reaching the dust SED peak. This prohibits us from constraining the full SED shape and dust temperature due to degenerate fitting parameters. Hence, further assume that T_dust = 47 K and β = 1.6, which are usually applied for quasar host galaxies at z ≳ 6 (e.g., Beelen et al. 2006; Decarli et al. 2018; Venemans et al. 2018). Scaling the SED model (Equation (17)) to the observed FIR photometry at 244 and 258 GHz results in ${M}_{\mathrm{dust}}=(4.88\pm 0.14)\times {10}^{8}\,{M}_{\odot }$. By integrating the SED, we obtain ${L}_{\mathrm{FIR}}=(1.22\pm 0.07)\,\times {10}^{13}\,{L}_{\odot }$ and ${L}_{\mathrm{TIR}}=(1.72\pm 0.09)\times {10}^{13}\,{L}_{\odot }$. This high luminosity means that PSO J083+11 can be classified as a hyperluminous infrared galaxy (HyLIRG, ${L}_{\mathrm{TIR}}\gt {10}^{13}\,{L}_{\odot }$).

The TIR luminosity can be converted to an SFR by utilizing the relation from Murphy et al. (2011) and Kennicutt & Evans (2012):

$$\begin{eqnarray}&&{\mathrm{SFR}}_{\mathrm{TIR}}[{M}_{\odot }\,{\mathrm{yr}}^{-1}]=1.49\times {10}^{-10}{L}_{\mathrm{TIR}}[{L}_{\odot }].\end{eqnarray} \tag{ 21 }$$

This gives us ${\mathrm{SFR}}_{\mathrm{TIR}}\sim 2560\,{M}_{\odot }\,{\mathrm{yr}}^{-1}$. Accounting for the factor of 3 of systematic uncertainty in the scaling relation, we obtain the 1σ possible range of ${\mathrm{SFR}}_{\mathrm{TIR}}=900\mbox{--}7600\,{M}_{\odot }\,{\mathrm{yr}}^{-1}$. This is consistent with SFR estimates based on the [C ii] luminosity above. A summary of the calculated parameters is shown in Table 2.

To compare the PSO J083+11 host galaxy properties with other z ≳ 6 quasars, we took ${L}_{[{\rm{C}}{\rm\small{II}}]}$ and L_FIR measurements from the literature (Walter et al. 2009; Venemans et al. 2012; Wang et al. 2013, 2016; Willott et al. 2013, 2015, 2017; Bañados et al. 2015; Venemans et al. 2016; Mazzucchelli et al. 2017; Decarli et al. 2018), which were recalculated by Decarli et al. (2018). We also added 10 young quasars studied by Eilers et al. (2020). As seen in Figure 13, PSO J083+11 shows values of ${L}_{[{\rm{C}}{\rm\small{II}}]}$ and L_FIR comparable to those high-z quasars with the highest SFR and FIR emission.

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One might expect that the BLR is less photoionized so the broad emission lines produced are weak due to soft ionizing continuum. How could this happen? The first proposed mechanism is an intrinsically soft continuum due to a cold accretion disk around a slowly accreting, very massive black hole (Laor & Davis 2011; Plotkin et al. 2015), although this possibility is rather low in the case of PSO J083+11. This is because the required critical mass is ${M}_{\mathrm{BH}}\gt 3.6\times {10}^{9}\,{M}_{\odot }$ for a nonrotating (a = 0) or ${M}_{\mathrm{BH}}\gt 1.4\times {10}^{10}\,{M}_{\odot }$ for a nearly maximally rotating (a = 0.998) black hole, while the mass estimated for PSO J083+11 is already a factor of ∼2 smaller than the lower of these limits (see Section 4.4). In addition, according to Volonteri et al. (2013), high-z SMBHs should be rapidly spinning, which would yet increase this difference.

The second proposed mechanism is a quasar with a high Eddington ratio that is expected to have an optically and geometrically thick inner accretion disk, a so-called slim disk (Luo et al. 2015). The scale height of this thick disk grows as a function of Eddington ratio and becomes a shielding component, which prevents the high-energy photons from the central region to reach and ionize the BLR, leading to the observed weak, high-ionization line emission (Ni et al. 2018). This model has been corroborated by observation of WLQs with weak X-ray emission that typically show harder X-ray spectra compared to normal quasars, indicating intrinsic X-ray absorption in these objects (Wu et al. 2011, 2012). However, we cannot completely rule out the possibility of the presence of shielding gas between the accretion disk and BLR due to the lack of X-ray spectroscopy for PSO J083+11.

The third proposed mechanism is that an extremely high accretion rate could lead to less efficient production of X-rays, which results in an SED that peaks in the ultraviolet (Leighly et al. 2007a, 2007b). This could be related to a quenched X-ray corona, making it smaller in size, or alternatively, X-ray photons that are trapped and advected into the black hole before they can diffuse out. In other words, in these scenarios, there would not be enough high-energy photons emitted from the continuum source to produce high-ionization potential lines like C iv. However, there would be no problems with producing lower-ionization lines (e.g., Hα, Hβ, and Mg ii).

The second and third aforementioned mechanisms suggest that on average the low-z WLQs have significantly higher Eddington ratios and luminosities compared to other normal low-z quasars (Meusinger & Balafkan 2014). In contrast, we do not see this behavior for PSO J083+11 compared to other high-z quasars because its Eddington ratio is based on the Mg ii line, and the underlying continuum is not particularly strong (L_bol/L_Edd ∼ 0.5; see Figure 3). However, we have to note that the virial-mass calculation relies on the empirical scaling relation derived via reverberation mapping, which might be inadequate in the case of WLQs due to the weakness of their emission lines (Luo et al. 2015). The calculated black hole mass tends to be underestimated, making the actual Eddington ratio potentially lower than currently estimated (Marculewicz & Nikolajuk 2019). Moreover, the typical systematic error of the single-epoch virial-mass approach is ≈0.4–0.5 dex (Shen 2013) and could be larger if the Mg ii BLR is complex and may not be virialized yet in these kinds of exceptional objects (Plotkin et al. 2015).

Another potential indicator of extremely high accretion rate is the shape of the continuum itself, as found in a super-Eddington (L_bol/L_Edd ≥ 9) quasar, PSO J006+39, at z ∼ 6.6 with a very blue continuum that was studied by Tang et al. (2019). This Eddington ratio was estimated by modeling the SED of the UV continuum, where they obtained a power-law slope index, α_λ = −2.94 ± 0.03. This is significantly bluer than the slope of α_λ = −2.33 predicted from the standard thin disk model (Shakura & Sunyaev 1973). Although PSO J006+39 is accreting at a super-Eddington rate, it is not a WLQ (EW ${({\rm{C}}{\rm\small{IV}})}_{\mathrm{rest}}\sim 84$ Å). Compared to that object, our measured power-law slope index is consistent with a typical quasar (α_λ = −1.66 ± 0.01) and not steeper than those aforementioned models, indicating the absence of super-Eddington accretion. Hence, overall, soft continuum models seem inadequate to explain the weak-line nature of PSO J083+11.

With the caveats of all the scenarios presented in Section 8.1, trying to explain the WLQ nature, another potential explanation for the weak-line nature might lie in gas-deficient or anemic BLR clouds, due to the quasar being in the beginning of its accretion phase. In the first scenario, there might be only a small amount of gas and/or covering factor in the BLR of WLQs (Shemmer et al. 2010; Nikołajuk & Walter 2012). If this is true, all broad lines should have small flux and equivalent widths. However, unlike normal quasars, WLQs have relatively small EW ${({\rm{C}}{\rm\small{IV}})}_{\mathrm{rest}}$/EW ${({\rm{H}}\beta )}_{\mathrm{rest}}$ line ratios, which are inconsistent with their BLRs simply having low gas content or small covering factors, although it could play a secondary role (Plotkin et al. 2015).

On the other hand, the second scenario proposed that in the beginning of an AGN phase, BLRs are still bare because the material from the accretion disk has not yet have sufficient time to reach the region where broad lines will later form (Hryniewicz et al. 2010). With the assumption that the wind from the disk has velocities of ≈100 km s^–1, the time needed to form the BLR is around ∼10³ yr (Hryniewicz et al. 2010). Therefore, if WLQs are indeed an evolutionary quasar phase, they should be rare. An interesting discovery is that the fraction of WLQs among quasars seems to increase with redshift (e.g., Diamond-Stanic et al. 2009; Bañados et al. 2016; Shen et al. 2019). In this picture of a still-forming BLR during the WLQ phase, higher-ionization species such as C iv would be the weakest because they originate in a region higher above the disk that is not fully formed yet. On the other hand, low-ionization species (e.g., Hβ, Mg ii) that are formed close to the accretion disk may look normal (Plotkin et al. 2015).

We should emphasize that most low-z WLQ studies do not have access to Lyα and the quasar lifetime could not be determined with the method that we explained in Section 6. At $z\gt 6$, the proximity zones are sensitive to the lifetime of the quasars as the intergalactic gas has a finite response time to the quasars' radiation. The accretion lifetime of PSO J083+11 as derived from its proximity zone gives us a range of ${t}_{{\rm{Q}}}={10}^{3.4\pm 0.7}$ yr, consistent with the BLR formation time at the lower end. Hence, there is a possibility that the BLR in this object is not fully formed yet, leading to the observed weak emission line signature in the spectrum (Hryniewicz et al. 2010). However, as pointed out by Eilers et al. (2018), this quasar could have a higher actual age of substantial accretion compared to the one estimated from its proximity zone size. In that case, it could have been growing in a highly obscured phase and the UV continuum radiation had only broken out of this obscuring medium ∼10³ to 10⁴ yr before (see also Hopkins et al. 2005; DiPompeo et al. 2017; Mitra et al. 2018). This might be caused by a huge dust and gas supply at high redshifts, funneled into the center of the host galaxy feeding both star formation and SMBH, but hiding the quasar within it. In line with that, even though tailored at lower-z systems, Sanders et al. (1988a) argued that ultraluminous infrared galaxies (ULIRGs) could be the initial stage of a quasar when heavy obscuration was present. Only at the end of this incipient dust-enshrouded phase would the quasar be revealed in the optical as an unobscured source (e.g., Sanders et al. 1988b; Sanders & Mirabel 1996; Hopkins et al. 2008). A model proposed by Liu & Zhang (2011) even predicts that the quasars that have just become unobscured should exhibit no broad emission lines because the BLR would form at a later stage when the dusty torus supplies the fuel to the accretion disk. The luminous FIR properties that we found in Section 7.4 would suggest that PSO J083+11 is just at the beginning of its unobscured quasar phase while the host galaxy still experiences highly active star formation. This picture is well consistent with the young accretion lifetime derived from the proximity zone size measurement.

Above, we attributed the small proximity zone of PSO J083+11 to a limited unobscured accretion lifetime. There is a hypothetical alternative, the truncation of PSO J083+11's proximity zone due to the presence of an absorption system within ≤10,000 km s⁻¹ (or z > 6.3) in front of the quasar, just around the edge of the proximity zone. Such a system, like a damped Lyα system (DLA) or Lyman limit system (LLS), would block ionizing radiation from the quasar to the IGM due to its optically thick nature at the Lyman limit (Eilers et al. 2017, 2018; D'Odorico et al. 2018; Bañados et al. 2019; Farina et al. 2019).

In Figure 14, we show hypothetical absorption systems at z = 6.295 (or z = 2.233, see below) that might be able to truncate the proximity zone. This redshift value is equivalent to a distance of R ≈ 2.43 pMpc from the central quasar. For comparison, the proximity zone size that we obtained in Section 6 is R_p = 1.17 ± 0.32 pMpc. The redshift of that system was chosen so that the associated N v $\lambda \lambda 1238,1242$ match the position of the two strong absorption lines observed at 9037 Å and 9066 Å. However, the positions of other lines that should be present in DLAs (e.g., Si ii λ1260, O i λ1302, Si iv λ1402, etc.) do not match any potential absorption lines seen in the observed spectrum. The presence of low-ionization lines is particularly important because they usually indicate an optically thick self-shielding absorption system (Eilers et al. 2020), which can truncate the proximity zone. Also, by considering that the inferred distance between the absorber and quasar is not really close, the proximity zone might be influenced by such a hypothetical absorber, but would unlikely be significant.

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In contrast, there is also an alternative possibility that the two strong absorption lines mentioned before are associated with Mg ii λλ2796,2803 from a lower-redshift (z = 2.233) absorption system instead. However, we have to note that searching for very weak metal absorption features with our current spectrum is difficult due to the low resolution. A more thorough analysis to put stringent constraints on potentially associated absorption systems will be done with VLT/MUSE data, and we report on this in our next paper (I. T. Andika et al. 2020, in preparation).