A Spatially-Resolved Survey of Distant Quasar Host Galaxies: II. Photoionization and Kinematics of the ISM
Andrey Vayner, Shelley A. Wright, Norman Murray, Lee Armus, Anna Boehle, Maren Cosens, James E. Larkin, Etsuko Mieda, Gregory Walth
DDraft version January 22, 2021
Typeset using L A TEX twocolumn style in AASTeX61
A SPATIALLY-RESOLVED SURVEY OF DISTANT QUASAR HOST GALAXIES:II. PHOTOIONIZATION AND KINEMATICS OF THE ISM
Andrey Vayner,
1, 2, 3
Shelley A. Wright,
1, 2
Norman Murray,
4, 5
Lee Armus, Anna Boehle, Maren Cosens,
1, 2
James E. Larkin, Etsuko Mieda, and Gregory Walth Department of Physics, University of California San Diego, 9500 Gilman Drive La Jolla, CA 92093 USA Center for Astrophysics & Space Sciences, University of California San Diego, 9500 Gilman Drive La Jolla, CA 92093 USA Department of Physics and Astronomy, Johns Hopkins University, Bloomberg Center, 3400 N. Charles St., Baltimore, MD 21218, USA Canadian Institute for Theoretical Astrophysics, University of Toronto, 60 St. George Street, Toronto, ON M5S 3H8, Canada Canada Research Chair in Theoretical Astrophysics Spitzer Science Center, California Institute of Technology, 1200 E. California Blvd., Pasadena, CA 91125 USA ETH Z¨urich Wolfgang-Pauli-Str. 27 8093 Z¨urich, Switzerland Department of Physics and Astronomy, University of California, Los Angeles, CA 90095 USA National Astronomical Observatory of Japan, Subaru Telescope, National Institutes of Natural Sciences, Hilo, HI 96720, USA Observatories of the Carnegie Institution for Science 813 Santa Barbara Street Pasadena, CA 91101 USA (Accepted January 15, 2021)
Submitted to ApJABSTRACTWe present detailed observations of photoionization conditions and galaxy kinematics in eleven z= 1 . − .
59 radio-loud quasar host galaxies. Data was taken with OSIRIS integral field spectrograph (IFS) and the adaptive opticssystem at the W.M. Keck Observatory that targeted nebular emission lines (H β ,[O III ] ,H α ,[N II ] ) redshifted into thenear-infrared (1-2.4 µ m). We detect extended ionized emission on scales ranging from 1-30 kpc photoionized by stars,shocks, and active galactic nuclei (AGN). Spatially resolved emission-line ratios indicate that our systems reside off thestar formation and AGN-mixing sequence on the Baldwin, Phillips & Terlevich (BPT) diagram at low redshift. Thedominant cause of the difference between line ratios of low redshift galaxies and our sample is due to lower gas-phasemetallicities, which are 2-5 × less compared to galaxies with AGN in the nearby Universe. Using gas velocity dispersionas a proxy to stellar velocity dispersion and dynamical mass measurement through inclined disk modeling we find thatthe quasar host galaxies are under-massive relative to their central supermassive black hole (SMBH) mass, with allsystems residing off the local scaling ( M • − σ , M • − M ∗ ) relationship. These quasar host galaxies require substantialgrowth, up to an order of magnitude in stellar mass, to grow into present-day massive elliptical galaxies. Combiningthese results with part I of our sample paper (Vayner et al. 2021) we find evidence for winds capable of causing feedbackbefore the AGN host galaxies land on the local scaling relation between black hole and galaxy stellar mass, and beforethe enrichment of the ISM to a level observed in local galaxies with AGN. Corresponding author: Andrey [email protected] a r X i v : . [ a s t r o - ph . GA ] J a n Vayner et al. INTRODUCTIONToday, feedback from supermassive black holes(SMBH) is an integral part of galaxy evolution mod-els. It is commonly used to explain the lack of observedbaryons in local massive galaxies (Behroozi et al. 2010),the enrichment of the circumgalactic medium with met-als (Prochaska et al. 2014) and the observed local scalingrelation between the mass of the galaxy/bulge and theSMBH (Ferrarese & Merritt 2000; Gebhardt et al. 2000;McConnell & Ma 2013).The latest observational and theoretical results pointto a critical question; at what points does the AGN drivean outflow powerful enough to clear the galaxy of its gasinto the surrounding CGM? (King & Pounds 2015) Ac-cording to theoretical work, this typically happens oncethe galaxy reaches the M • − σ relationship (Zubovas& King 2014). However, there has been growing ev-idence for galaxies with massive SMBH and powerfuloutflows that are offset from the local scaling relation-ship (Vayner et al. 2017). The origin and evolution ofthe local scaling relationships with redshift have beenan active debate topic over the last decade. When arethe local scaling relations established? Are the localscaling relationships the end product of galaxy evolu-tion? Meaning, as galaxies form and evolve, do they fallin and out of the relationships due to rapid growth orfeedback process? Do galaxies eventually end up on thelocal scaling relations once the galaxy or SMBH catchup and finish growing (Volonteri 2012)? Alternatively,is there an inherent evolution in the scaling relationshipwith redshift and a symbiosis between the galaxy andSMBH growth? (i.e., evolution in slope, offset, and scat-ter). Finally, there is still an open question regardingthe role of quasar feedback in establishing the relation-ship and whether the merging of galaxies can producethe M • − σ relation following the central limit theorem(Jahnke & Macci`o 2011). From a sample of AGN in theCOSMOS field (Merloni et al. 2010) finds an offset inthe local scaling relationship between redshift 0 and 2.These authors use SED decomposition with numerousspectral bands to measure the stellar mass of the AGNhost galaxy in the redshift range of 1 < z < .
2. From asample of lensed quasars at 1 < z < . ∼ . − ∼ . − > M (cid:12) )are systematically more massive relative to their hostgalaxies (Mart´ın-Navarro et al. 2018). Fields such asCOSMOS or the Extended Chandra Deep Field-Southare relatively small in the sky; hence, the number ofluminous quasars with massive SMBH is small. Stud-ies that explored the evolution of the local scaling re-lationships have generally focused on lower-mass blackholes with masses < M (cid:12) . Furthermore, a largefraction of these studies used broadband HST imag-ing to study the host galaxies of their quasars/AGN.It is often difficult to disentangle the bright AGN emis-sion from the host galaxy at smaller angular separations( < (cid:48)(cid:48) ). These studies have a limited number of filtersto measure the stellar population’s age and the massto light ratio. Alternatively, mm-interferometry obser-vations have become an essential tool in measuring thedynamical masses of quasar host galaxies across differ-ent redshift ranges. At the highest redshifts (z > µ m line has been the most commonly usedtracer of the dynamics of the ISM. There is growing evi-dence that the most massive ( > M (cid:12) ) SMBH in thehighest redshift quasars known to date (z >
6) appearto be over massive for the mass of their host galaxies(Wang et al. 2013; Venemans et al. 2016; Decarli et al.2018), indicating that the most massive SMBHs residingin high redshift quasars grow first relative to their hostgalaxies. At more intermediate redshifts 1 < z <
3, somesystems also appear to have overly massive SMBH rela-tive to their stellar/dynamical mass (Shields et al. 2006;Trakhtenbrot et al. 2015; Vayner et al. 2017). While asignificant fraction of galaxies with lower SMBH < M (cid:12) appear closer or within the scatter of the local scal-ing relations, galaxies with the luminous quasars andmassive SMBH appear to be under massive relative tothe mass of their SMBH. As outlined by (Lauer et al.2007; Schulze & Wisotzki 2014), the offset from the localscaling relations for the systems with more massive blackholes is biased due to the steep decline in the galaxymass function at the massive end.Integral field spectroscopy (IFS) behind adaptive op-tics is another method with which it is possible to dis-entangle the bright quasar emission from the extendedemission of the host galaxy. A point spread functioncan be constructed using data channels confined to thebroad emission line of the quasar. After the point spreadfunction is normalized, it is subtracted from the restof the data channels in the cube. This technique wasfirst shown to be able to resolve host galaxies of lowredshift ( z < .
2) luminous type-1 quasars in seeinglimited observations (Jahnke et al. 2004) and extendedLy α emission around high redshift quasars (Christensen inematics and Energetics of Ionization Gas (cid:48)(cid:48) , compared to ∼ (cid:48)(cid:48) for HST(Vayner et al. 2016). Although today’s near-infraredIFSs are not sensitive enough to detect the stellar con-tinuum from the quasar/AGN host galaxies, they canstill detect extended ionized emission, enabling us toextract the dynamical properties of the galaxy (Inskipet al. 2011; Vayner et al. 2017) and compare systems tothe local scaling relation. However, today, the largestfraction of quasar host galaxy masses still come fromHST and mm-interferometric observations. Most likely,selection effects play an important role in determiningwhether there is a systematic offset from the local scal-ing relations among the different studies.Besides measuring the host galaxies and SMBHmasses, there are vital open questions regarding thegas phase properties. Galaxies exhibit a correlationbetween the stellar mass and metallicity across a wideredshift range (Erb et al. 2006a; Sanders et al. 2015). Itis often difficult to place galaxies with bright AGN onthe mass-metallicity relationship due to limited contrastand the fact that the AGN has a strong impact on theISM’s ionization state. What are the metallicities ofthe gas in quasar hosts? How does the metallicity inquasar host galaxies evolve with redshift? What is thedominant source of ionization in quasar hosts? Whatare the star formation rates? One of the best ways tomeasure the ionization properties of the gas in galax-ies is through the BPT (Baldwin, Phillips & Terlevich)diagram (Baldwin et al. 1981; Veilleux & Osterbrock1987). The traditional BPT diagram plots the ratioof log([O III ]/H β ) vs. log([N II ]/H α ) and contains twoclearly defined sequences: the star-forming sequence andthe mixing sequence. The star-forming sequence pro-vides information about the metallicity of HII regions,the stellar ionizing radiation field as well as informa-tion on the gas condition in star-forming regions. Onthe other hand, the mixing sequence consists of gasphotoionized by hot stars, AGN, and shocks. It canpotentially provide information on the hardness of theAGN ionizing radiation and the metallicity of the gasphotoionized by the quasar/AGN (Groves et al. 2006) and shocks. Studies of high redshift star-forming galax-ies have shown evidence for elevated line ratios relativeto low redshift galaxies. At z ∼
2, the observed elevatedline ratios have been attributed to denser ISM condi-tions (Sanders et al. 2016) and harder ionizing radiationfields at fixed N/O and O/H abundances relative totypical z=0 galaxies (Strom et al. 2017). EvolutionaryBPT models by Kewley et al. (2013a) are consistentwith these observations. The evolutionary BPT modelsalso provide a prediction on the evolution of the mix-ing sequence between z=0 and 3. The location of themixing sequence moves to lower log([N II ]/H α ) valueat a relatively fixed log([O III ]/H β ) value, primarilydue to lower on average gas-phase metallicity at higherredshift (Groves et al. 2006; Kewley et al. 2013a). Thereis tentative evidence that gas photoionized by AGN isconsistent with this picture, as there are several galax-ies with AGN, which have emission line ratios offsetfrom the local mixing sequence (Juneau et al. 2014; Coilet al. 2015; Strom et al. 2017; Nesvadba et al. 2017b;Law et al. 2018). Given the presence of the AGN, youngstars and shocks in quasar host galaxies, it is crucial tospatially resolve the quasar host galaxy to understandthe various contributions to gas ionization. In the dis-tant Universe, this generally requires observations withan IFS and adaptive optics. Resolved BPT diagnosticsin both nearby and distant AGN/quasar host galaxieshave found regions with distinct photoionization mech-anisms (Davies et al. 2014; Williams et al. 2017; Vayneret al. 2017). The question remains whether the ISMcondition in the most luminous high redshift quasarhost galaxies is different from local AGN and wherethey lie relative to the mass metallicity relationship.We have begun a survey to study the host galaxies ofz= 1 . − . β ,[O III ] ,H α ,[N II ],[S II ] ) redshifted in the near-infrared bands (1 − . µ m), at the distance of our sample, the angular res-olution of the OSIRIS/LGS-AO mode corresponds toapproximately 1.4 kpc in projection. Vayner et al.
This paper is part two of two papers focusing on un-derstanding the photoionization mechanisms of gas inradio-loud quasar host galaxies and weigh the mass ofthe galaxy and SMBH to compare them to the localscaling relations. Refer to paper I (Vayner et al. 2021)for details on the sample selection, properties, and datareduction. Details on archival HST imaging data setare presented in §
2. Blackhole masses are presented in §
3, we describe how we identify spatially-resolved dy-namically quiescent regions in each quasar host galaxyin § §
5, dynamical massesof the quasar host galaxies and their place relative tothe local scaling relations is presented in § §
6, wediscuss our results in broader context of massive galaxyevolution in § §
9. Noteson individual sources are presented in §
9. Throughoutthe paper we assume a Λ-dominated cosmology (PlanckCollaboration et al. 2014) with Ω M =0.308, Ω Λ =0.692,and H o =67.8 km s − Mpc − . All magnitudes are onthe AB scale unless otherwise stated. ARCHIVAL
HST
IMAGINGThe sources within our sample have a rich set of multi-wavelength space and ground-based data sets. To as-sist in our analysis and interpretation of distinct regionswithin these quasar host galaxies, we utilize high angu-lar resolution images from the
Hubble Space Telescope .We download fully-reduced data from the Barbara A.Mikulski Archive for Space Telescopes (MAST). Table1 list the archival HST observations used in this study.We construct a model of the PSF using stars in thevicinity of the quasar within the FOV of each instru-ment. Images centered on each star are extracted in abox region of roughly 5 (cid:48)(cid:48) x 5 (cid:48)(cid:48) . We then subtract the localbackground for each star and median combine the stel-lar images into a normalized “master” PSF. This PSF isthen re-scaled to the quasar’s peak flux and subtractedout at the spatial location of the quasar. In cases wherethe quasar was saturated, we scale the flux in the diffrac-tion pattern of the PSF. BLACK HOLE MASS MEASUREMENTBlackhole masses are calculated using the broad-H α line width and luminosity using the scaling relation fromGreene & Ho (2005) for a single epoch SMBH mass es-timate. We describe the details of the nuclear spectrumfitting in Vayner et al. (2021), which comprises of multiGaussian models with a broad component for the BLRemission, a narrow Gaussian for the narrow-line region,and an intermediate width Gaussian for the case wherethere is an outflow. We use the flux and width of the broadest Gaussian to compute the black hole mass. For3C9, 3C298, there are strong telluric/filter transmissionissues that prevent accurate measurement of the FWHMfor the emission line. For these targets, we use the Mg IIsingle epoch black hole mass estimate from Shen et al.(2011). The black hole masses are provided in Table2. We assume an uncertainty of 0.4 dex on the SMBHmasses. DISTINCT REGIONS WITHIN EACH QUASARHOST GALAXYIn this section we outline how we define various regionswithin the data cube of each individual object.4.1.
Spatially-Resolved Dynamically “Quiescent”Regions
In the first survey paper, we outline our methodol-ogy for fitting the emission lines in individual spaxels ofour data cubes. From these fits, we derive integratedintensity, velocity, and velocity dispersion maps. Theerrors on the radial velocity and dispersion maps comedirectly from the Least-squares Gaussian model fit. Theflux map’s errors come directly from integrating a noisespectrum in quadrature over the same wavelength rangewhere the emission line is integrated. The fits are pre-sented in the appendix of (Vayner et al. 2021). Here weutilize the radial velocity and dispersion maps to selectregions with low-velocity dispersion to search for gas ingravitational motion and search for regions where starformation may have recently happened.We define a dynamically “quiescent” region of ourdata set that contains gas with a velocity dispersion( V σ ) less than 250 km s − . A quiescent region thatbelongs to the host galaxy of the quasar must have aradial velocity <
400 km s − as we expect the maximumrotational velocity for a given host galaxy to be at most400 km s − . The maximum rotational velocity foundfor the most massive galaxies studied with IFS at z ∼ − (F¨orster Schreiber et al. 2018). Wedefine gas with V r > | | km s − and V σ <
250 km s − belonging to a merger system. A system is defined as amerger if there are components with V r > | | km s − or more than one distinct kinematic component. For ex-ample, in the 3C298 system, two galactic disks are foundto be offset by less than 400 km s − . All radial velocityand velocity dispersion measurements are relative to theredshift of the quasar. The redshifts for the individualquasars are calculated in Vayner et al. (2021) and aretaken from the fit to the narrow-line region. For sourceswith no spatially unresolved narrow emission, we use theredshift of the broad-line region. We label quiescent re-gions in the following manner: source name + direction inematics and Energetics of Ionization Gas Table 1.
Archival HST imagingObject Proposal ID Instrument Filter Exposure time(s)3C446 12975 ACS-WFC F814W 22003C298 13023 WFC3-UV F606W 11003C268.4 13023 WFC3-UV F606W 11004C09.17 5393 WFPC2 F555W 21003C9 13945 ACS-WFC F814W 2040
Table 2.
QUART Sample propertiesName RA DEC z L bol L M BH J2000 J2000 (10 erg s − ) (10 erg s − ) M (cid:12)
3C 9 00:20:25.22 +15:40:54.77 2.0199 8.17 ± ± ± ± ± ± ± . ± ± ± + component A or B where A = component associatedwith the quasar, B = component associated with thegalaxy merging with the quasar host galaxy. We followthese with a one or two-word comment about the re-gion. Examples of description words are clump, diffuse,or tidal feature. Where clump referrers to a typical fewkpc in size compact ionized emission typically seen inhigh redshift star-forming galaxies. Diffuse referrers togas that has a surface density of less than typical clumpystar-forming regions. A tidal feature refers to ionized gasassociated with a tidal tail in a merging system, contain-ing both diffuse and clumpy ionized gas morphology.For each dynamically quiescent region, we construct a1D spectrum by integrating over its spaxels. We show anexample of this for 4C09.17 in Figure 1, spectra of dis-tinct regions for the rest of the sources are presented inthe appendix (Figures 11-18). The emission lines in eachspectrum are fit with multi-Gaussian profiles. In theseplots, we also present the outflow regions from (Vayner et al. 2021), to illustrate the location of dynamicallyquiescent regions relative to turbulent regions in thesequasar hosts. From these fits, we derive integrated inten-sity and velocity dispersion that are presented in Tables3 and 5.4.2. Spatially unresolved narrow-line regions
We search for narrow spatially unresolved emission ineach object. To do so, we first subtracted a model of theextended emission from our fits to each emission line inindividual spaxels. We then perform aperture photom-etry on the spatially unresolved emission and extract aspectrum. The emission lines are fit with multiple Gaus-sian profiles. The fluxes of the narrow emission ( σ < − ) lines from unresolved regions are presented inTable 4. For sources where no unresolved narrow emis-sion line is detected, we place a 1 sigma upper limit onthe line flux. Based on the average angular resolution of Vayner et al. −2000−1500−1000 −500 0 500 1000 1500 2000−2−1012345 ∆ λ F λ [ e r g / s / c m ] S/E component A outflow [OIII] −2000−1500−1000 −500 0 500 1000 1500 2000−1.0−0.50.00.51.01.52.0
S/E component A outflow Hα+[NII]
N E −2000−1500−1000 −500 0 500 1000 1500 2000−0.4−0.20.00.20.40.60.81.01.2
SW component A [OIII]
Total[OIII] 500.7nm −2000−1500−1000 −500 0 500 1000 1500 2000−2−1012345
SW component A Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm −2000−1500−1000 −500 0 500 1000 1500 2000−0.50.00.51.01.52.02.53.03.5
W component B clumps [OIII]
Total[OIII] 500.7nm −2000−1500−1000 −500 0 500 1000 1500 2000−0.20.00.20.40.60.81.01.21.4
W component B clumps Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm −2000−1500−1000 −500 0 500 1000 1500 2000 km/s −0.20.00.20.40.60.81.0
W component B d ffuse [OIII]
Total[OIII] 500.7nm −2000−1500−1000 −500 0 500 1000 1500 2000 km/s −1.0−0.50.00.51.01.52.02.53.03.5
W component B d ffuse Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm
Figure 1.
On the left, we present the spectra of distinct regions and fits to individual emission lines for the 4C09.17 system. Onthe right, we present the three-color composite where H α is color-coded to red, [O III ] to green and [N II ] to blue. The contouroutlines the spatial location of the region. The bar on the right in each stamp represents 1 (cid:48)(cid:48) or approximately 8.6 kpc at thesystem’s redshift. about 0.1 (cid:48)(cid:48) , the unresolved narrow line emitting regions’sizes are < NEBULAR EMISSION LINE DIAGNOSTICSAND SOURCES OF GAS EXCITATIONIn this section, we explore the photoionization mech-anism in all distinct regions of each quasar host galaxy.The Baldwin, Phillips & Terlevich (BPT) diagram isused to differentiate between different gas photoioniza-tion sources (Baldwin et al. 1981). Here, we use thelog([O
III ]/H β ) and log([N II ]/H α ) line flux ratios todistinguish heating from young stars, AGN, and shocks.To construct the BPT diagram for our sources, we in-tegrated each emission line over the same velocity width(∆V) and velocity offset relative to the redshift derivedfrom the [O III ] emission line at each spaxel. We inte-grate the maps relative to [O
III ] since it is typically thebrightest emission line in any given spaxel. The highersignal-to-noise [O
III ] emission line leads to a smallerspaxel-spaxel variation in the radial velocity and disper-sion maps, creating a more consistent log([O
III ]/H β )and the log([N II ]/H α ) ratio between neighboring spax-els. We find that for the entire sample, the standarddeviation on the log([O III ]/H β ) ratio decreases by 0.2 dex compared to when integrating the cubes relative tothe H α line.A resolved BPT diagram allows us to investigate thesource of ionization throughout each quasar host galaxy.Due to sensitivity and, in some cases, wavelength cov-erage, we cannot create an integrated emission-line mapfor H β on a similar scale to H α , [O III ] , or [N II ] maps.For our BPT diagrams, we construct our H β map byassuming case B recombination (H β =H α /2.86) with agas temperature of 10 K and an electron density of 10 cm − . Assuming other recombination cases and ISMconditions with reasonable temperatures and densitieswould not change our results by a significant amount asthe ratios between H β and H α would only change atmost by a factor of ∼ α and H β we find a maxi-mum V band extinction of 1 mag, however in most cases,line ratios consistent with case B recombination. In re-gions where gas extinction is present, the log([O III ]/H β ) ratios are preferentially lower.Only spaxels where at least H α and [O III ] were de-tected are analyzed and presented here. Typically [N II ] inematics and Energetics of Ionization Gas Table 3.
Fluxes of distinct “dynamically quiescent” regions in individual sourcesSource Region F [OIII] F H α F [NII] − erg s − cm − − erg s − cm − − erg s − cm − ±
20 65 ± ±
2N component B 127 ±
13 40 ± ± ± ± ± ± ± ± ± ± ± ±
25 51 ± ±
17C 1354+2552 component A 46 ± ± ± ± < ±
65 188 ±
20 65 ± ± ± ± ± ± ± ± ± ± ± ± ± ± < ±
10 48 ± ± Table 4.
Fluxes of spatially unresovled narrow emission line regions in indi-vidual sourcesSource F [OIII] F H α F [NII] − erg s − cm − − erg s − cm − − erg s − cm − ± ± ± ±
70 239 ±
20 76 ± ±
200 102 ±
10 3.5 ± ± ± ± ± ± < ± < < < < < < < < < . < < Vayner et al. is detected in far fewer spaxels compared to H α and[O III ] . For spaxels where only H α and [O III ] are de-tected, we calculate a limit on [N II ] by integrating a skyspectrum over the same velocity width as [O III ] at theexpected wavelength location of [N II ] . In Figure 2 weplot the ratios from each spaxel. Diamonds are regionswhere [N II ] , H α , and [O III ] were detected, and trian-gles are regions where only H α and [O III ] were detectedwith a limit on the [N II ] flux. A total of 3160 spaxelsare plotted corresponding to 21 distinct galactic regions.For each distinct regions identified in section 4.1 andfrom (Vayner et al. 2021) we over plot their line ratiosand label them with a star. Individual spaxels typicallyhave high uncertainties in their ratios but tend to clustertogether on the BPT diagram. Integrating over distinctregions and re-calculating the ratios from a high SNRspectrum confirms that region’s true line ratio.To conserve space, we do not over-plot the error barson points from individual spaxels in Figure 2, we onlyshow the error bars of ratios computed for integratedvalues of the distinct regions. In Figure 3, we plot pointsof individual spaxels along with the error bars.5.1. Ionization Diagnostic Models
We find that a large portion of our line ratios valueslies outside the two typical sequences of the BPT dia-gram (Figure 2). At a fixed log([N II ]/H α ) nearly, allvalues are above both the local mixing and star-formingsequence. At a fixed log([O III ]/H β ) value, nearly allvalues are outside the local mixing sequence. A largeportion of points falls between the star-forming and mix-ing sequence, with relatively high log([O III ]/H β ) val-ues. Metallicity, electron density, the hardness of theionization field, and the ionization parameter determinesthe location of the galaxy/region on a given sequence.With changing conditions in the ISM between z=0 andthe median redshift of our sample, the locus of both thestar-forming and mixing sequence can change locations(Kewley et al. 2013a).Galaxies at a fixed stellar mass have lower metallicitiesat high redshift compared to galaxies today (Yuan et al.2013). Near the peak of the star formation rate den-sity at z ∼ −
3, the ISM conditions and star formationhistories of star-forming galaxies may differ from localsystems. Star formation appears to happen in denser en-vironments in the distant Universe with higher electrondensities (Sanders et al. 2016), akin to conditions seen inlocal ULIRGs. According to Steidel et al. (2014); Stromet al. (2017) ISM in high redshift galaxies experiencesharder ionization at a fixed N/O and O/H abundancesthan z=0 star-forming galaxies. On the other hand,galaxies at higher redshift have elevated N/O rations (Shapley et al. 2015). Taken together, Kewley et al.(2013a) shows that such changes in ISM conditions canalter the location of the star formation sequence betweenz=0 and z=2.5. Notably, the combination of harder ion-ization, electron density, and ionization parameter canshift the locus of the star-forming sequence to higherlog([N II ]/H α ) and log([O III ]/H β ) values. It appearsthat UV/emission line selected galaxy samples tend toshow a more significant shift from the SDSS star for-mation locus, as evident in a large sample of 377 star-forming galaxies explored by Strom et al. (2017). Nearlyall their galaxies have a higher log([O III ]/H β ) value ata fixed log([N II ]/H α ) compared to local galaxies stud-ied in SDSS. Various galaxy selection techniques maylead to samples of galaxies with inherently different ion-ization properties. However, the overall conclusion fromstudying star-forming galaxies in the distant Universe isthat the line ratios of these systems lie on different starformation locus compared to the local Universe.Changes in the ISM conditions of distant galaxies mayalso lead to changes in the location of the mixing se-quence. Kewley et al. (2013a) and Groves et al. (2006)show that for galaxies with lower metallicities, the mix-ing sequences shifts to lower log([N II ]/H α ) values withrelatively small changes in the log([O III ]/H β ) value.We explore the various evolutionary models of the star-forming and mixing sequence with redshift and ISM con-ditions proposed by Kewley et al. (2013a). The best fitmodel to our sample is the one where the ISM of highredshift galaxies have more extreme condition (higherelectron density, harder ionization field, and larger ion-ization parameters) and the metallicity of the gas pho-toionized by the quasar is at a lower metallicity com-pared to the gas ionized by local AGN in the SDSS sam-ple. The median log([N II ]/H α ) value is about 1.0 dexlower than that of the mixing sequence at z=0. If theprimary source in the shift of the mixing sequence fromz=0 to z=1.5-2.5 is a change in the gas phase metallicity,then the gas photoionized by the quasar in our samplehas a metallicity a 1/2-1/5 of that in narrow line regionsof z=0 AGN on the Kewley & Dopita (2002) metallicityscale. One of the consequences of the shift in the mix-ing sequence is that it becomes harder to distinguishbetween gas photoionized by AGN vs. star formation,especially in systems with potentially multiple ioniza-tion sources. Changes in the photoionization conditionlikely also play a role in the offset from the local mixingsequence. In a sample of local type-1 quasars, (Stern &Laor 2013) found that systems with higher bolometricluminosities and higher Eddington ratios are systemat-ically offset to lower log([N II ]/H α ) ratios. inematics and Energetics of Ionization Gas −2.0 −1.5 −1.0 −0.5 0.0 0.5 log([NII]/Hα) −1.0−0.50.00.51.01.5 l og ( [ O III ] / H β ) Model0 z=2.5 Mixing Se.2ence Lower Bo2ndar5 K13a model: 4z=2.5 Mixing Sequence Upper Boundary K13a model: 4z=2.5 Upper Boundry Star Forming Sequence K13a model: 4z=2.5 Lower Boundry Star Forming Sequence K13a model: 4 z=0 BPT sequences
Regions
Regions
Figure 2.
Line ratio diagnostics of individual resolved distinct regions. In grey, we plot the line ratios of individual spaxelswhere at least [O
III ] and H α was detected at an SNR >
3. Uncertainties on these line ratios are generally large; hence, we alsointegrate over all spaxels in individual regions to increase the SNR and lower the uncertainties on the line ratios. We showregion-integrated line ratios with the colored symbols where each object has the same symbol, and each region has a differentcolor. The names of the distinct region are present in the lower-left corner, and these match the names given in Table 3. Wepresent the evolutionary models of the mixing and star-forming sequence with red and green curves from Kewley et al. (2013a).We show the upper limit of a sequence with a straight line and the lower boundary of each sequence with a dashed curve. Tealcurves represent the bounds of the two sequences where the majority of the line ratio in low redshift galaxies fall. Our lineratios are consistent with a model where gas photoionized by the quasar is denser, has lower metallicity, and experiencing harderionization compared to the gas photoionized by AGN in nearby galaxies.
Most of the gas in our quasar host galaxies lies inthe mixing sequence where the gas is photoionized by acombination of quasar ionization and radiative shocks.In Figure 3, a significant fraction of the points in indi-vidual objects match the predicted location of radiativeshocks on the BPT diagram. The radiative shock mod-els assume solar metallicity and a preshock density of215 cm − . With the present data, it is difficult to dis-tinguish the percentage of photoionization from shocksvs. AGN. However, given the overlap of both line ratiosand kinematics with shock models, we cannot rule themout to contribute to gas photoionization.A number of our distinct regions appear to have lowlog([N II ]/H α ) values ( < σ <
250 km s − ). Morphologically these re-gions appear to be clumpy in their H α maps reminiscent of typical star-forming regions in galaxies at z >
1. Theline ratios of these points do not coincide with regions offast or slow shocks photoionization on the BPT diagram(Allen et al. 2008; Newman et al. 2014). Archival HSTdata of 3C9, 3C298, 4C09.17, 3C268.4, 3C446 all show-case that the dynamically “quiescent” regions in thesegalaxies have clumpy morphology in rest-frame UV con-tinuum data, similar to those of star-forming galaxies atthese redshifts. In Figure 4, we overlay the H α emis-sion from dynamically quiescent regions onto archivalHST observations at rest-frame UV wavelength. Com-bining these clues suggests that the quasar does not en-tirely photoionize the gas in these regions. The elevatedlog([O III ]/H β ) in these regions compared to local anddistant star-forming regions may be from the mixingof photoionization from massive stars and the quasar.0 Vayner et al. −1.0−0.50.00.51.01.5 3C3183C318 4C04.81 3C9 3C446−0.80.00.8 l og ( [ O III ] / H β ) log([NII]/Hα) −0.80.00.8 4C09.17 −2.0−1.5−1.0−0.5 0.0 0.53C298 Limit on [NII]Hα,[NII] & [OIII] detected
Figure 3.
We present line ratio diagnostics for spaxels in each source where at least [O
III ] and H α were detected at SNR greatthan 3. We show the uncertainties on the line ratio, which were omitted from figure 2 to conserve space. The dashed red line ineach panel shows the theoretical separation between gas photoionized by star formation and AGN or shocks from Kewley et al.(2013a). Points above the line are photoionized by the quasar, while regions below are photoionized by O and B stars. Solidblack mesh represents the location of radiative shocks from the astrophysical plasma modeling code MAPPINGS V (Alarie &Morisset 2019). The shock model uses solar abundances from Gutkin et al. (2016). Either shocks or the quasar photoionizesthe majority of the gas within our systems. There is some evidence for this based on the morphol-ogy of the ionized gas and their respective log([O
III ]/H β ) ratios. For example, in 4C09.17, we see that more dif-fuse emission with low-velocity dispersion tends to havea higher log([O III ]/H β ) value compared to clumpier re-gions where there is evidence for recent star formationactivity and potentially more shielding from the quasarradiation field.Using the empirical star formation rate H α luminosityrelationship from Kennicutt (1998), we convert the H α luminosity of the distinct extended quiescent regions tostar formation rates. Most likely, the majority of theseregions are photoionized by a combination of AGN andstar formation, hence the derived star formation ratesare upper limits. Regions “3C298 SE component BTidal feature” and “4C09.17 W component B clumps”have line rations most consistent with photoionizationby O/B stars, the star formation rates derived in theseregions are closer to their actual value. To partially ad-dress the contribution from AGN to photoionization indynamically quiescent regions, we also derive star for-mation rates only using spaxles that fall within the line ratio error inside the star-forming region of the BPTdiagram based on the Kewley et al. (2013b) maximumseparation between star formation and AGN photoion-ization. Generally, we find lower (1/2 - 1/10) star for-mation rates when using the BPT diagram criteria. Wealso measure the metallicities of these regions using thePettini & Pagel (2004) empirical gas-phase metallicity -log([N II ]/H α ) relationship. Given that log([N II ]/H α )is elevated in the presence of an AGN/quasar ionizationfield, the metallicities for the majority of the regionsare also upper limits. We also calculate the metallicityusing theoretically derived chemical abundance calibra-tion for narrow-line regions of AGN (Storchi-Bergmannet al. 1998). We present quantitative values of theseregions in Table 5, where we show the H α luminosityof each quiescent region, along with the star formationrate, metallicities, and velocity dispersion. Since nearlyall of the unresolved narrow line regions are consistentwith quasar photoionization, we do not include them inTable 5 with the exception of 3C318. In this object, theline ratios are consistent with star formation, indicatinga nuclear starburst on scales < inematics and Energetics of Ionization Gas (cid:12) yr − for the four objects (4C05.84,4C04.81, 4C57.29, and 7C1354) where no strong narrownuclear emission was detected. SMBH-GALAXY SCALING RELATIONSHIPSIn this section, we place our galaxies on the veloc-ity dispersion and galaxy mass vs. SMBH mass plots,comparing their locations to the local scaling relations( M • − σ and M • − M ∗ ). We calculate the SMBH massesfrom the broad H α luminosity and line width using themethodology presented in Greene & Ho (2005). TheSMBH masses span a range of 10 . − . M (cid:12) . The ve-locity dispersions are taken from dynamically quiescentregions, while the galaxy masses are calculated from thevirial equation and from modeling the radial velocity oftargets with rotating disks and extracting a dynamicalmass. 6.1. Host Galaxy Velocity Dispersion
We identify several dynamically quiescent regionswithin most of the quasar host galaxies in our sam-ple. These regions show lower log([N II ]/H α ) line ratiosand typically have clumpy morphology, reminiscent ofthe general star-forming regions seen in nebular emis-sion and UV continuum in high redshift galaxies. Inmost galaxies, these regions lie away from any galactic-scale outflows. Hence their observed dynamics could bea probe of the galactic gravitational potential. Theseregions can be used to measure the velocity dispersionof our quasar host galaxies. In combination with themeasured black hole masses, we can compare them tothe local scaling relation between the SMBH mass andthe velocity dispersion of the galaxy/bulge. In Figure5, we plot the mass of the SMBH presented in Table 2against the velocity dispersion of distinct quiescent re-gions measured with the H α line. Also, we include thevelocity dispersion measured from CO (3-2) emission for3C 298 from Vayner et al. (2017). We find a significantoffset from the local scaling relation between the SMBHmass and the velocity dispersion of the galaxy/bulge( M • − σ ) (G¨ultekin et al. 2009; McConnell & Ma2013). To address the issue that the velocity dispersionmay be systematically lower in dynamically quiescentregions offset from the quasar (3C446) or regions wherethe surface area of the narrow emission is significantlylower than the outflow (4C09.17, 3C298), we recalculatethe velocity dispersion in a larger aperture that includes outflows and narrow emission. We see no strong system-atic difference in the velocity dispersion of the narrowgas. The source integrated narrow velocity dispersionfor 3C298, 3C446 and 4C09.17 are 100.7 ± ± ± − , respectively. In the nearbyuniverse, the velocity dispersion is typically measuredinside the effective radius of the galactic bulge. Thedifference within our observations is that we do notknow the bulges’ true sizes for our galaxies as we haveno way to measure them with current data and tele-scope/instrument technology. However, the extents ofthe dynamically quiescent regions are in the range ofthe effective radii for bulges in massive (10 . − . M (cid:12) ) galaxies studied in the CANDELS survey (Bruce et al.2014).There have seen numerous discussions in the literatureabout whether the velocity dispersion measured fromgas traces the stellar velocity dispersion. The gas andstars might not have the same uniform distribution, andwinds can broaden the nebular emission lines. Further-more, the line of sight absorption and emission lines fromwhich the velocity dispersion is calculated are luminos-ity weighted subject to galactic dust extinction. Becauseof the different light distribution between stars and gas,the measured velocity dispersion can differ. These dif-ferences can lead to increased scattering in any corre-lation between σ ∗ and σ gas . Data-sets that spatiallyresolve the gas and stellar components and have enoughresolving power to separate multi-component emissionfrom different regions (e.g., outflowing/in-flowing gasvs. galactic disk rotation) are important when mak-ing a comparison between σ ∗ and σ gas . In Bennertet al. (2018), for a large sample of local AGN, theyfind when fitting a single Gaussian component to the[O III ] emission line, they can overestimate the stellarvelocity dispersion by about 50-100%. Only by fittingmultiple Gaussian components to account for both thenarrow core and the broader wings of the [O
III ] lineprofile can they adequately match the velocity disper-sion from the narrow component of the [O
III ] line tothat of the stellar velocity dispersion. For their entiresample, the average ratio between the velocity dispersionof narrow Gaussian component and the stellar velocitydispersion is ∼
1. The 1 σ scatter on the ratio between σ [ OIII ] ,narrow and σ ∗ is about 0.32 with a maximummeasured ratio of about a factor of 2 which translatesto a scatter in ∆ σ = σ [ OIII ] − σ ∗ of 43.22 km s − witha maximum difference of about ±
100 km s − . How-ever, only a few sources show such drastic velocity dif-ferences ( ∼ .
5% of the entire sample, 82% of the sourcesshow | σ [ OIII ] − σ ∗ | <
50 km s − ). When fitting for the M • − σ relationship with the narrow [O III ] emission2
Vayner et al.
Table 5.
Star formation rates and metallicities of distinct dynamically quiescent regionsSource Region L H α SFR a L H α BPT
SFR
BPTb σ gas erg s − M (cid:12) yr − erg s − M (cid:12) yr − PP04 SB98 km s − ± ±
16 0.20 ± ± ± ± ±
10 0.06 ± ± ± ± ± ± ± ±
3W B clumps 0.35 ± ± a a ± ± ± a a ± ± ± ± ± ±
57C 1354+2552 A 0.37 ± ± ± ± < ± d ± ± a a ± ± ± b a < < ± ± ± ± ± ± ± ± ± ± ± d ± ± ± ± < < ± ± ± ± ± ± a Star formation rate derived using the H α luminosity of the entire distinct quiescent region b Star formation rate derived using spaxels that fall within the star formation sequence on the BPT diagram. c Value indistinguishable from the integrated value over the entire dynamically quiescent region d Tidal tail as a proxy for stellar velocity dispersion, the resultantfit agrees with that of quiescent galaxies reverberation-mapped AGNs. These results indicate that for the sam-ple as a whole Bennert et al. (2018) finds that both thestars and gas follow the same gravitation potential.Given the Bennert et al. (2018) results that demon-strates that the gas velocity dispersion can be used asa proxy for the stellar velocity dispersion, we follow asimilar analysis using our IFS data sets to explore thelocation of our galaxies relative to the M • − σ relation athigh redshift. We attempted to the best of our ability toseparate regions that contain galactic scales winds fromthose with more quiescent kinematics both spectrallyand spatially with OSIRIS. Hence similar to Bennertet al. (2018) we think that the measured velocity disper-sions in quiescent regions are good tracers of the galacticpotential on average. Throughout the paper we use thenarrow velocity dispersion of [O III ] and H α emissionlines of dynamically quiescent regions as a proxy for thestellar velocity dispersion. We still find a significant off-set for our sample after applying the observed scatter inthe difference between σ ∗ and σ gas . This is also truewhen applied to the more distant quasar host galaxies studied with 158 µ m [CII] emission. In nearby galaxies,there is a dependence on the velocity dispersion withthe radius from the galaxy center (Bennert et al. 2018;Ene et al. 2019). However, based on the local galaxy ob-servations, the velocity dispersion is unlikely to increaseby ∼
200 km s − that is necessary to bring the galaxiesonto the local scaling relations.Using N-body smoothed-particle hydrodynamics sim-ulations Stickley & Canalizo (2014) examines how thestellar velocity dispersion evolves in a binary galaxymerger. At various stages in the merger (e.g., a closepassage, nucleus coalescence), they measure the stellarvelocity dispersion along 10 random lines of sight. Neareach close passage and during coalescence, they find thatthe scatter on the velocity dispersion significantly in-creases from ∼ −
11 km s − to about 60 km s − withthe average velocity dispersion a factor of ∼ inematics and Energetics of Ionization Gas Figure 4.
Detection of dynamically quiescent regions in archival
Hubble Space Telescope observation. In the background,we show PSF-subtracted images of rest-frame UV emission in the quasar host galaxy. Overlaid in contours is the extendedH α emission of the dynamically quiescent regions detected with OSIRIS. Note the similarities in both morphology and extent,indicating massive young stars photoionize at least a portion of the gas. The bar represents a spatial scale of 1 (cid:48)(cid:48) or about 8.5kpc. certainty on the velocity dispersion of 60 km s − giventhat the majority of our mergers are near coalescence ora close passage (∆ R <
10 kpc).It should be noted that this is near the maximum scat-ter seen in the simulations on σ . These simulations alsofind that for merging galaxies at their maximum sepa-ration, the measured velocity could be a factor of ∼ . ∼ a few km s − ). However, we still apply an additional60 km s − uncertainty in these regions.Even after these corrections are made to approximatethe stellar velocity dispersion from the [O III ] emissionlines in our sample, we still find that all of our sys-tems are offset from the local scaling relation betweenthe mass of the SMBH and the velocity dispersion of the bulge/galaxy. Given that we are dealing with rela-tively small sample size, we performed statistical tests toconfirm the offset between the local scaling relation andour sample. We measure the offset between the observedand predicted velocity dispersion for the SMBH mass ofour systems for each object. We use the local scalingrelation fit from (McConnell & Ma 2013), and H α mea-sured SMBH masses. We construct a data set consistingof velocity differences. From bootstrap re-sampling ofthe velocity difference data set, we find that the averageoffset of 188.7 km s − is significant at the 3.25 σ level.Using Jackknife re-sampling similarly, we find that theoffset is significant at the 3.3 σ level with the 95% con-fidence intervals of 154.4 km s − to 223.0 km s − on thevelocity dispersion offset. Performing similar statisticaltests on the Decarli et al. (2018) sample, we find an av-erage offset of 178.8 km s − with a significance of theshift at 2.7 σ and 2.8 σ for Jackknife and bootstrap re-sampling, respectively from the local relationship. Wealso measure the offsets of massive BCGs in the localUniverse from the M • − σ relationship. Using a two-sided Kolmogorov-Smirnov test, we can ask if the ob-served offsets of the local and high redshift data sets aredrawn from the same continuous distribution. We find a4 Vayner et al.
Figure 5.
The location of our galaxies on the velocity dispersion vs. SMBH mass plot compared to the local M • − σ relationship.We use the narrow H α emission line velocity dispersion of dynamically quiescent regions as a proxy for the stellar velocitydispersion. Red stars are the measurements from our sample, where we measure the velocity dispersion from the narrow H α emission line. We measure the black hole masses using the broad H α line from the broad-line-region discussed in section 3. Thetwo blue stars represent the velocity dispersion measured in the disk of the host galaxy of 3C 298 and the tidal tail feature21 kpc away from the quasar (Vayner et al. 2017). Blue circles are quasars from the Shields et al. (2006) sample, where theymeasure the velocity dispersion from CO emission lines. The yellow points are from quasars at z >
6, where they measure thevelocity dispersion from the 158 µ m [CII] emission line (Decarli et al. 2018). Green points represent the local sample of all thebright cluster galaxies with SMBH greater than 10 M (cid:12) taken from McConnell & Ma (2013). The blue curve represents thebest fit to the entire galaxy sample from McConnell & Ma (2013) with the blue shaded region represents the intrinsic scatter onthe M • − σ relationship from the fit while the green curve is the fit to the sample studied in G¨ultekin et al. (2009). We find asignificant offset between the galaxies in our sample and local BCG and the general local M • − σ relationship. This large offsetindicates that the host galaxies appear to be under-massive for their SMBHs. p-value of 5.7 × − , indicating that the two populationsare not drawn from the same distribution. Applying theKolmogorov-Smirnov test to the velocity dispersion off-sets from our sample and in the higher redshift quasars,we find a p-value of 0.84, indicating that these two datasets could be drawn from the same continuous distribu-tion. We find similar results by comparing the Shieldset al. (2006) sample at z ∼ DYNAMICAL MASS MEASUREMENTSWe can also test whether these systems lie off the lo-cal scaling relationship between the SMBH mass andthe dynamical mass of the bulge/galaxy. First by usinga virial estimator for the dynamical mass of the galaxy M virial = C σ rG where C=5 for a uniform rotating sphere(Erb et al. 2006b). We assume 7 kpc for the radius,which is the median effective radius of massive quiescentgalaxies in the local Universe (Ene et al. 2019). Here σ is derived from a Gaussian fit to the integrated spectraover the distinct region. For galaxies with multiple dis-tinct regions, we derive two or more dynamical massesas there may be a dependence on the velocity dispersionas a function of position with the galaxy. For systemsin a clear merger, the galactic component belonging tothe quasar is used to estimate the dynamical mass sincewe are interested in the correlation between the SMBHand the velocity dispersion of the quasar host galaxy.For systems with velocity shear in the 2D radial ve-locity map, we fit a 2D inclined disk model to the kine- inematics and Energetics of Ionization Gas Figure 6.
We present the location of individual galaxies compared to the local scaling relation between the mass of the SMBHand mass of the galaxy/bulge shown with a blue curve. Blue points represent systems with virial dynamical masses. Red pointsrepresent systems where we calculate the dynamical mass by modeling the radial velocity maps with an inclined disk model.Gray points show the location of galaxies at z > .
5, with lower SMBH masses and lower AGN luminosity compared to oursample. The blue curve represents the local scaling relationship as measured in McConnell & Ma (2013), with the shaded regionrepresenting the intrinsic scatter. We find the majority of our points are offset from the local scaling relationship, outside theobserved scatter. matics data to measure the dynamical mass. The modelis a 2D arctan function V ( r ) = 2 π V max arctan (cid:16) rr dyn (cid:17) , (1)where V(r) is rotation velocity at radius r from the dy-namical center, V max , is the asymptotic velocity, and r dyn is the radius at which the arc-tangent functiontransitions from increasing to flat velocity. The mea-sured line-of-sight velocity from our observations relatesto V(r) as V = V + sin i cos θV ( r ) , (2)where cos θ = (sin φ ( x − x )) + (cos φ ( y − y )) r . (3)Radial distance from the dynamical center to each spaxelis given by r = (cid:114) ( x − x ) + (cid:16) y − y cos i (cid:17) , (4)where x , y is spaxel location of the dynamical center,we quote the value relative to the centroid of the quasar, V is velocity offset at the dynamical center relative tothe redshift of the quasar, φ is position angle in spaxel space, and i is the inclination of the disk. V max is notthe true “plateau” velocity of the galaxy’s disk. V max can have arbitrarily large numbers, especially when r dyn is very small (Courteau 1997). To fit the data we use theMCMC code emcee . We construct the model in a gridwith a smaller plate scale than the observed data, whichgets convolved with a 2D Gaussian PSF with an FWHMmeasured from the quasar PSF image. The image is thenre-sized to the plate scale of the data. We construct thepriors on each of the seven free parameters. The prioron V max is 300 < V max < − the prior onboth x , y is the boundary of the FOV of the imagedarea, the prior on the position angle is 0 < φ < π , theprior on the inclination angle is 0 < i < π/
2, the prioron the radius is 0 . < r dyn <
10 pixels and the prioron V is − < V <
100 km s − . We then samplethis distribution with emcee . We initialize 1000 walkersfor each free parameter using the best fit values from leastsquares fitting as the starting point, with a smallrandom perturbation in each walker. We run MCMCfor 500 steps starting from the perturbed initial value.The best-fit parameters, along with their confidence in-tervals, are presented in 6 for the quasar host galaxiesof 7C 1354+2552, 3C9. For 3C 298 we do not see thedisk in the ionized emission with the OSIRIS data, it issolely detected in CO (3-2) observations from ALMA,6 Vayner et al.
Table 6.
Best fit values for each inclined disk model param-eterParameters 7C 1354+2552 3C9 3C298 V max [km s − ] 449.67 +0 . − . +23 . − . +65 − x [kpc] -2.37 +0 . − . +2 − +0 . − . y [kpc] -0.93 +0 . − . -4.8 +1 . − . +0 . − . φ [ ◦ ] 75.68 +0 . − . +3 . − . +1 . − . i [ ◦ ] 47.6 +0 . − . +5 . − . +6 . − . r [kpc] < +0 . − . +0 . − . V [km s − ] -93.9 +1 . − . -9.22 +30 . − . -13.0 +3 . − . ∆ v obs / − ] 309.84 ± ± ± σ [km s − ] 61.3 ± ± ± here we present the best fit values from Vayner et al.(2017). Also, we present ∆ v obs /
2, the average betweenthe maximum and the minimum velocity along the kine-matic major axis as determined by the position angle( φ ). We also present the intrinsic velocity dispersion( σ ), measured along the kinematic major axis, towardsthe outskirts, away from the steep velocity gradient nearthe center of the disk.In addition, we measure V rot /σ to gauge whetherthese systems are dynamically supported by rotation ordispersion. We measure a value of 6.8 ±
1, 2.7 ± ± M ( R ) = 2 . × rV r / sin( i ) (5)Where V r is the radial velocity, i is the inclinationangle from the disk fit. For the radial velocity we use∆ v obs /
2. Similarly, we assume a radius that is the me-dian value of nearby BCGs (7.1 kpc). The selected ra-dius should give us an absolute upper limit on the dy-namical mass of the galaxy/bulge as this radius is muchlarger than the typical size of a galactic bulge at thisredshift and is larger than the observed extent of thegalactic disks. The reason for choosing a larger radius isto address the case where the quasar host galaxy extendsto a larger radius and is not captured in our OSIRIS ob-servations because they are not sensitive enough to lowsurface brightness emission at larger separation from thequasar. Virial and dynamical masses are presented in7. However, it is not guaranteed that the extent of the
Table 7.
Virial and dynamical mass values.Source Virial Mass Disk-fit Dynamical mass × M (cid:12) × M (cid:12) ± ± ± ± ± ± ± a ± ± ± ± ± ± a computed from CO 3-2 velocity dispersion ionized gas will match the stellar. We attempted tomeasure the size of the stellar continuum from the HSTobservations but were unsuccessful. Using the Galfitpackage, we were unable to constrain the radius due tothe sources’ complex morphologies and the increased in-ner working angle due to the quasars’ brightness andsaturated counts in the HST observations.Due to the limited sensitivity of OSIRIS to lower sur-face brightness emission, we are missing an accuratemeasurement of the plateau velocity for the galacticdisks at large separations from the quasar. Hence, ourfitting routine is unable to constrain V max for 3C9 and7C 1354. Also, it appears that the turn over radius isvery small for these two systems, smaller than the reso-lution element of our observations. For this reason, weare unable to constrain the turn over radius, and weonly provide a limit. For 7C 1354, there is a degeneracybetween the maximum velocity, turn over radius, andinclination; thus, the values that we provide are thosethat give the smallest velocity residual.Using the measured virial and disk fit dynamicalmasses and the SMBH masses, we can now compare ourgalaxies to the local M • − M ∗ relationship. Not only arethese galaxies offset from the local M • − σ relationship,but we also find that these galaxies are on an averageoffset from the local M • − M ∗ relationship. The galax-ies need about an order of magnitude of stellar growth ifthey are to evolve into the present-day massive ellipticalgalaxies.We note that we have used two different methods forexploring the scaling relationship for galaxy mass vs.SMBH. Both the gas velocity dispersion method anddynamical measurement imply that the SMBH is over- inematics and Energetics of Ionization Gas −1.5 −1.0 −0.5 0.0 0.5 1.0 1.5ΔarcsecΔ["]−3−2−10123 −1.5 −1.0 −0.5 0.0 0.5 1.0 1.5ΔarcsecΔ["]−3−2−10123 −1.5 −1.0 −0.5 0.0 0.5 1.0 1.5−3−2−10123 −400−300−200−1000100200300400 Figure 7.
Fitting an inclined disk model to the radial velocity map of the 3C9 quasar host galaxy. Far-left we plot the isolatedradial velocity structure belonging to the quasar host galaxy of 3C9, middle left shows the best fit model overlaid as contourson top of the radial velocity map, middle right is the best fit model. On the right, we plot the residuals. Larger blue-shiftedresiduals at − (cid:48)(cid:48) south from the quasar are from the outflow (3C9 SE component A outflow A). −1.0 −0.5 0.0 0.5 1.0ΔarcsecΔ["] −1.0 −0.5 0.0 0.5 1.0ΔarcsecΔ["] −1.0 −0.5 0.0 0.5 1.0 −400−300−200−1000100200300400 Figure 8.
Fitting an inclined disk model to the radial velocity map of the 7C1354 quasar host galaxy. Far-left we plot theisolated radial velocity structure belonging to the quasar host galaxy of 7C1354, middle left shows the best fit model overlaidas contours on top of the radial velocity map, middle right is the best fit model, and on the right, we plot the residuals. massive compared to their host galaxies when exploringthe local scaling relationship. It will be important tofurther compare these methods with larger samples, aswell as future observations with the James Webb SpaceTelescope that will be able to directly measure the stellarvelocity dispersion. DISCUSSIONOur survey aimed to study host galaxies of redshift1.4 - 2.6 radio-loud quasars through rest frame nebularemission lines redshifted into the near-infrared.We place distinct regions of each quasar host galaxyon the traditional BPT diagram (log([O
III ]/H β ) vs.log([N II ]/H α ) ). The majority of the points for our8 Vayner et al.
NE −0.4 −0.2 0.0 0.2 0.4∆arcsec ["]−0.4−0.20.00.20.4 −0.4 −0.2 0.0 0.2 0.4∆arcsec ["]−0.4−0.20.00.20.4 −0.4 −0.2 0.0 0.2 0.4−0.4−0.20.00.20.4 V(CO (3→2)) [kms −1 ] −200−160−120−80−4004080120160200 Figure 9.
Fitting an inclined disk model to the radial velocity map of the 3C298 quasar host galaxy. Far-left we plot theisolated radial velocity structure belonging to the quasar host galaxy of 3C298, middle left shows the best fit model overlaid ascontours on top of the radial velocity map, middle right is the best fit model, and on the right, we plot the residuals. sources lie outside the two local sequences (the mixingand star-forming sequence). In section 5, we introduceevolutionary BPT models from Kewley et al. (2013a)that indicate changes in the photoionization and metal-licity conditions of the gas can shift both of the star-forming and mixing sequences. We fit these models toour data and find that the best-fitting model is the onewhere the gas in our quasar host galaxies is at leasttwo to five times less metal-rich compared to the nar-row line regions of nearby (z < < z < > z > . II ]/H α ) value of our sampleseems to be lower than that of Nesvadba et al. (2017a);this could be due to the lower metallicity of our sam-ple. On the other hand, a different approach to howwe compute our line ratios can cause the discrepancy.Nesvadba et al. (2017a) only presents source integratedline ratios, while we explore ratios of distinct regions be-cause we typically have a factor of 5-10 better angularresolution due to adaptive optics and hence can resolvethe different ionized/kinematics structures of our galax-ies. In the majority of our sources, we see significantvariations in log([N II ]/H α ) and log([O III ]/H β ) valuesacross each system, hence why we explore distinct re-gions. Line ratios from integrated spectra that includeregions with various ionization sources and from multi-ple components of a merger system may shift towardshigher log([N II ]/H α ) , and log([O III ]/H β ) values asthe regions photoionized by the quasar/AGN tend tobe brighter. Line ratios of galaxies with lower luminos-ity AGN compared to quasars/radio galaxies studied inStrom et al. (2017) are nearly all outside the local mix-ing sequence. These points overlap with the locationof our line ratios and that of the radio galaxy sample.The MOSDEF survey finds similar results for their AGNsample at a range of bolometric luminosities (Coil et al.2015). The ubiquity of elevated line ratios in host galax-ies of AGN, meaning they are typically above the localmixing or star-forming sequence on the traditional BPTdiagram (log([O III ]/H β ) vs. log([N II ]/H α ) ), indicates inematics and Energetics of Ionization Gas ∼ Star formation and dynamically “quiescent”regions in the host galaxies
In 9/11 quasar host galaxies within our sample, wesee the morphology of clumpy star-forming regions seenin other galaxies at these redshifts. These regionsalso typically show lower velocity dispersion and lowerlog([N II ]/H α ) values. We described them in more de-tail in section 4.1. These regions lie 1 - 21 kpc fromthe quasar and generally do not coincide with the loca-tion of galactic outflows. For sources with available HSTimaging of rest-frame UV continuum, these regions ap-pear bright and clumpy (see Figure 4). Taking thesetwo results together indicates that O and B stars couldphotoionize a non-negligible fraction of the gas in theseclumpy regions. In section 5, we derive an upper limiton their star formation rates and gas-phase metallicities.Taking this together, there is evidence for very re-cent star formation activity in 9/11 quasars within oursample. We find an average star formation rate of 50M (cid:12) yr − for the star-forming regions within our sample.The average dynamical mass of our quasar host galax-ies of ∼ M (cid:12) , indicates that the galaxies sit nearthe galaxy star formation rate - stellar-mass sequence atz ∼ ∼ ∼ Herschel Space Ob-servatory (Podigachoski et al. 2015; Barthel et al. 2017)for 4C04.81, 4C09.17, 3C318, and 3C298. The mostlikely explanation is that the quasar itself could par-tially heat the dust, H α misses a significant fraction ofthe obscured star formation, or the dust traces a differ-ent star formation history. Interestingly for 3C298 and3C318, where both high spatial resolution imaging ofthe dust and H α emission is present, there is a signifi-cant misalignment between the maps. In places wherewe see evidence for recent star formation based on neb-ular emission-line ratios in 3C 298 and 3C 318, Barthelet al. (2018); Barthel & Versteeg (2019) does not seeany dust emission. For the case of 3C 298 in the loca-tion where we see recent star formation traced by H α ,we also detect a molecular reservoir; however, no dustemission is present there. Furthermore, in the placeswhere dust emission exists in the case of 3C 298, themolecular gas at that location is stable against gravita-tional collapse and has been on a time scale longer thanthe propagation of the observed outflow. For the caseof 4C09.17 and 4C04.81, no high-resolution dust mapsare available. The dust emission could originate at anylocation within the ∼ (cid:48)(cid:48) Herschel SPIRE beam, whichtranslates to a physical scale of about 150 kpc. Futurehigh spatial resolution dust and molecular gas emissionmaps are necessary for proper comparison between theobscured and unobscured star formation traces and themolecular gas dynamics.8.2.
Offset from local scaling relations
The majority of our systems appear to be offset fromboth local scaling relationships between the mass of theSMBH and mass and the velocity dispersion of the bulge(see Figures 5, 6). To explain the large offset from thelocal M • − σ and M • − M ∗ relationship, we could invokea significant error in the estimated SMBH masses. Thebolometric luminosities of some of our quasars are fargreater than those used for reverberation mapping in thenearby Universe, which is used in calibrating the singleepoch SMBH mass (Greene & Ho 2005). The SMBH0 Vayner et al. masses would have to be off by 2-3 orders of magnitudeto explain the observed offsets. By assuming that theSMBH grows primarily through gas accretion, we canuse the Eddington luminosity formula to estimate theSMBH mass. Given that our quasars are most likelynot all accreting at or close to the Eddington limit, thisderived mass is effectively a lower limit.M
SMBH , min = L Eddington . × M (cid:12) (6)For the derived bolometric luminosities in Table 2 wefind a range of minimum SMBH of 10 . − M (cid:12) , con-sistent with what we measure from single epoch SMBHmasses using the H α emission line. Hence, ther is likelyno significant error in our measured black hole masses.In Figure 10, we plot the offset from the local scal-ing relation against the redshift of each object from oursample, the local galaxies sample with SMBH > M (cid:12) and higher redshift quasars. Quasars with SMBH > M (cid:12) appear to be offset from the local scaling re-lationship, which indicates that SMBH growth appearsto outpace that of stars in these systems. The SMBHsmay grow rapidly up to a mass of several times 10 M (cid:12) as early as 690 Myr after the Big Bang (Ba˜nados et al.2018), matching in mass to some of the most massiveSMBH seen today. Some galaxies with lower luminos-ity AGN and lower mass SMBH also appear to be offsetfrom the local scaling relation at z > > (cid:12) yr − from z=2 to z=0.In the host galaxy of 3C 298, there is currently in-sufficient molecular gas for the galaxy to grow in stel-lar mass to match the mass predicted by the local scal-ing relationship. Furthermore, the quasar 3C 298 doesnot appears to live in an over-dense environment basedon the number count of galaxies seen with the Spitzerspace telescope imaging data (Ghaffari et al. 2017). Theopen question is, how do these galaxies obtain the stel-lar mass necessary to grow into the massive galaxies wesee today? Are minor mergers responsible for growingthese galaxies? Alternatively, is the accretion of coolgas from the CGM responsible for providing the fuelnecessary for future star formation? The results we findfor the host galaxy of 3C 298 favor the scenario wherecold accretion flows from the CGM will supply most ofthe fuel necessary for future star formation. Anotherscenario could be that the Spitzer observations are tooshallow to see lower mass galaxies. If these systems aregas-rich, they can supply future fuel for star formationfrom merging the gas in their CGM and ISM with thequasar’s host. Indeed in recent hydrodynamical simu-lation (Angl´es-Alc´azar et al. 2017) found that for darkmatter halos with masses > . M (cid:12) majority of themass build up happens from gas accreted from the CGMand transfer/exchange of gas from CGM and ISM of can-nibalized low mass galaxies. These simulations also findthat stellar build-up from dry mergers and just accretionof stars from merging galaxies is not significant to growthe stellar mass of galaxies in massive halos. If this isthe case for the majority of our galaxies, it implies thatthey have enormous amounts of gas inside their CGM.Our results can be in stark contrast to the predictedevolutionary paths of massive galaxies. In today’s theo-retical framework (Di Matteo et al. 2005; Hopkins et al.2008; Zubovas & King 2012, 2014), feedback from theSMBH is predicted to happen once the galaxy reachesthe local M • − σ relationship. However, our systemsare showcasing outflows that are capable of causing ma-jor feedback when the mass of the galaxies is a fractionof their predicted final mass from the local scaling rela-tions. Also, the gas-phase metallicities are far lower thanthose observed in nearby AGN. The kinetic luminositiesfor half of the outflows in our sample are far lower thanthe values predicted in simulations for the bolometricluminosities of our quasars (Vayner et al. 2021). Ionizedoutflows in other samples show similar results, whereabout half the objects lie below the predicted minimumenergy-coupling between the quasar and the outflow of inematics and Energetics of Ionization Gas Figure 10.
Measured offset of galaxies from the local M • − σ scaling relationship (McConnell & Ma (2013), log ( M BH /M (cid:12) ) =8 .
32 + 5 .
64 log ( σ/
200 km s − )). On the y-axis, we quantify the offset as the difference between the observed and predictedvelocity dispersion from the local scaling relation based on the observed SMBH mass. We plot the observed offset from the localscaling relation against the redshift for individual targets. The labels are similar to 5. The shaded blue region represents theintrinsic scatter in the M • − σ relationship for black holes with a mass of 10 . M (cid:12) . There is an overall offset for galaxies withmassive SMBH at z > M • − σ relationship. We find no statistically significant difference in the offset betweenany of the high redshift samples, while there is a statistically significant offset from the local BCG points (green). ∼ M • − σ relationship. We might need to recon-sider our theoretical framework for massive galaxy for-mation, where the gas is not cleared from the galaxy in a single “burst” of feedback once the galaxies reachthe M • − σ relationship. Instead, the SMBH grows firstin massive dark matter haloes, followed by a delayedgrowth of the host galaxy with regulatory feedback fromthe SMBH and near-continuous accretion of gas from theCGM and nearby satellite galaxies. In such a scenario,the coupling efficiency might be lower per outflow event,compared to a single burst model where a single outflow-event clears all the gas. At later times, maintenancemode feedback from jets can heat the CGM, preventinggas from cooling and accreting onto the galaxy, keepingthe galaxies “quenched”. CONCLUSIONSWe have conducted a near diffraction-limited surveyof 11 quasar host galaxies to study the distribution,kinematics, and dynamics of the ionized ISM using theOSIRIS IFS at the W.M. Keck Observatory. This surveypaper aimed to understand the source of gas ionization,the physical and chemical conditions of the ISM and to2
Vayner et al. estimate the masses of radio-loud quasar host galaxiesat z ∼
2. We detected extended emission in all objects onscales from 1-30 kpc and found that: • The AGN photoionizes the majority of the ex-tended gas. A significant fraction of emission-line ratios are found to reside between the twosequences on the traditional BPT diagram. Byapplying evolutionary models of the mixing andstar-forming sequence from z=0 to z=2.5, we findthat the gas within our systems is denser and haslower metallicity compared to the gas photoionizedin local AGN. • In 9 objects, we find dynamically quiescent re-gions, with lower average log([O
III ]/H β ) ratios.For systems where Hubble Space Telescope imag-ing is available, their morphologies are consistentwith clumpy star-forming regions commonly ob-served in the distant Universe, indicating the pres-ence of recent star formation. We find these sys-tems to be forming stars at a maximum rate of9-160 M (cid:12) yr − based on the H α luminosity. • For nine objects, we are able to measure the massof the SMBH, the stellar velocity dispersion usingthe narrow component of H α emission line as aproxy, and galaxy mass. We compare these nineobjects to the local scaling relation between themass of the SMBH and the mass or velocity dis-persion of the galaxy. Our systems are both off-set from the M • − σ and M • − M ∗ relationship.Substantial growth is still necessary if these sys-tems are to evolve into the present-day massiveelliptical galaxies. Gas accretion from the CGMand gas-rich minor mergers are necessary to growthe stellar mass and increase the metallicity of theISM. On average, the galaxies need to grow by atleast an order of magnitude in stellar mass if theyare to assemble onto the local scaling relations. A near-constant mass growth rate of ∼
100 M (cid:12) yr − is necessary within a radius of 10 kpc from thequasar from z ∼ • Combining the results of this paper with (Vayneret al. 2021) we find evidence for the onset of feed-back before the galaxies are on the local M • − σ relationship. Luminous type-1 quasars are notthe end phase of massive galaxy formation.The authors wish to thanks Jim Lyke, Randy Camp-bell, and other SAs with their assistance at the tele-scope to acquire the Keck OSIRIS data sets. We wantto thank the anonymous referee for their constructivecomments that helped improve the manuscript. Thedata presented herein were obtained at the W.M. KeckObservatory, which is operated as a scientific partner-ship among the California Institute of Technology, theUniversity of California and the National Aeronauticsand Space Administration. The Observatory was madepossible by the generous financial support of the W.M.Keck Foundation. The authors wish to recognize andacknowledge the very significant cultural role and rev-erence that the summit of Maunakea has always hadwithin the indigenous Hawaiian community. We aremost fortunate to have the opportunity to conduct ob-servations from this mountain. This research has madeuse of the NASA/IPAC Extragalactic Database (NED)which is operated by the Jet Propulsion Laboratory, Cal-ifornia Institute of Technology, under contract with theNational Aeronautics and Space Administration. Software:
OSIRIS DRP (Larkin et al. 2013), Mat-plotlib (Hunter 2007), SciPy (Virtanen et al. 2020),NumPy (Harris et al. 2020), Astropy (Astropy Collab-oration et al. 2018), MAPPINGS (Alarie & Morisset2019), emcee (Foreman-Mackey et al. 2013)APPENDIX A.
3C 93C9 is a luminous quasar at z = 2.019922 with a prominent blue rest-frame UV continuum. For this source,we identify three distinct regions. “SE-SW component A” is a region with a ring-like morphology associated withthe 3C9 quasar host galaxy. We measure a velocity dispersion from a Gaussian fit to the nebular emission lines of407.6 ± − and the kinematics resembling a rotating disk. “SE component A” is classified as an outflow regionwith a very high emission line FWHM of 1362.7 ± − and elevated log([O III ]/H β ) and log([N II ]/H α ) ratiosrelative to the rest of the system. “N component B” is the merging galaxy in the 3C9 system showcasing a line FWHMof 472.15 ± − and a velocity offset of ∼
200 km s − from the quasar. The projected spatial separation betweenthe two apparent nuclei is 9 kpc. The quasar lies in the galaxy with a ring-like morphology showing the kinematicstructure of a disk. Archival HST imaging of rest-frame UV continuum shows the ring morphology as well (see Figure inematics and Energetics of Ionization Gas −2000 −1500 −1000 −500 0 500 1000 1500 2000−1.5−1.0−0.50.00.51.01.52.0 ∆ λ F λ [ e r g / s / c m ] SE component A outflow A [OIII]
Total[OIII] 500.7nm −2000 −1500 −1000 −500 0 500 1000 1500 2000−3−2−10123456
SE component A outflow Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm NE +2000 +1500 +1000 +500 0 500 1000 1500 2000+0.50.00.51.01.52.0
SE-SW component A [OIII]
Total[OIII] 500.7nm −2000 −1500 −1000 −500 0 500 1000 1500 2000−2−101234
SE-SW component A Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm −2000 −1500 −1000 −500 0 500 1000 1500 2000 km/s −0.50.00.51.01.5
N component B [OIII]
Total[OIII] 500.7nm −2000 −1500 −1000 −500 0 500 1000 1500 2000 m/s −2−1012345
N component B Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm
Figure 11.
Spectra of distinct regions along with fits to individual emission lines for the 3C 9 system. −2000−1500−1000−500 0 500 1000 1500 2000−2−10123456 ∆ λ F λ [ e r g / s / c m ] SW component A outf ow [OIII] −2000−1500−1000−500 0 500 1000 1500 2000−1.5−1.0−0.50.00.51.01.5
SW component A outf ow Hα+[NII]
Tota Hα[NII] 658.4 nm[NII] 654.8 nm NE −2000−1500−1000−500 0 500 1000 1500 2000 km/s −0.50.00.51.01.52.02.5
SW component B [OIII]
Tota [OIII] 500.7nm −2000−1500−1000−500 0 500 1000 1500 2000 km/s −2−10123456
SW component B Hα+[NII]
Tota Hα[NII] 658.4 nm[NII] 654.8 nm
Figure 12.
Spectra of distinct regions along with fits to individual emission lines for the 3C268.4 system. Vayner et al. −2000−1500−1000 −500 0 500 1000 1500 2000−3−2−101234 ∆ λ F λ [ e r g / s / c m ] c mp nent A [OIII] T tal[OIII] 500.7nm −2000−1500−1000 −500 0 500 1000 1500 2000−0.50.00.51.01.52.0 c mp nent A Hα+[NII]
T talHα[NII] 658.4 nm[NII] 654.8 nm NE −2000−1500−1000 −500 0 500 1000 1500 2000−2.0−1.5−1.0−0.50.00.51.01.52.02.5 c mp nent A [OIII]
T tal[OIII] 500.7nm −2000−1500−1000 −500 0 500 1000 1500 2000−0.50.00.51.01.52.0 c mp nent A Hα+[NII]
T talHα[NII] 658.4 nm[NII] 654.8 nm −2000−1500−1000 −500 0 500 1000 1500 2000 km/s −0.50.00.51.01.52.0
E c mp nent B [OIII]
T tal[OIII] 500.7nm −2000−1500−1000 −500 0 500 1000 1500 2000 km/s −0.50.00.51.01.52.0
E c mp nent B Hα+[NII]
T talHα[NII] 658.4 nm[NII] 654.8 nm
Figure 13.
Spectra of distinct regions along with fits to individual emission lines for the 7C 1354+2552 system. −5000 −4000 −3000 32000 31000 0 1000 2000 30000123 λ F λ [ × − e − g / s / c m ] W componen/ A o0/flo2 [OIII]
Componen/ A [OIII] 500.7 nmComponent A [OIII] 495.9 nmComponent B [OIII] 500.7 nmComponent B [OIII] 495.9 nmComponent C [OIII] 500.7 nm −1000 0 1000 2000 30003202468 λ F λ [ × − e − g / s / c m ] W componen/ A o0/flo2 Hα+[NII]
TotalComponent A HαComponent A [NII] 658 nmComponent B HαComponent B [NII] 658 nm −5000 −4000 −3000 32000 31000 0 1000 2000 30000.00.5 1e316
E component A o0tflo2 [OIII]
Tot l[OIII] 495.9nm[OIII] 500.7nm
E component A o0tflo2 Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm −5000 −4000 −3000 32000 31000 0 1000 2000 30000.00.5 λ F λ [ × − e − g / . / c m ] SE componen/ B o0/flo2 [OIII] [OIII] 500.7 AGN O0/flo2 [OIII] 495.9 AGN O0/flo2 [OIII] 500.7 AGN ENLR[OIII] 500.7 AGN ENLRHβ AGN ENLR
SE componen/ B o0/flo2 Hα+[NII]
TotalHα AGN ENLR[NII] 658.4 AGN ENLR[NII] 654.8 AGN ENLRHα AGN Driven Outflow −5000−4000−30003200031000 0 1000 2000 3000 km/s
SE component B Tid l fe t0−e [OIII] [OIII] 500.7 nm[OIII] 495.9 nm −1000 0 1000 2000 3000 km/s
SE componen/ B Tid l fe /0−e Hα+[NII]
Hα[NII] 658.4 nm
Figure 14.
Spectra of distinct regions along with fits to individual emission lines for the 3C 298 system. inematics and Energetics of Ionization Gas −2000−1500−1000−500 0 500 1000 1500 2000−2−10123 ∆ λ F λ [ e r g / s / c m ] NW component A t dal feature A [OIII]
Total[OIII] 500.7nm −2000−1500−1000−500 0 500 1000 1500 2000−3−2−1012345
NW component A t dal feature Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm NE +2000+1500+1000+500 0 500 1000 1500 2000 km/s +1.5+1.0+0.50.00.51.01.5 ∆ λ F λ [ e r g / s / c m ] E-W component B [OIII]
Total[OIII] 500.7nm −2000−1500−1000−500 0 500 1000 1500 2000 km/s −2−101234
E-W component B Hα+[NII]
TotalHα[NII] 658.4 nm[NII] 654.8 nm NE
Figure 15.
Spectra of distinct regions along with fits to individual emission lines for the 3C 446 system. −2000−1500−1000−500 0 500 1000 1500 2000−1.0−0.50.00.51.01.52.0 ∆ λ F λ [ e r g / s / c ] NE co ponent A [OIII] −2000−1500−1000−500 0 500 1000 1500 2000 k /s −6−4−20246
NE co ponent A Hα+[NII]
TotalHα[NII] 658.4 n [NII] 654.8 n NE −2000−1500−1000−500 0 500 1000 1500 2000 k /s −1.0−0.50.00.51.0
N co ponent A/B(?) [OIII]
Total[OIII] 500.7n
Figure 16.
Spectra of distinct regions along with fits to individual emission lines for the 4C 57.29 system.B.
4C 09.174C 09.17 is a luminous quasar at z=2.117 with a blue UV continuum. For this source, we identify four distinct regions.“SW component A” is a star-forming clump associated with the quasar host galaxy. The spectrum of this region showsa single narrow emission line with an FWHM of 312.0 ± − . “S/E component A” is an outflow region drivenby the quasar, the nebular emission lines for this region have an FWHM of 887.2 ± − . A second narrowcomponent is required for a good fit for each emission line in this region, with a line FWHM of 290.4 ± − . “W component B clumps” is a region part of the merging galaxy within the 4C09.17 system. The region consistsof clumpy emission selected by isolating spaxels with an H α line surface density > × − erg s − cm − arcsec − .“W component B diffuse” is emission associated with “diffuse” ionized emission in the merging galaxy selected by6 Vayner et al. −2000−1500−1000−500 0 500 1000 1500 2000 km/s −0.50.00.51.01.52.02.53.03.54.0 ∆ λ F λ [ e r g / s / c m ] N,S compo e t A [OIII]
Total[OIII] 500.7 m −2000−1500−1000−500 0 500 1000 1500 2000 km/s −0.50.00.51.01.52.02.53.0
N,S compo e t A Hα+[NII]
TotalHα[NII] 658.4 m[NII] 654.8 m NE
Figure 17.
Spectra of distinct regions along with fits to individual emission lines for the 4C 22.44 system. −2000−1500−1000−500 0 500 1000 1500 2000−0.50.00.51.01.5 ∆ λ F λ [ e r g / s / c ] S co ponent A Outflow [OIII]
Total[OIII] 500.7n −2000−1500−1000−500 0 500 1000 1500 2000−20246
S co ponent A Outflow Hα+[NII]
TotalHα[NII] 658.4 n [NII] 654.8 n NE −2000−1500−1000−500 0 500 1000 1500 2000−1.0−0.50.00.51.01.5
NE co ponent A Outflow [OIII]
Total[OIII] 500.7n −2000−1500−1000−500 0 500 1000 1500 2000−0.4−0.20.00.20.40.60.8
NE co ponent A Outflow Hα+[NII]
TotalHα[NII] 658.4 n [NII] 654.8 n −2000−1500−1000−500 0 500 1000 1500 2000 k /s −0.4−0.20.00.20.40.60.81.0
SW co ponent A clu p [OIII]
Total[OIII] 500.7n −2000−1500−1000−500 0 500 1000 1500 2000 k /s −2−1012345
SW co ponent A clu p Hα+[NII]
TotalHα[NII] 658.4 n [NII] 654.8 n
Figure 18.
Spectra of distinct regions along with fits to individual emission lines for the 4C 05.84 system. isolating spaxels with an H α spatial line surface density < × − erg s − cm − arcsec − . The diffuse region showshigher log([O III ]/H β ) and log([N II ]/H α ) line ratios associated with both AGN and star formation photoionizationwhile the clumpy regions of the merging galaxy showcase lower ionization levels consistent with photoionization bystar formation. This region is associated with bright UV emission in HST imaging of this object (Lehnert et al.1999). “S/E component A outflow” shows high log([N II ]/H α ) and log([O III ]/H β ) values relative to the rest of thesystem, indicating this region is predominantly photoionized by the quasar. The 4C09.17 system is a merger of twogalaxies with velocity offsets of ∼ − and a projected separation of ∼ C.
3C 268.43C 268.4 is a luminous quasar at z=1.39, with a slightly reddened UV continuum compared to the average type-1quasar. For this target, we identified two distinct regions. “SW component A” is an outflow driven by the quasar. TheFWHM of the emission lines is 2075 ±
354 km s − as measured from the Gaussian fit to the [O III ] line. The spectrum inematics and Energetics of Ionization Gas ± − , most likely signalingemission from an extended narrow-line region close to the quasar. Because of issues with miss-assignment of flux inthe OSIRIS pipeline (Lockhart et al. 2019), the rows below and above the centroid of the quasar do not have properlyextracted spectra in the H band observations of this object. Hence we do not have a good spectrum of the extendedemission in a 0.2-0.3 (cid:48)(cid:48) radius around the quasar in the H band, which covers the H α and [N II ] emission lines of theionized outflow. “SW component B” is a region associated with the merging galaxy, showcasing clumpy morphology inionized gas emission. The emission lines have an FWHM of 367.7 ± − and an offset of −
300 km s − relativeto the redshift of the quasar. The log([O III ]/H β ) line ratios are lower for this region compared to the rest of thesystem, consistent with a mixture of AGN and star formation photoionization. This region is also associated withbright rest-frame UV continuum emission, seen with HST observations of this target Hilbert et al. (2016). D.
7C 1354+25527C 1354+2552 is a luminous quasar at z=2.0064 with a blue UV continuum. For this target, we identify two distinctregions. “Component A” is the extended emission associated with the quasar host galaxy. The kinematics show asmooth velocity gradient, indicating the presence of a galactic disk. The size, morphology, and kinematics of the diskare similar to that of star-forming galaxies on the more massive end of the star formation main sequence at z ∼ ± − on the redshifted sideof the disk and 497.7 ± − on the blue-shifted side of the disk. Although this region only has a single label(“component A”), in Figure 13 rows one and two show the fits to the red and blue-shifted sides of the disk that arepart of this region. This region is selected based on the location where H α emission is detected. This is done to boostthe SNR in the H α line as it appears to be clumpier, more compact, and less extended than [O III ] . In Table 3 weprovide values integrated over the entire galactic disk. “E component B ” is a region associated with the merginggalaxy at a projected separation of 6-7 kpc. The kinematics are consistent with a dispersion dominated galaxy. Theentire “component A” is consistent with quasar photoionization. The gas in “E component B” is photoionized by starformation. E.
3C 2983C298 is a luminous quasar at z=1.439 with a slightly reddened UV continuum. For this target, we identify fivedistinct regions. We present a detailed discussion of each region in Vayner et al. (2017). “W/E component A” areoutflow regions with a bi-conical morphology, where the western (W) is the redshifted receding cone, and the eastern(E) is the approaching cone. In Vayner et al. (2017), they are referred to as the red(blue) shifted outflow region.The emission lines over the outflows are very broad, with FWHM up to ∼ − . A combination of shocksand quasar activity is likely responsible for photoionizing the gas. “SE component B outflow” is an outflow regionbelonging to a merging galaxy. “SE component B ENLR” is an extended narrow-line region belonging to the disk ofthe merging galaxy, with gas photoionized by the quasar or secondary AGN. “SE component B Tidal feature” is aregion of the merging galaxy with active/recent star formation as evident by lower log([N II ]/H α ) and log([O III ]/H β ) values compared to the rest of the regions. F. ± (cid:12) yr − . This star formation rate is farlower than the far infrared derived rate of 580M (cid:12) yr − . The extinction towards the nuclear region measured fromWillott et al. (2000) alone cannot explain the mismatch between the H α and far-infrared derived SFR. Either a largerfraction of the far-infrared emission is from dust that is being heated by the AGN, or the far-infrared emission tracesa different star formation history than H α (Calzetti 2013). No narrow extended emission is detected in this object.The merger status of this object is unclear. Two nearby galaxies to the north and west of the quasar are visiblein archival HST imaging (Willott et al. 2000). We do not detect the western object that is 2 (cid:48)(cid:48) away from the quasarin our OSIRIS observations in any emission line. Willott et al. (2007) studied this object with PdBI through CO2-1 emission at a fairly coarse ( ∼ Vayner et al. G.
4C 57.294C 57.29 is a luminous quasar at z=2.1759 with a blue UV continuum. For this target, we identify two regions.Region “NE component A” belongs to the host galaxy of the quasar. The relatively high log([O
III ]/H β ) value indicatesthat this region is consistent with being photoionized by the quasar. The 500.7 nm [O III ] is the only emission linedetected for this region. The region is marginally resolved, making it hard to measure the kinematic structure. Werequire a double Gaussian fit to the [O
III ] emission in this region to obtain a good fit, and we measure an FWHMof 474.3 and 502.5 km s − with offsets of 35.0 km s − and -1050.1 km s − relative to the redshift of the quasar. Weidentify a second region north of the quasar. It is unclear if it belongs to a merging galaxy or the quasar host galaxy.There is a ∼
100 km s − offset from the quasar, and the line has an FWHM of 550.13 ∼
19 km s − . This region is alsoonly detected in [O III ] . The SNR is too low to measure any kinematics structure. H.
4C 22.444C22.44 is a luminous quasar at z=1.5492 with a reddened UV continuum. Similar to 3C318, we do not detect anyevidence for a merging galaxy for this system. For this target, we identify a single region, “N, S component A”. Thekinematics of this region may be consistent with a galactic disk belonging to the quasar host galaxy. We see evidencefor a smooth gradient in the radial velocity map, however, the region is marginally resolved. We measure an emissionline FWHM of 434.8 km s − . The region is consistent with being ionized by star formation with some contributionfrom quasar photoionization. I.
4C 05.844C05.84 is a luminous quasar at z=2.323 with a slightly reddened UV continuum. For this target, we identifythree distinct regions. Regions “S component A” and “NE component A” are the blue(red) shifted outflow regionsresembling a bi-conical outflow. They showcase broad extended emission with a line FWHM of ∼
800 km s − . Thequasar photoionizes these regions. Region “SW component A clump”, shows a line FWHM of 467.9 ± − andis photoionized by a combination of star formation and the quasar. This clump is also detected in NIRC2 imaging ofthis object studied by Krogager et al. (2016), where they consider this clump to be associated with a damped Ly α system. However, here we confirm that this objected is part of the quasar host galaxy. We find no evidence for amerging galaxy within our OSIRIS observations. J.
3C 4463C446 is a quasar at z = 1.404. For this target, we identify two regions, “N component A tidal feature” is a regionbelonging to the quasar host galaxy, resembling a tidal feature that is most likely induced by the merger. We measurean FWHM of 395.14 ± − for this region. “E-W component B” belongs to the merging galaxy, a portion of itresembles a tidal feature, counter to the tidal arm of “N component A tidal feature.” For this region, we measure aline FWHM of 558.5 ±
63 km s − however, it appears to be a blend of two velocity components. It is unclear where thenucleus of the merging galaxies resides. It could be that it has already merged with that of the quasar host galaxy.The two galaxies appear to be offset by at least 500 km s − in velocity, and there is a possibility that a portion of themerging galaxy lies on top of the quasar host galaxy. K.
4C 04.814C04.81 is a luminous quasar at z=2.5883 with a reddened UV continuum. For this target, we identify a singleregion, “E component A outflow”. The kinematics show blue and red-shifted broad (FWHM ∼
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