# The Diverse Morphology, Stellar Population, and Black Hole Scaling Relations
of the Host Galaxies of Nearby Quasars

Yulin Zhao Kavli Institute for Astronomy and Astrophysics, Peking University,
Beijing 100871, China  Department of Astronomy, School of Physics, Peking
University, Beijing 100871, China  Luis C. Ho Kavli Institute for Astronomy
and Astrophysics, Peking University, Beijing 100871, China  Department of
Astronomy, School of Physics, Peking University, Beijing 100871, China  Jinyi
Shangguan Max-Planck-Institut für Extraterrestrische Physik (MPE),
Giessenbachstr., D-85748 Garching, Germany  Kavli Institute for Astronomy and
Astrophysics, Peking University, Beijing 100871, China  Minjin Kim Department
of Astronomy and Atmospheric Sciences, Kyungpook National University, Daegu
702-701, Republic of Korea  Korea Astronomy and Space Science Institute,
Daejeon 305-348, Republic of Korea  Dongyao Zhao Beijing Planetarium, Beijing
Academy of Science and Technology, Beijing, 100044, China  Kavli Institute for
Astronomy and Astrophysics, Peking University, Beijing 100871, China  Hua Gao
Kavli Institute for Astronomy and Astrophysics, Peking University, Beijing
100871, China  Department of Astronomy, School of Physics, Peking University,
Beijing 100871, China  Kavli Institute for the Physics and Mathematics of the
Universe (WPI),The University of Tokyo Institutes for Advanced Study,
The University of Tokyo, Kashiwa, Chiba 277-8583, Japan

###### Abstract

We present rest-frame $B$ and $I$ imaging of 35 low-redshift ($z<0.5$)
Palomar-Green quasars using the Hubble Space Telescope Wide Field Camera 3. We
perform multi-component two-dimensional image decomposition to separate the
host galaxy from its bright active nucleus, characterize its morphology, and
measure its photometric properties. Special care is devoted to quantifying the
structural parameters of the galaxy bulge, determine its $B-I$ color, and
estimate its stellar mass. Roughly half of the sample, comprising the less
luminous ($\mbox{$L_{5100}$}\lesssim 10^{45}\,\mathrm{erg\,s^{-1}}$) but most
high Eddington ratio quasars, reside in disk galaxies that are often barred
and possess pseudo bulges. The large stellar masses, large effective radii,
and faint surface brightnesses suggest that the host galaxies of the most
luminous quasars are mostly ellipticals. Major mergers constitute only a
minority ($\lesssim 20\%$) of our sample. Our quasar sample roughly obeys the
scaling relations between black hole mass and host galaxy (bulge, core, total)
stellar mass. Hosts with black holes more massive than $\sim
10^{8}\,M_{\odot}$ behave similarly to classical bulges and early-type
galaxies, while those with less massive black holes, particular the narrow-
line Seyfert 1s, are consistent with pseudo bulges in late-type galaxies. The
host galaxy bulges, irrespective of whether they are classical or pseudo,
follow the relatively tight inverse relation between effective radius and mean
effective surface brightness of inactive classical bulges and ellipticals. We
argue that pseudo bulges experience recent or ongoing nuclear star formation.

galaxies: evolution — galaxies: formation — galaxies: active — galaxies:
bulges — galaxies: photometry — quasars: general

††journal: ApJ

##  1 Introduction

Supermassive black holes (BHs) are commonplace in the centers of massive
galaxies, and their masses strongly correlate with the properties of the
bulges of their host galaxies, including luminosity (Kormendy & Richstone,
1995), stellar mass ($M_{\rm bulge}$; Magorrian et al. 1998), and stellar
velocity dispersion ($\sigma_{\star}$; Ferrarese & Merritt 2000; Gebhardt et
al. 2000). These scaling relations split into two branches (Kormendy & Ho,
2013): classical bulges and elliptical galaxies form a tight relation, while
the relation for pseudo bulges exhibits markedly larger scatter and a
different zero point. The distinctly different trends observed for the two
bulge types in terms of their BH-host correlations likely reflect their
different evolutionary pathways (Kormendy & Kennicutt, 2004). The exact origin
of these BH-host scaling relations is still a subject of lively research
(Menci et al., 2016; Thomas et al., 2019; Li et al., 2020; Terrazas et al.,
2020). From an observational point of view, a number of pressing issues remain
to be clarified. Which came first, BH or galaxy? Must the growth of the BH and
its host be finely synchronized so as to preserve the small intrinsic scatter
observed in the local scaling relations? What is the cosmic evolution of the
zero point and scatter of the scaling relations? Do they depend on the
morphological type of the host or environment?

An important step toward addressing these issues can be made by extending the
BH-host scaling relations to active galaxies, wherein the BHs are still
growing. Toward this end, two ingredients are needed: BH masses and host
galaxy parameters. Fortunately, BH masses can now be estimated with reasonable
($\sim 0.3-0.4$ dex) accuracy from single-epoch spectra for broad-line (type
1) active galactic nuclei (AGNs), particularly for $z\lesssim 0.75$ when the
more reliable rest-frame optical lines can be accessible readily (e.g.,
Vestergaard & Peterson 2006; Ho & Kim 2015). By contrast, host galaxy
parameters, especially for the bulge, are more challenging to obtain because
of the bright glare of the AGN. With considerable effort, stellar kinematics
have been measured successfully for type 1 AGNs (e.g., Greene & Ho 2006;
Bennert et al. 2015; Shen et al. 2015; Caglar et al. 2020), but the samples
are small, and $\sigma_{\star}$ becomes increasingly difficult, if not
impossible, to measure for the most powerful quasars, especially at higher
redshift. The bulge stellar mass offers a promising alternative replacement of
$\sigma_{\star}$ for investigating BH-host scaling relations. At least insofar
as classical bulges and ellipticals are concerned, $M_{\rm bulge}$ correlates
just as tightly with $M_{\rm BH}$ as does $\sigma_{\star}$ (Kormendy & Ho,
2013). Unlike stellar kinematics, however, $M_{\rm bulge}$ is more
observationally tractable. Securing a robust measurement of $M_{\rm bulge}$
for AGNs requires (1) high-resolution imaging with a stable point-spread
function (PSF), (2) an effective analysis technique to deblend the host galaxy
from the active nucleus and to decompose the bulge itself from other possible
substructure in the host, and (3) sufficient color information to constrain
the mass-to-light ratio ($M/L$). Detailed two-dimensional (2D) image
decomposition enables quantitative analysis of the structural components of
quasar host galaxies (e.g., Kim et al. 2008b; Veilleux et al. 2009; Schramm &
Silverman 2013; Kim et al. 2017; Bentz & Manne-Nicholas 2018; Li et al. 2021).

In this current work, we present Hubble Space Telescope (HST) observations of
a small, carefully selected sample of nearby quasars designed to yield
reliable photometry of the host galaxy, with the intent of estimating stellar
masses. Special care is devoted to deriving accurate photometric and
structural parameters for the bulge, for which it is of paramount importance
to properly model strong bars, spiral arms, and other nonaxisymmetric
structures. Unlike earlier works (Kim et al., 2008b; Kim et al., 2017), here
we provide rest-frame optical colors in order to mitigate possible variations
in $M/L$, which can be considerable in light of the evidence for recent or
ongoing star formation (e.g., Husemann et al. 2017; Kim & Ho 2019; Zhao et al.
2019; Jarvis et al. 2020; Shangguan et al. 2020b; Zhuang & Ho 2020; Xie et al.
2021). We focus on low-redshift quasars in order to best characterize the
local scaling relations for AGNs, laying the foundation for future
investigations of possible cosmic evolution. It is imperative to establish the
true intrinsic scatter of the scaling relations, as well as to ascertain
whether and how they might depend on the properties of the AGNs and the host
galaxies. For example, if substantial BH growth occurs during the AGN phase
without synchronized star formation in the bulge, we expect the zero point or
scatter of the scaling relations to depend on the mass accretion rate ($L_{\rm
bol}/L_{\rm Edd}$) and/or color of the bulge. Or, if galaxy major mergers play
a role in establishing the relations, their scatter might correlate with the
presence of tidal features or local environment. The dependence on bulge type
and galaxy morphology has been well-established for inactive galaxies (e.g.,
Kormendy & Ho 2013; Greene et al. 2020, and references therein), and a similar
effort should be extended to their active counterparts. While the sample
considered here is far too small to definitively address all of these issues,
it serves as a useful starting point.

Our HST observations and data reduction are described in Section 2. We build a
library of empirical PSFs performed bulge-disk decomposition in Section 3.
Section 4 analyses the main results on morphological classification, stellar
mass estimation, and calculation of BH masses. Implications are discussed in
Section 5, and the main conclusions are summarized in Section 6. This work
adopts the following parameters for a $\Lambda$CDM cosmology:
$\Omega_{m}=0.308$, $\Omega_{\Lambda}=0.692$, and
$H_{0}=67.8\,\mathrm{km\,s^{-1}\,Mpc^{-1}}$ (Planck Collaboration et al.,
2016).

Figure 1: The distribution of $V$-band absolute magnitude versus redshift for
the 87 low-redshift ($z<0.5$) PG quasars. The filled red circles are the 35
targets observed by our snapshot survey.

##  2 Data

###  2.1 Sample and Observations

Our parent sample consists of the 87 low-redshift ($z<0.5$) optically/UV-
selected quasars from the Palomar-Green (PG) survey (Schmidt & Green, 1983;
Boroson & Green, 1992). Although the PG survey is incomplete due to large
photometric error (Goldschmidt et al., 1992), this representative sample of
nearby, bright quasars has been extensively studied for decades. The
availability of line widths for the broad H$\beta$ emission line from single-
epoch spectroscopy (Boroson & Green, 1992; Ho & Kim, 2009) and
spectrophotometry for the optical continuum (Neugebauer et al., 1987) enables
BH masses to be estimated from conventional empirical calibrations (e.g.,
Vestergaard & Peterson 2006; Ho & Kim 2015). Robust bolometric luminosities,
and hence Eddington ratios, can be derived from comprehensive broadband
spectral energy distributions (Ho, 2008). Figure 1 shows the distribution of
$V$-band absolute magnitude versus redshift for the sample.

We undertook a snapshot program (proposal 12903; PI: L. C. Ho) using the HST
Wide Field Camera 3 (WFC3; Dressel 2019) to obtain high-resolution rest-frame
optical images for a significant subset of the PG quasars. Snapshot programs
cannot guarantee that any specific target will be observed, instead drawing
from a pre-designated list of candidates to fill in gaps in the HST schedule
to obtain relatively short observations of typical duration 45 minutes or less
(including all overheads). Our program achieved a success rate of $\sim 40\%$,
securing useful data for 35 sources randomly selected from the original sample
of 87 quasars. The final sample forms a representative subset of the original
sample (Figure 1). Details of the observed targets and the observations are
summarized in Table 1.

Each target was observed with two filters, which roughly correspond to rest-
frame $B$ band and $I$ band, with the intent of providing a sufficiently wide
color leverage to constrain the $M/L$ of the stellar continuum in order to
estimate the stellar mass of the host galaxy. Each pair of filters was
customized carefully according to the redshift of the source to avoid, to the
extent possible, contamination from strong emission lines. In general each
object was observed with both the UVIS and IR channels of WFC3, except for the
nine sources with $z\lesssim 0.05$, for which both filters fall in the UVIS
channel. The UVIS detector has a pixel scale of $0\farcs 0400\times 0\farcs
0398$ and a point-spread function (PSF) with a full-width at half maximum
(FWHM) of $\sim 0\farcs 08$. The IR channel has a pixel scale of $0\farcs
135\times 0\farcs 121$ and a PSF with ${\rm FWHM}\approx 0\farcs 2$. To reduce
the readout overheads, we employ the UVIS2-M1K1C-SUB subarray for the UVIS
detector and the IRSUB512 subarray for the IR detector, which yield,
respectively, a field-of-view of $40\arcsec\times 40\arcsec$ and
$69\arcsec\times 62\arcsec$, adequate to cover the target and sufficient
nearby field for background subtraction. Exposure times varied. Every target
was observed with a long exposure with a total integration time of $\sim
200-500$ s in rest-frame $B$ and $\sim 150-300$ s in rest-frame $I$. We use a
three-point and four-point dither pattern for the UVIS and IR observations,
respectively, to facilitate rejection of cosmic rays and bad pixels, as well
as to better sample the PSF. Given the brightness of the quasar at short
wavelengths, saturation in the center is unavoidable in the long UVIS
exposure. We acquired an additional short ($\sim$ 6–100 s) UVIS exposure for
the purposes of replacing the saturated nuclear region of the long exposure.
The IR images, taken in MULTIACCUM mode with STEP25 timing sequence
($\mathrm{NSAMP}=6$), are not saturated.

Most of the observations were executed in the standard FINE LOCK mode, but
four targets (PG 0050+124, 1202+281, 1226+023, and 1244+026) were observed in
GYRO mode after a failure in guide star acquisition. This resulted in
significant degradation of the pointing precision during the observations,
prompting us to apply corrective procedures described in Appendix A.

###  2.2 Data Reduction

We reprocessed the calibrated, flat-fielded science files using the auxiliary
data sets acquired from the Hubble Data Archive111https://hla.stsci.edu/ on 25
January 2016. We used AstroDrizzle (Gonzaga et al., 2012), largely with
default parameters, to process the single-frame short exposure and to combine
the dithered frames of the long exposure to remove cosmic rays and correct for
geometric distortion. The dithered sub-exposures are weighted by their maps of
effective exposure time during the drizzle combination. The pixel scale of the
combined image is set to $0\farcs 06$ for the IR channel and $0\farcs 03$ for
the UVIS channel to achieve Nyquist sampling. The final images have a size of
$54\arcsec\times 54\arcsec$ for the IR channel and $36\arcsec\times 36\arcsec$
for the UVIS channel. The saturated nucleus of the long UVIS exposure was
replaced by the corresponding pixels from the unsaturated short exposure,
which was aligned to the former using AstroDrizzle. Depending on the level of
saturation of the cores, we usually replace a circular region with radius
$0\farcs 3$.

###  2.3 Background Estimation

Although AstroDrizzle performs background removal, we take additional steps to
verify the accuracy of the background determination and quantify its
uncertainty. While the pixel-to-pixel variation of the background emission is
dominated by Poisson noise, often the main uncertainty of the background comes
from larger scale fluctuations. We first use the deeper $I$-band image to
create a source mask from the segmentation image produced by SExtractor
(Bertin & Arnouts, 1996). The $I$-band mask is then applied to the $B$-band
image. As in Li et al. (2011), the mask is expanded according to the flux of
the source. Special care is paid to mask the region of the quasar and its host
galaxy, whose total radial extent is determined from inspection of its one-
dimensional (1D), azimuthally averaged surface brightness profile. If the
quasar is sufficiently strong, the mask for its diffraction spikes needs to be
adjusted manually. In order to identify the spatial scale and amplitude of the
background variation, we randomly sample the unmasked pixels $10^{4}$ times
with boxes of increasing size. For each box size, we calculate the RMS of the
sampled median background values. The scale of the background fluctuation is
approximated by the box size above which the RMS converges, and the
uncertainty of the background is given by the RMS value (Li et al., 2011;
Huang et al., 2013). A constant background, the median of the sampled values,
is subtracted from the image.

Figure 2: Using a single star to mimic the WFC3 $I$-band observations of a
quasar without a host galaxy, we fit the mock quasar with a PSF generated from
(top) TinyTim, (middle) a single star, and (bottom) stacking a large number of
stars. From left to right, the panels show the 1D profile, the observed image
of the mock quasar, the best-fit GALFIT model, and the residuals. The 1D
azimuthally averaged profile shows the original data (black circles) and the
best-fit model (magenta solid line). The reduced $\chi^{2}$ from GALFIT as
well as the RMS of the 1D surface brightness residuals at radius below and
above 0$\farcs$6 are shown in the lower-left corner. All images are on an
asinh stretch.

##  3 Image Decomposition

The bright nucleus of type 1 AGNs poses a severe challenge for quantitative
measurements of the stellar light of the underlying host galaxy, especially so
for the bulge component, which is affected most acutely by contamination from
the central point source (e.g., McLure & Dunlop, 2002; Peng et al., 2006;
Greene et al., 2008; Kim et al., 2008a, b; Jiang et al., 2011; Kim et al.,
2017). We use the 2D fitting code GALFIT (version 3.0; Peng et al., 2002,
2010) to simultaneously model the active nucleus and the host galaxy with
parametric models. GALFIT, a nonlinear fitting algorithm that uses Levenberg-
Marquardt least-square minimization (Press et al., 1992), has been widely used
to study the morphologies of galaxies, by decomposing their substructures
using various PSF-convolved parametric functions to account for the bulge,
disk, bar, and even spiral arms and tidal features (Peng et al., 2010). A
point-like nucleus can be added when an AGN is present.

We model the bulge, disk, and bar with the Sérsic (1968) function

$\Sigma(r)=\Sigma_{e}\,\exp\left[-\kappa\left(\frac{r}{R_{e}}\right)^{1/n}-1\right],$
(1)

where $\Sigma_{e}$ is the surface brightness at the effective radius $R_{e}$,
$n$ is the Sérsic index, which controls the concentration of the profile, and
$\kappa$ is related to $n$ by the incomplete-gamma function
$\Gamma\left(2n\right)=2\gamma\left(2n,\kappa\right)$ (Graham & Driver, 2005).
When $n$ is large, the Sérsic function has a steep inner profile and an
extended outer wing; when $n$ is small, the Sérsic function has a shallow
inner profile and a steep truncation at large radius. When fit as single-
component entities, elliptical galaxies have $n\approx 3-8$ (e.g., Huang et
al. 2013), classical bulges exhibit $n\approx 1.5-4$ (e.g., Gao et al. 2018,
2020), and pseudo bulges have $n\approx 1-3$ (Gao et al., 2020). Disks are
well-described by an exponential ($n\approx 1$; Freeman 1970), at least over
the radial extent not modified strongly by inner or outer breaks (e.g., Pohlen
& Trujillo, 2006; Li et al., 2011; Gao & Ho, 2017; Gao et al., 2019), while
bars usually can be fit with an $n=0.5$ profile (Freeman, 1966; de Jong,
1996). GALFIT employs coordinate rotation to produce realistic-looking spiral
arms, and Fourier modes can mimic lopsidedness and a variety of non-
axisymmetric tidal features (Peng et al., 2010).

###  3.1 Point-spread Function

Accurate knowledge of the PSF is critical to decomposing the bright nucleus of
an active galaxy, a prerequisite to deriving reliable measurements of the
host. While the PSF of HST is much more stable than that of ground-based
telescopes, it still varies with time due to telescope breathing and changes
in instrument focus, and with position across the field due to optical
distortion and the properties of the detector. Ideally, an empirical PSF model
can be obtained by observing a bright star with a spectrum similar to that of
the science target, observed close in time to and using the same observing
strategy as the science target (Biretta, 2012; Mechtley, 2014). This option
was not possible for our observations, in view of the severe scheduling
limitations of the snapshot program. Moreover, our decision to use subarrays
to minimize readout time further significantly reduced the field-of-view of
the observations, to such an extent that hardly any of the images contains
useful foreground stars that can serve as a PSF model. In the absence of a
dedicatedly observed PSF, many studies have utilized synthetic PSFs (TinyTim;
Krist et al. 2011) to analyze HST images of AGN host galaxies (e.g., Greene et
al. 2008; Kim et al. 2008a; Jiang et al. 2011; Kim et al. 2017). Here, we
explore another strategy, by constructing an empirical PSF using stars
extracted from archival images observed non-simultaneously in other HST
programs, but otherwise employing the same detector-filter combination and
dithering strategy. We stack a number of stars to boost the signal-to-noise
ratio. For each detector-filter combination used in our program, we searched
the HST archive for useful images that contain suitably isolated, unsaturated
stars that are at least one-third as bright as the corresponding quasar. We do
not include stars with close contamination (extended sources or bright point
sources) within $\sim 3\arcsec$, where the diffraction spike is prominent. We
use AstroDrizzle to process the images of the stars with exactly the same
procedure and pixel scale as the science images. Each image, typically $\sim
30\arcsec$ on a side, is inspected carefully for possible faint contaminating
sources, which, when present, are replaced manually by random Gaussian noise
using the standard deviation measured from the nearby background. The final
PSF for each filter, a median-stacked combination of typically $\sim 10-65$
stars, is only slightly ($\lesssim 10\%$) broader than the profile of
individual stars.

To test the effectiveness of our strategy, we choose a single, bright star
having roughly similar brightness as a typical quasar in our sample to serve
as the mock image of a “pure” quasar nucleus (one lacking any underlying host
galaxy), and we fit it with different PSF models to evaluate how well each can
match the point source. Figure 2 compares the performance of (1) a standard
synthetic PSF generated from TinyTim, (2) an empirical PSF derived from a
single star in the same field-of-view as the mock quasar with similar flux,
and (3) an empirical PSF produced from stacking a large number of stars.222The
star used to represent the mock quasar is not included in the stack. The image
of the mock quasar and the PSF models are convolved with a Gaussian kernel to
achieve Nyquist sampling, as discussed in Section 3.2. Using GALFIT to fit the
mock quasar, we find that the stacked-star PSF performs slightly better than
the single-star PSF, and both empirical PSFs are always superior to the
TinyTim PSF, as judged by the reduced $\chi^{2}$ reported by GALFIT as well as
the RMS of the residuals of the 1D profile. This conclusion holds for tests
with different sets of stars on images of different bands (e.g., F438W, F814W,
F105W, and F110W). We also explored the influence of the number of stars used
in the stacking, finding, as a rule of thumb, that combining $\gtrsim 10$
stars produces the most optimal empirical PSF.

###  3.2 Fitting Procedure

PSF mismatch in the central pixels of bright nuclei presents a critical
challenge for the analysis of the host galaxies of AGNs, especially those
bright enough to be considered quasars. The problem is exacerbated when the
PSF is undersampled, as is the case for HST. While employing dithering
strategies can alleviate the problem, in practice this still proved
insufficient for our purposes, especially in the absence of dedicated PSF
observations (Section 3.1). Under these circumstances, one can mitigate the
PSF mismatch by artificially broadening the PSF image and science image by a
modest amount, by convolving them with a Gaussian kernel so that the PSF is
Nyquist-sampled by the pixel scale of the detector (Kim et al., 2008a). We
adopt the same strategy for this study. GALFIT requires a “$\sigma$ image” to
properly calculate $\chi^{2}$. Following Casertano et al. (2000), the $\sigma$
image takes into account noise correlation due to the drizzling process of
combining images and the different exposure times if the core of the nucleus
is replaced. We nevertheless confirm that the decomposition results, after
accounting for systematic uncertainties, are not significantly different if we
adopt the “$\sigma$ image” generated automatically by GALFIT.

We first focus on the $I$-band data, which have higher signal-to-noise ratio
than the $B$-band data. The $I$ band also offers higher contrast between the
host galaxy and the AGN. Once a final, viable model is constructed for the $I$
band, we adopt the best-fit parameters and apply them to the $B$ band,
adjusting only the flux amplitude of each component. This approach, similar to
that of Li et al. (2021), ensures a consistent bulge-disk decomposition in
both bands and mitigates severe PSF mismatch in the $B$ band, even if the
$B$-band model may not be optimized (see Section 3.3). We use the same source
mask to estimate the background. We start with the simplest possible model—a
point-like nucleus plus a single-component host galaxy—and slowly include more
components (e.g., a disk or a bar) to account for additional substructure of
the host galaxy if the data require. The centers of all of the components are
tied together. Close companions, when sufficiently nearby, need to be fit
simultaneously. Whenever possible, we allow the Sérsic index of the spheroidal
component (elliptical galaxies and the bulges of disk galaxies) to be fit
freely. However, as the solutions are not always meaningful because of
parameter degeneracies or strong systematic effects arising from PSF mismatch,
we often need to fix the Sérsic index to discrete values between $n=1$ to
$n=4$ and choose that which produces the lowest residuals, with the range of
acceptable solutions serving to bracket the uncertainties. Given the many
systematic uncertainties in model selection, it is difficult to rely on formal
statistical criteria (e.g., reduced $\chi^{2}$) to determine the goodness-of-
fit. Instead, we utilize the 1D and 2D residuals to inform us of the relative
merits of different models and to decide on whether and in which manner more
complex models need to be considered. For example, spiral structure betrays
the presence of a disk, and bisymmetric residuals often indicate that a bar is
present. We invoke coordinate rotation to capture obvious spiral arms, even if
they do not significantly affect the photometric parameters of the bulges of
inactive galaxies (Gao & Ho, 2017). Bending modes can mimic tidal features
effectively, and Fourier modes of order up to 3 can accommodate a variety of
non-asymmetric structures in the outermost extended regions (e.g., Kim et al.,
2008b). In particular, the absolute amplitude of the $m=1$ Fourier mode
($A_{L}$) can effectively gauge the global asymmetry (Peng et al., 2010).
Quasar hosts with $A_{L}>0.2$ tend to be highly disturbed systems that are
ongoing or late-stage mergers (Kim et al., 2017).

Figures 3 and 4 show the fits of the $I$-band and $B$-band images,
respectively, and Appendix B provides notes on the individual objects. The
best-fit parameters are reported in Table The Diverse Morphology, Stellar
Population, and Black Hole Scaling Relations of the Host Galaxies of Nearby
Quasars.

###  3.3 Uncertainties

The formal uncertainties reported by GALFIT are usually small. In practice,
the error budget for most of the derived parameters are dominated by several
sources of systematic uncertainty. First, as described in Section 3.2, many of
the fitting results, especially those that pertain to the bulge, are extremely
sensitive to the PSF. While we have made every effort to secure a reliable
empirical PSF based on stacking a number of observed stars (Section 3.1), of
course, the final PSF depends on the exact number and choice of stars. To
account for this factor, we repeat the fits with different realizations of the
empirical PSF constructed from various combinations of observed stars. Next,
due to the limited field-of-view of our images, we opt to fix the background
to its predetermined value (Section 2.3) in order to facilitate the
convergence of the fit (Kim et al., 2008a), but we explore the impact of
large-scale fluctuations of the background by repeating the fits with the
background value perturbed around its RMS value. Finally, as documented in
Appendix B, a major source of uncertainty occasionally arises from the
ambiguity of model selection. Because of the complexity of AGN host galaxies,
the robustness of the fits must be scrutinized individually for each source,
and often the “best fit” is based less on any rigorous statistical criteria
than on a reasonable choice bracketed by a range of plausible choices. The
adopted range of acceptable solutions then defines the true uncertainty range.
The final uncertainties reported in Table The Diverse Morphology, Stellar
Population, and Black Hole Scaling Relations of the Host Galaxies of Nearby
Quasars reflect the quadrature sum of the statistical uncertainties from
GALFIT and the systematic uncertainties due to the (1) the PSF, (2) background
subtraction, and (3) model selection.

Following Kim et al. (2008a), we cross-check the decomposition results by
comparing the parametrically derived total magnitudes with those measured
nonparametrically (Figure 5). After removing the nuclear emission and close
contaminating sources based on the overall GALFIT model, we sum up all the
remaining flux to estimate the total magnitude of the host galaxy. The
uncertainty of the nonparametric magnitude is estimated by adopting different
PSF models and background subtraction. The nonparametric magnitudes usually
agree well with the sum of the parametric components, with the median
difference between the parametric and nonparametric magnitudes and their
standard deviation being $\Delta m=0.01\pm 0.19$ mag in the $I$ band. PSF
mismatch poses a more serious problem in the $B$ band and strongly affects the
nonparametric magnitudes if the host galaxy is faint. For the 20 sources for
which a meaningful comparison can be made, the $B$-band parametric magnitudes
are fainter than the nonparametric magnitudes by $\Delta m=0.16\pm 0.13$ mag.
The parametric measurements, constrained by the structural parameters based on
the $I$-band decomposition, may miss more extended emission in the $B$ band.
Nevertheless, the difference is much smaller than the uncertainty of the
integrated magnitudes, and in terms of stellar mass it amounts to only 0.07
dex, which is much smaller than the mass uncertainty.

Figure 3:  GALFIT decomposition of $I$-band images, showing (a) the 1D
profile, (b) the observed image, (c) the best-fit 2D model of the host (AGN
component excluded), and (d) the residuals. The 1D azimuthally averaged
profile shows the original data (black circles), the best-fit model (magenta
solid line), and the subcomponents (PSF: black dotted line; bulge: red dashed
line; disk: blue dashed line; bar: green dashed line; companion: gray dashed
line). All images are on an asinh stretch. The complete figure set (35 images)
is available in the online journal.

Figure 4:  GALFIT decomposition of $B$-band images, showing (a) the 1D
profile, (b) the observed image, (c) the best-fit 2D model of the host (AGN
component excluded), and (d) the residuals. The 1D azimuthally averaged
profile shows the original data (black circles), the best-fit model (magenta
solid line), and the subcomponents (PSF: black dotted line; bulge: red dashed
line; disk: blue dashed line; bar: green dashed line; companion: gray dashed
line). All images are on an asinh stretch. The complete figure set (35 images)
is available in the online journal.

Figure 5: Comparison of the integrated magnitude of the host galaxy based on
parametric and nonparametric measurement for (a) $I$ band and (b) $B$ band.
Different morphological types are given in the legend. The dashed line is the
one-to-one relation, with the dotted lines indicating a difference of $\pm
0.3$ mag. The $I$-band magnitudes show no systematic difference between the
two methods, while in the $B$ band the parametric method yields median
brightnesses that are 0.16 mag fainter than those from the nonparametric
method (see text).

##  4 Results

###  4.1 Morphological Classification

On account of the strong contamination by the AGN core, coupled with the large
distances of the sources and the relative shallowness of the HST images, the
morphology of the host galaxies is challenging to ascertain. We base our
morphological classification in part on our quantitative decomposition and in
part on visual inspection of the original images and the residual images after
subtracting the best-fit parametric model (Figures 3 and 4). Except for a
single case (Appendix B), the 11 galaxies (31% of the sample) that can be fit
adequately with a single Sérsic component are classified as ellipticals, even
if in deep observations nearby ellipticals often exhibit more than one major
photometric component (Huang et al., 2013; Zhu et al., 2021). Although the
visual classification is usually challenging when a source is at $z\gtrsim
0.2$, particularly when the nuclear emission overwhelms the starlight, the
host galaxies, robustly detected, are sufficiently luminous and massive to be
regarded securely as ellipticals. Disk galaxies are recognizable as two-
component systems with a Sérsic bulge plus an exponential disk, which is often
accompanied by an additional non-axisymmetric feature that can be readily
identified with a bar. Spiral features in the residual maps offer further
supporting evidence for the presence of a disk. Nearly half of the sample
(49%, 17/35) resides in disks, among them $\sim 35\%$ (6/17) that are barred.
Six objects (17%) exhibit irregular substructure indicative of unambiguous
ongoing or recent merger activity. Throughout the paper, we mainly focus on
galaxy major mergers with stellar mass ratio typically $\lesssim 0.3$. PG
1202+281 has a close projected neighbor, but in the absence of detectable
tidal features or redshift information for the nearby source, we do not
consider it a merger system. The single-component hosts are systematically
more luminous (median $M_{I}=-23.71$ mag) than those decomposed into two
(bulge and disk) or three (bulge, bar, and disk) components (median
$M_{I}=-22.66$ mag), consistent with the known luminosity difference between
nearby ellipticals and spirals (e.g., Burstein et al. 2005). This gives us
some confidence that the ellipticals have been properly classified. The
mergers are the most luminous, with median $M_{I}=-24.21$ mag, not too
dissimilar from the ellipticals.

Figure 6: Comparison of the total stellar mass of the host galaxies derived in
this work with (a) the stellar mass of the bulges of the disk galaxies and (b)
the total stellar mass from Zhang et al. (2016). The dashed line indicates the
1:1 relation. In panel (a), the solid line indicates the 3:1 relation; the
bulge mass is always smaller than the total stellar mass of the system. In
panel (b), the dotted lines mark $\pm 0.3$ dex; the two sets of estimates are
mostly consistent within the 0.3 dex uncertainty.

###  4.2 Stellar Mass Estimation

We estimate the total stellar mass of the host galaxy, and separately for the
bulge alone if present, from its rest-frame, extinction-corrected $I$-band
magnitude and $B-I$ color, using the color-dependent $M/L$ prescribed by Bell
& de Jong (2001). To convert the measured HST-based magnitudes to rest-frame
magnitudes in the Johnson $B$ and $I$ bands, we adopt Bruzual & Charlot (2003)
stellar population synthesis models covering a grid of stellar ages provided
by HST calibration reference
data.333http://etc.stsci.edu/etcstatic/users_guide/1_ref_2_spectral_distribution.html
The models assume an exponentially decaying star formation history with an
e-folding time of 0.6 Gyr, solar metallicity, and a Chabrier (2003) stellar
initial mass function (IMF). The model spectra are redshifted to the observed
frame of each target, corrected for foreground extinction using the extinction
values of Schlafly & Finkbeiner (2011) and the Galactic extinction curve of
Cardelli et al. (1989), and then convolved with the respective HST filter
response function using synphot444https://synphot.readthedocs.io. The model
spectrum that provides the closest match to the observed color and its
associated uncertainty is used to calculate the K-correction. Typical ages for
the bulge component range between 1 and 12 Gyr, overlapping with but
systematiacally larger than the ages of 1–5 Gyr that characterize the entire
galaxy. The uncertainty of the K-correction is typically small ($\lesssim
0.05$ mag; Kim et al. 2017). The derived quantities are not sensitive to the
star formation history of the adopted stellar synthesis model. Whether we use
the simple stellar population model or the exponential decaying model with an
e-folding time 1.0 Gyr, our derived bulge and host galaxy masses varies by
typically $\lesssim 0.1$ dex.

The calculations of Bell & de Jong (2001) are based on the IMF of Salpeter
(1955). Scaling the stellar masses by a factor of 0.61 (Madau & Dickinson,
2014) to match the IMF of Chabrier (2003),

$\mathrm{\log}\left(\frac{M_{*}}{M_{\odot}}\right)=-0.4\,(M_{I}-M_{I,\odot})+0.439\,(B-I)-0.609,$
(2)

where $M_{I,\odot}=4.07$ mag is the $I$-band absolute magnitude of the Sun
(Blanton & Roweis, 2007). The final uncertainties of the stellar masses are
significant (median $\sim 0.5$ dex; Table The Diverse Morphology, Stellar
Population, and Black Hole Scaling Relations of the Host Galaxies of Nearby
Quasars), accounting not only for the $\sim$0.3 dex uncertainty of $M/L$
(Conroy, 2013), but also for the significant uncertainty associated with the
photometric decomposition. Our subsequent analysis will use the results based
on the parametric magnitudes, except for the hosts of objects classified as
mergers, for which the nonparametric magnitudes should be more reliable when
the nuclei are properly removed. The concept of a “bulge,” of course, is not
meaningful in an ongoing merger. As all the mergers in our sample appear to be
quite substantial, for simplicity we assume that the final product of the
merger will be an elliptical, or at the very least a highly bulge-dominated
galaxy, even if a disk can reform (e.g., Peschken et al. 2020).

Despite the large uncertainties of the stellar masses, it is reassuring that
the bulge masses are always smaller than the total masses of the host (Figure
6a). The median bulge-to-total stellar mass ratio is $\sim 0.4$. While this
may seem an obvious point, it does offer some measure of reassurance that our
masses are not unreasonable despite the challenges associated with estimating
$M/L$ for the bulge and the galaxy separately, which can occasionally produce
bulge masses that formally exceed those of the whole galaxy (Bentz & Manne-
Nicholas, 2018; Li et al., 2021). Prior to this study, the work of Zhang et
al. (2016) has provided the largest collection of (total) stellar masses for
the PG quasars, derived from assembling heterogeneous, literature-based HST
photometry of 55 sources assuming a fixed $M/L$ for the optical or near-IR. We
compare our host galaxy stellar masses with those of Zhang et al. (2016) in
Figure 6b, for the subset of 22 overlapping sources.555If two stellar masses
are given in Table 1 of Zhang et al. (2016), we use their average value. We
scale their masses to our adopted IMF (Chabrier, 2003) and cosmology (Planck
Collaboration et al., 2016). The two sets of stellar masses show overall
agreement, with most sources lying within an envelope of $\pm 0.3$ dex. The
most deviant points are PSF-dominated sources and mergers. The most striking
example is PG 0953+414, for which our stellar mass of $M_{*}=1.9\times
10^{10}\,M_{\odot}$ is $\sim 10$ times lower than that reported by Zhang et
al. (2016). The optical photometry used by Zhang et al. (2016) originated from
the analysis of an HST WFPC2/F675W image by Hamilton et al. (2002), who seem
to have severely overestimated the flux of the host (Kim et al., 2017), likely
on account of the extremely dominant nuclear emission of PG 0953+414.
Similarly, our stellar mass for PG 1700+518 is a factor of $\sim 3$ lower than
that reported by Zhang et al. (2016), who based it on an HST NICMOS/F160W
observation. Bentz & Manne-Nicholas (2018) analyzed the same image for this
source and found a stellar mass similar to ours.

###  4.3 Black Hole Mass Estimation

Following Ho & Kim (2015), the BH mass is derived as

$\mathrm{\log}\left(\frac{M_{\rm
BH}}{M_{\odot}}\right)=\log\left[\left(\frac{\rm{FWHM}_{\rm
H\beta}}{1000\,\rm{km\,s^{-1}}}\right)^{2}\left(\frac{L_{5100}}{10^{44}\,\rm{erg\,s^{-1}}}\right)^{0.533}\right]+7.03,$
(3)

where we assume, for concreteness, the virial factor for classical
bulges.666Our main conclusions are qualitatively the same if we use the virial
factor for pseudo bulges (Ho & Kim, 2014). The uncertainty of the BH masses is
0.35 dex. In Equation 3, ${\rm FWHM}_{\rm H\beta}$ gives the line width of
broad H$\beta$ derived from the single-epoch spectra of Boroson & Green
(1992), and $L_{5100}\equiv\lambda L_{\lambda}(5100\,\rm{\mathring{A}})$, with
$L_{\lambda}(5100\,\rm{\mathring{A}})$ the specific luminosity of the
nonstellar continuum at 5100 Å.

We obtain $L_{5100}$ from our image analysis, converting the nuclear magnitude
from the $B$-band GALFIT decomposition to rest-frame 5100 Å flux density
assuming a power-law spectrum $f_{\nu}\propto\nu^{-0.44}$ (Vanden Berk et al.,
2001) and correcting for Galactic extinction (Schlafly & Finkbeiner, 2011).
Uncertainties arising from variations in PSF model and background subtraction
(Section 3.2) are typically less than 0.05 dex. As shown in Figure 7, our
values of $L_{5100}$ agree within $\sim 0.3$ dex with those given by
Vestergaard & Peterson (2006), who made use of the ground-based
spectrophotometry from Neugebauer et al. (1987). The lower luminosity sources
with $L_{5100}\lesssim 10^{44.5}\,\rm{erg\,s^{-1}}$, predominantly narrow-line
Seyfert 1s (NLS1s; $\mathrm{FWHM}_{\rm H\beta}\leq 2000\>\mathrm{km\>s^{-1}}$;
Osterbrock & Pogge 1985), are systematically brighter in Neugebauer et al.’s
measurements, no doubt because of the significant host galaxy light included
in the large, $15\arcsec$-diameter aperture used for the ground-based
observations. The extreme outlying point for PG 1149$-$110 requires another
explanation. Its nucleus is quite modest in strength compared to the host in
the $B$\- and $I$-band HST images, both much fainter than reported in earlier
observations (e.g., Neugebauer et al. 1987; Jahnke & Wisotzki 2003).
Inspection of multi-epoch photometry in the W1 (3.4 $\mu$m) and W2 (4.6
$\mu$m) bands recorded during 2011–2019 by WISE (Wright et al., 2010) reveals
that the light curves of PG 1149$-$110 reached a minimum in 2014, close in
time to when our HST observations were conducted, and then quickly rose by
more than 0.5 mag. This level of near-IR variability, large even by the
standards of “changing-look” AGNs, likely induces a brightnening of $\gtrsim
1$ mag in the optical (Yang et al. 2018). Therefore, we suspect that for PG
1149$-$110 intrinsic variability is the culprit for the large difference
between our measurement of its nuclear brightness and that recorded more than
20 years earlier. The BH mass for this source was computed using the value of
$L_{5100}$ from Vestergaard & Peterson (2006).

Figure 7: The 5100 Å continuum luminosity measured by our image decomposition
is compared with the spectrophotometric measurements collected by Vestergaard
& Peterson (2006). The dashed line is the 1:1 relation, with the dotted lines
indicating $\pm 0.3$ dex. The 0.05 dex uncertainties of $L_{5100}$ are small
and not plotted. Our measurement of $L_{5100}$ for PG 1149$-$110 is $\sim 20$
times lower than the previous measurement, likely because the source is
variable (see text).

##  5 Discussion

Figure 8: The relation between BH mass and the stellar mass of (a) the bulge
and (b) the entire host galaxy. In panel (a), the dashed line is the
$M_{\mathrm{BH}}-M_{\mathrm{bulge}}$ relation of classical bulges and
elliptical galaxies from Kormendy & Ho (2013; KH13), while the solid line is
the revised version from Zhu et al. (2021; ZHG21) that considers classical
bulges and only the “cores” of elliptical galaxies. In panel (b), the red and
blue dotted lines are the $M_{\mathrm{BH}}-M_{\mathrm{total}}$ relation for
early-type and late-type galaxies, respectively (Greene et al. 2020; GSH20).
The typical uncertainty of $M_{\mathrm{BH}}$ is given in the bottom-left
corner. The shaded regions indicate the intrinsic scatter of ZHG21 (0.41 dex),
GSH20 ETG (0.65 dex), and GSH20 LTG (0.65 dex). For claroty, the intrinsic
scatter of KH13 (0.28 dex) is not displayed.

###  5.1 The $M_{\rm BH}$–$M_{\rm bulge}$ and $M_{\rm BH}$–$M_{\rm total}$
Relations

Figure 8a shows the relation between BH mass and bulge stellar mass for our
sample of PG quasars. For comparison, we overplot the $M_{\rm BH}$–$M_{\rm
bulge}$ relation of nearby, inactive elliptical galaxies and classical bulges
from Kormendy & Ho (2013), as well as the revised version of this relation
that replaces ellipticals by their “core,” the component most closely linked
to the compact, high-redshift progenitors of present-day ellipticals, which
were most directly related to the growth phase of the BH (Zhu et al., 2021).
As with Kormendy & Ho (2013), the zero point of the $M/L$ used by Zhu et al.
(2021) was calibrated dynamically and does not depend on the choice of IMF.
Within the relatively large scatter of the quasar sample, both of the
relations for inactive galaxies intersect our data points for the PG quasars,
especially for the group BHs more massive than $M_{\rm BH}\gtrsim
10^{8}\,M_{\odot}$. By contrast, the objects with $M_{\rm BH}\lesssim
10^{7.5}\,M_{\odot}$, which are predominantly NLS1s hosted by disk, and
especially barred, galaxies, depart notably from the
$M_{\mathrm{BH}}-M_{\mathrm{bulge}}$ or $M_{\mathrm{BH}}-M_{\mathrm{core}}$
relations of their more massive counterparts. This is qualitatively consistent
with the known behavior of pseudo bulges in inactive galaxies (Kormendy & Ho,
2013), and suggests that the lower mass BHs reside in pseudo bulges.

While BH mass correlates most closely with bulge stellar velocity dispersion
and stellar mass (Kormendy & Ho, 2013), its relation with the total galaxy
stellar mass (Reines & Volonteri, 2015), albeit looser, nonetheless still
offers a pragmatic and observationally less demanding empirical tool to
investigate BH-galaxy coevolution. Figure 8b illustrates that the PG quasars
obey the $M_{\mathrm{BH}}-M_{\mathrm{total}}$ relation of inactive galaxies,
especially when separately divided into early-type galaxies and late-type
galaxies, as recently calibrated by Greene et al. (2020). Note that Greene et
al. (2020) derive stellar masses using the color-based $M/L$ from Bell et al.
(2003), which, on account of the “diet Salpeter” IMF assumed, need to be
scaled by 0.06 dex to be consistent with our adopted Chabrier (2003) IMF.

Figure 9: Bulge-to-total ratio distribution of our sample (red hatched
histogram) is in general consistent with that of the large sample of
$z\lesssim 0.35$ AGNs from (Kim et al., 2017, black open histogram). Figure
10: The distribution of the bulge-to-total ratio ($B/T$) and bulge Sérsic
index ($n$) measured in the $I$ band. Morphological types are given in the
legend. NLS1s are highlighted with circles. Elliptical and systems with
uncertain morphology are designated $B/T=1$, as are mergers, for illustration
purposes; for the sake of clarity, these points are moved slightly and
randomly along the $x$ axis. Shown for comparison in light grey symbols are
the 320 disk galaxies with morphological decomposition from Gao et al. (2019).

###  5.2 Diversity of Morphologies

The most luminous quasars in our sample are mainly at $z\gtrsim 0.2$ and are
thus not as well-resolved. However, so long as they are detected robustly,
their massive hosts can be securely classified as elliptical galaxies (Section
4.1). Moreover, these host galaxies usually display faint surface brightnesses
and large effective radii that follow the Kormendy (1977) relation of inactive
elliptical galaxies (Section 5.3). Therefore, we conclude that most of the
luminous ($M_{B}<-23.5$ mag or $\mbox{$L_{5100}$}\gtrsim
10^{45}\,\mathrm{erg\,s^{-1}}$) quasars in our sample likely reside in
elliptical galaxies, consistent with early HST observations (Disney et al.,
1995; Bahcall et al., 1997; McLure et al., 1999; Dunlop et al., 2003).

Figure 9 compares the distribution of bulge-to-total ratio of our sample with
that of the 235 type 1 AGNs studied by Kim et al. (2017), which, while
heterogeneous, is the largest and most representative sample to date of nearby
($z\lesssim 0.35$) AGNs with detailed morphological decomposition based on HST
imaging. The Anderson-Darling test (Scholz & Stephens 1987) cannot reject the
null hypothesis that the samples are drawn from the same parent distribution
($p$-value $>0.25$). Our sample contains a moderately higher fraction of
elliptical ($B/T=1$) galaxies, likely because our sample extends to slightly
higher redshift ($\sim 0.4$) than that of Kim et al. (2017), and it comprises
among the most luminous sources within its redshift range. As in Kim et al.
(2017), mergers make up only a minority (17%; 6/35) of the sample, emphasizing
that mergers are not a prerequisite for triggering quasar-level activity in
galaxies (Villforth et al. 2014; Hewlett et al. 2017).

Quasars hosted by disk galaxies exhibit relatively modest bulge fractions
($B/T\lesssim 0.5$) and low bulge Sérsic indices ($n\lesssim 2$). Those that
are barred, in particular, have the lowest $B/T$ (Figure 10), consistent with
the overall trend observed among nearby inactive galaxies (gray symbols) from
the Carnegie-Irvine Galaxy Survey (Ho et al., 2011) with detailed
morphological decomposition (Gao et al., 2019).777Although the analysis of Gao
et al. (2019) was based on the $R$ band, differences betweem the $R$ and $I$
band should not be large (e.g., Tasca & White 2011). In contrast, the
elliptical galaxies and major mergers have large Sérsic indices. While
rigorous distinction between pseudo and classical bulges is still
controversial (Gao et al. 2020 and references therein), pseudo bulges are
generally characterized by low $B/T$ and low Sérsic $n$. Adopting, for
simplicity, $n<2$ as a criterion for recognizing pseudo bulges, 82% (14/17) of
the PG quasars hosted by disk galaxies contain a pseudo bulge.

Our sample contains a total of nine NLS1s (Figure 10). All but one reside in
disk galaxies, among which 50% (4/8) are barred. This fraction is consistent
with that of Kim et al. (2017) but lower than that of earlier studies of
(lower luminosity) Seyfert galaxies (e.g., 65%–85%; Crenshaw et al. 2003; Ohta
et al. 2007). The bar fraction of inactive disk galaxies decreases with
increasing galaxy stellar mass (Díaz-García et al., 2016; Li et al., 2017;
Erwin, 2018; Zhao et al., 2020), averaging $\sim$50% for stellar masses $\sim
10^{10}-10^{11}\,M_{\odot}$, similar to our finding. The higher bar fraction
reported in previous NLS1 samples likely can be attributed to their lower
typical galaxy masses. The lower bar fraction of 22% (2/9) among the targets
with ${\rm FWHM}_{\rm H\beta}>2000\,\mathrm{km\>s^{-1}}$ (“broad-line”
objects), while in agreement with previous studies (Crenshaw et al., 2003;
Ohta et al., 2007; Kim et al., 2017), nevertheless may be underestimated by
observational bias (Erwin, 2018) due to the larger distances of the sample
(median $z\approx 0.14$ compared to $z\approx 0.05$ for the NLS1s).

With $R\equiv L_{\nu}({\rm 6~{}cm})/L_{\nu}(B)>100$, six of the quasars in our
sample (Table 1) can be considered radio-loud (Kellermann et al. 1989). Apart
from a single exception in a merger, the rest of the radio-loud sources are
hosted by massive ellipticals, which are systematically more massive than
their radio-quiet counterparts. Dunlop et al. (2003) and Kim et al. (2017)
reached similar conclusions based on more extensive statistics.

Figure 11: The Kormendy relation between effective radius ($R_{e}$) and mean
surface brightness within $R_{e}$ ($\langle\mu_{e,I}\rangle$) for the bulges
of AGN host galaxies. The best-fit relations from Kim & Ho (2019) are shown
for classical bulges and elliptical galaxies (solid red line), pseudo bulges
(dashed line), and AGNs (dash-dotted line). We converted the relations of Kim
& Ho (2019) from the $R$ band to the $I$ band assuming $R-I=0.7$ mag (Fukugita
et al., 1995).

###  5.3 Kormendy Relation

The effective radius ($R_{e}$) forms a strong inverse correlation with the
mean surface brightness within $R_{e}$ in elliptical galaxies (Kormendy, 1977)
and classical bulges (Carollo, 1999). By contrast, pseudo bulges depart
markedly from the Kormendy relation by manifesting lower
$\langle\mu_{e}\rangle$ at a given $R_{e}$ (Kormendy & Kennicutt, 2004;
Gadotti, 2009; Fisher & Drory, 2010; Gao et al., 2020). Early HST observations
revealed that quasar host galaxies, when properly decomposed from the bright
nucleus, follow the Kormendy relation of inactive galaxies (Dunlop et al.,
2003). However, with the vastly expanded and more diverse sample of type 1
AGNs of Kim et al. (2017), Kim & Ho (2019) found that the bulges of AGN host
galaxies tend to be more luminous than those of inactive galaxies of the same
bulge size, concluding that their bulge surface brightnesses are likely
enhanced by newly formed stars associated with the current episode of BH
accretion.

We calculate the mean effective surface brightness by averaging the flux
inside the effective radius of the bulge,

$\langle\mu_{e,I}\rangle=M_{I}+5\log\left(\frac{R_{e}}{10\,\mathrm{pc}}\right)+28.57,$
(4)

where $M_{I}$ is absolute magnitude of the bulge. As shown in Figure 11, the
bulges of PG quasars generally overlap with the distributions defined by the
classical bulges of inactive galaxies and AGNs studied by Kim & Ho (2019). At
the same time, the quasar hosts designated by us as ellipticals in this study
indeed possess large effective radii and correspondingly faint surface
brightnesses similar to their inactive counterparts (Kormendy et al., 2009;
Huang et al., 2013; Gao et al., 2020), which lends confidence in our
morphological assignment. Major mergers tend to have systematically higher
surface brightnesses than the rest of the sample, likely a reflection of
boosted central brightness from merger-induced starbursts, by analogy with
ultra-luminous infrared galaxies (Rothberg et al., 2013). At a given $R_{e}$,
pseudo bulges in inactive galaxies display lower surface brightness than
classical bulges or elliptical galaxies (Kormendy & Kennicutt, 2004; Gadotti,
2009; Fisher & Drory, 2010; Gao et al., 2020). Judging by a variety of
evidence (disk morphologies, presence of a bar, low $B/T$, low $n$), a number
of the PG quasars in our sample, particularly those classified as NLS1s,
should possess pseudo bulges, but, intriguingly, they are located mostly above
the Kormendy relation of pseudo bulges. As for the mergers, the simplest
explanation for this trend is to posit that the AGNs hosted by pseudo bulges
have central surface brightnesses enhanced by recent or ongoing star
formation. Zhao et al. (2019) reached the same conclusion studying the HST
images of nearby type 2 quasars.

The results presented here, in concert with the analysis of Kim & Ho (2019),
paint a consistent picture. The host galaxies of quasars, most notably those
with accretion rates high enough to be NLS1s, which tend to be relatively
late-type, often barred, disk galaxies with pseudo bulges, show evidence of
enhanced young stellar populations relative to their inactive counterparts.
Mergers do as well—whether they be active or not—when compared to classical
bulges and ellipticals. Nearby AGNs, notably those powerful enough to qualify
as quasars, lack neither gas nor the ability to form stars (e.g., Husemann et
al. 2017; Shangguan et al. 2018, 2020a; Jarvis et al. 2020; Yesuf & Ho 2020;
Zhuang & Ho 2020; Xie et al. 2021). BH accretion and host galaxy star
formation are tightly linked by the molecular gas (Shangguan et al., 2020b;
Zhuang et al., 2021), which may be preferentially more centrally concentrated
(Molina et al., 2021).

##  6 Conclusions

We investigate the host galaxies of 35 low-redshift ($z<0.5$) Palomar-Green
(PG) quasars with high-resolution HST/WFC3 images in the rest-frame $B$ and
$I$ bands. With the aid of multi-component 2D image decomposition, we classify
the morphological types of the host galaxies, measure the structural
parameters of their bulge, constrain the stellar population, and estimate the
stellar mass and its relation to the central BH. Our main results are as
follow.

  * •

About half of our quasar host galaxies possess a disk, and hence
morphologically can be classified as spiral or S0 galaxies, among which $\sim
1/3$ are barred. Most of the lower luminosity quasars are NLS1s, and they tend
to reside in barred disks galaxies with pseudo bulges. The more powerful
quasars ($M_{B}<-23.5$ mag or $\mbox{$L_{5100}$}\gtrsim
10^{45}\,\mathrm{erg\,s^{-1}}$) are hosted by elliptical galaxies. Mergers
account for less than 20% of the sample, with no apparent connection to the
radio-loudness of the AGN.

  * •

The host galaxy-decomposed $B$-band magnitude of the nucleus agrees well with
early spectrophotometric measurements. We use the new optical nuclear
magnitudes to update the single-epoch virial BH masses.

  * •

Using the $B-I$ color to constrain the stellar mass, we find that PG quasars
generally track the scaling relations between BH mass and stellar mass
observed in local inactive galaxies. BHs more massive than $\sim
10^{8}\,M_{\odot}$ follow the $M_{\rm BH}$–$M_{\rm bulge}$ and $M_{\rm
BH}$–$M_{\rm core}$ relations of classical bulges and ellipticals, as well as
the $M_{\rm BH}$–$M_{\rm total}$ relation of early-type galaxies; less massive
BHs, particular those in NLS1s, behave like pseudo bulges and late-type
galaxies. The correlations between BH mass and bulge stellar mass or total
galaxy stellar mass are well established in low-$z$ quasar host galaxies.

  * •

The bulges of the host galaxies of PG quasars display a relatively tight
correlation between effective radius and mean surface brightness, one that
coincides with the Kormendy relation of inactive elliptical galaxies and
classical bulges, in spite of the fact that $\sim 50\%$ of the disk and barred
galaxies likely contain pseudo bulges. In agreement with findings based on
other samples of nearby quasars, we argue that a sizable fraction of the disk
host galaxies, in particular those that possess pseudo bulges, exhibit central
regions brightened by recent or ongoing star formation.

We thank the anonymous referee for useful suggestions to improve the clarity
of the paper. This research was supported by the National Science Foundation
of China (11721303 and 11991052) and the National Key R&D Program of China
(2016YFA0400702). DZ was supported by the Scholar Program of Beijing Academy
of Science and Technology (BS202002). MK was supported by the National
Research Foundation of Korea (NRF) grant funded by the Korea government (MSIT)
(No. 2020R1A2C4001753). YZ thanks Yunfeng Chen and Niankun Yu for advice on
data analysis. JS thanks Qian Yang for discussion of quasar variability. HST

##  Appendix A WCS Correction for GYRO-mode Observations

HST pointing is fine-locked when there are two guide stars available during
the exposure, in which case the pointing can reach an accuracy of $\sim 2-5$
mas during an orbit. However, when only one guide star is available and gyros
are used to control the telescope in GYRO mode, the exposure can incur a
typical drift rate of 1–2 $\mathrm{mas\,s^{-1}}$ ($\sim
5\,\mathrm{mas\,s^{-1}}$ in rare cases; Gonzaga et al. 2012), and the final
image will be smeared if we combine the dithered frames without correcting the
world coordinate system (WCS).

Four of our targets were observed in GYRO mode (Table 1). For the typical
integration times in our program, the WCS mismatch in GYRO mode is up to $\sim
1\farcs 4$. While the drift during the exposure of each dither cannot be
removed, the relative drift between the individual dithered frames can be
corrected. The standard method to correct WCS offset is to align the images
with TweakReg (Gonzaga et al., 2012) using reference sources in the field.
However, the small field-of-view of our observations and their relatively
shallow depth conspired to yield very few reference sources in the images
suitable to for proper coordinate registration. This compelled us to align the
dithered images manually. We use the bright nucleus to adjust the WCS
reference in the image header. For the UVIS observations, we use the
unsaturated short exposure as the reference image, and the first sub-exposure
is used as the reference image for the IR observations. We subtract the other
sub-exposures from the reference image, iteratively fine-tuning their WCS to
minimize the asymmetric residuals that result from astrometric misalignment.
We estimate that this method can reach an aligment accuracy of $\sim
1.5-3\,\mathrm{mas}$. Figure A1 shows the drizzle-combined images of the
nucleus of PG 0050+124 without and with manual WCS correction.

Figure A1: The drizzle-combined WFC3/F105W image of the nucleus of PG 0050+124
(a) before and (b) after manual WCS correction.

##  Appendix B Notes on Individual Objects

PG 0050+124: The host features a grand-design spiral, with one of the arms
mildly perturbed by a small, nearby companion, which as masked during the fit.
We model the galaxy with a Sérsic bulge (free $n$) and an exponential ($n=1$)
disk. A weak, bar-like feature may be present, but adding a bar component does
not significantly improve the residuals, and therefore we do not include it.

PG 0052+251: A bar component is clearly present, but to account for it the
Sérsic index of the bulge component has to be fixed. The smallest residuals
are achieved using a bulge with $n=1$. The best-fit model leaves a residual
outer ring structure, which we do not attempt to fit.

PG 0804+761: The host has an obvious extended disk with a bar along the
direction of the diffraction spike. We fit it with three components, one for
the bulge (free $n$), one for the disk ($n=1$), and one for the bar ($n=0.5$).

PG 0923+129: The host has an obvious extended disk with a very diffuse outer
ring. We fit it with three free Sérsic components, one for a bulge, one for a
disk, and a third for an outer ring. Spiral structure can be seen in the
central $\sim 3\arcsec$, but it is too faint to be properly fit.

PG 0934+013: The host galaxy contains an obvious disk with a prominent bar and
spiral arms. The bulge is fit with $n$ set free. The bulge component in the
F438W filter is very faint and difficult to fit; its magnitude is estimated by
iteratively adjusting the amplitude of the F814W model until the residuals
appear acceptable. The $B$-band magnitude of the bulge is likely more
uncertain than in other sources, although the final $B-I$ color ($2.20\pm
0.70$ mag) does not seem unusual.

PG 0953+414: The host galaxy is very faint compared to the nucleus. We can
obtain a useful fit only with a single-component model with a fixed Sérsic
index. The best fit favors $n=1$. We deem the morphology too uncertain to be
classified.

PG 1011$-$040: The host galaxy contains an obvious disk with a prominent bar
and spiral arms. The bulge is fit with $n$ set free.

PG 1012+008: This is an obvious merging system of two, near-equal brightness
galaxies, with another possible small companion projected to the north. The
host galaxy of the AGN can be fit with two Sérsic components modified by an
$m=1$ Fourier mode. A similar model is used for the main interacting
companion. The small northern companion is fit simultaneously with a single
Sérsic component. We masked the northern companion when calculating the
nonparametric magnitude of the system.

PG 1022+519: The host galaxy can be well fit with a bulge component (free
$n$), bar, and disk modified by an $m=1$ Fourier mode.

PG 1049$-$005: The host galaxy shows obvious tidally-disrupted feature. We use
two Sérsic components modified by an $m=1$ Fourier mode to fit the host and an
additional component to account for the north tidal tail structure.

PG 1100+772: We fit the galaxy with a single component. The Sérsic index is
somewhat low ($n\approx 1.1$), and there may be mild evidence for additional
structure.

PG 1103$-$006: We fit the host galaxy with a single Sérsic model with $n$ set
free. A faint projected companion to the west of the quasar is fit
simultaneously. Other sources in the field are masked.

PG 1114+445: The galaxy is fit with two free Sérsic components for the bulge
and disk. An outer ring-like structure is not considered.

PG 1119+120: We adopt a bulge with $n=1$ and an exponential disk, plus an
additional Sérsic component to account for the faint companion to the north.
We tried to add a bar, but it did not improve the fit.

PG 1149$-$110: The host galaxy is fit with a bulge (free $n$) and disk. There
is a faint, diffuse ring in the outskirts of the galaxy.

PG 1202+281: The host galaxy can be well fit by a single Sérsic component. The
projected companion is modeled with another Sérsic component.

PG 1216+069: The nuclear emission is extremely dominant. We fit the host with
a single Sérsic component, masking the nearby, highly saturated star, which
cannot be fit simultaneously.

PG 1226+023: The nuclear emission is extremely dominant. We fit the host
galaxy with a single Sérsic component, fixing $n$ to achieve a stable
solution. Both $n=3$ and $n=4$ yield acceptable fits, and we adopt $n=4$ model
as the final model, with the difference between $n=3$ and $n=4$ models
bracketing the flux uncertainty. The well-known jet is seen in the residuals.

PG 1244+026: The host galaxy can be well fit with a bulge, bar, and disk.
Fixing the bulge Sérsic index to 1 gives the lowest residuals. The disk
contains faint spiral structure, but we do not attempt to model it. The
nucleus is slightly elongated likely due to the pointing drift (Appendix A).

PG 1259+593: The nuclear emission is very bright. We fit the galaxy with only
a single Sérsic component with $n$ fixed to 4, even though there is mild
evidence for additional faint, extended emission in the residuals.

PG 1302$-$102: There is a plume of extended emission to the north that may be
remnant of a merger. Two small, compact sources are projected toward the
center of the galaxy. Apart from accounting for the projected sources, we fit
the quasar host with two Sérsic components including Fourier modes.

PG 1322+659: This is an obvious spiral galaxy. The bulge was fit with a Sérsic
component with $n=1$, and the spiral arms were modeled with coordinate
rotation and Fourier modes.

PG 1351+640: The host galaxy can be well fit with a single Sérsic component.

PG 1354+213: The host galaxy requires two components, which together can be
interpreted to be a bulge plus disk system. However, the bulge has a very
small Sérsic index ($n\approx 0.2$); fixing it to a larger value produces
notably worse residuals.

PG 1444+407: The host galaxy requires two components, which together can be
interpreted to be a bulge plus disk system. However, the bulge has an
obviously elongated morphology and a Sérsic index suspiciously close to that
of a bar ($n\approx 0.2$). For the sake of concreteness, we attribute it to
the bulge.

PG 1448+273: The host galaxy is clearly a late-stage merger. We adopt three
Sérsic components to fit it, but there remains substantial structure in the
residual image.

PG 1534+580: The host galaxy can be well fit with a bulge (free $n$) and an
exponential disk. Very faint spiral features are seen in the residuals.

PG 1535+547: The host galaxy is fit with an $n=1$ bulge and an exponential
disk. Allowing the Sérsic index to be free yields a best fit value of
$n\approx 0.5$, but the residuals are similar to the case of $n=1$, and we
finally adopt the latter.

PG 1613+658: The host galaxy is a dramatic late-stage merger (see also Hong et
al. 2015), with a projected compact source that may be the second nucleus. We
adopt two free Sérsic components to fit the main body of the host and another
component to account for the projected companion. The distant galaxy to the
north is masked.

PG 1626+554: The host galaxy can be well fit with an $n=1$ bulge and an
exponential disk. Allowing the Sérsic index to be free yields a best fit value
of $n\approx 0.2$, but the residuals are similar to the case of $n=1$, and we
finally adopt the latter. Faint spirals are evident in the residuals.

PG 1700+518: The nuclear emission is very bright. We fit the host, which
appears to be a late-stage merger, with a single Sérsic component fixed to
$n=4$ plus and additional component to account for the arc-like feature to the
north.

PG 1704+608: The nuclear emission is very bright. We adopt a single Sérsic
component ($n=4$) to fit the host galaxy.

PG 2112+059: The nuclear emission is very bright. We adopt a single Sérsic
component ($n=4$) to fit the host galaxy.

PG 2214+139: The host galaxy can be well fitted with a single Sérsic
component. There are faint, ring-like features in the residuals, which are
related to the shells seen in the deeper ground-based image of Hong et al.
(2015).

PG 2251+113: The nuclear emission is very bright. We adopt a single Sérsic
component ($n=4$) to fit the host galaxy.

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Table 1: Object Information

Name | Alternative name | Observation | Filter | Exposure Time | $z$ | $D_{L}$ | $A_{V}$ | $M_{V}$ | $\log\,R$ | $\log\,L_{5100}$ | $\mathrm{FWHM_{H\beta}}$ | $\log\,M_{\mathrm{BH}}$
---|---|---|---|---|---|---|---|---|---|---|---|---
|  | Date | ($B$/$I$) | (s) |  | (Mpc) | (mag) | (mag) |  | ($\mathrm{erg\,s^{-1}}$) | ($\mathrm{km\,s^{-1}}$) | ($M_{\odot}$)
PG 0050+124**Observed in GYRO mode. | I Zw 1 | 2013-11-02 | F438W/F105W | 300/147 | 0.061 | 282.1 | 0.171 | -23.77 | $-0.48$ | 44.23 | 1171 | 7.29
PG 0052+251 |  | 2013-07-31 | F475W/F105W | 360/147 | 0.155 | 762.5 | 0.125 | -24.64 | $-0.62$ | 44.83 | 5187 | 8.90
PG 0804+761 |  | 2013-03-10 | F438W/F105W | 450/147 | 0.100 | 475.0 | 0.095 | -24.44 | $-0.22$ | 44.96 | 3045 | 8.51
PG 0923+129 | Mrk 705 | 2012-12-16 | F438W/F814W | 300/180 | 0.029 | 131.1 | 0.109 | -21.59 | $0.32$ | 43.38 | 1957 | 7.28
PG 0934+013 | Mrk 707 | 2013-01-29 | F438W/F814W | 450/300 | 0.050 | 229.4 | 0.184 | -21.43 | $-0.42$ | 43.51 | 1254 | 6.96
PG 0953+414 |  | 2013-02-03 | F475W/F105W | 300/147 | 0.239 | 1234.6 | 0.032 | -25.65 | $-0.36$ | 45.63 | 3111 | 8.88
PG 1011$-$040 |  | 2013-03-29 | F438W/F814W | 300/210 | 0.058 | 267.7 | 0.097 | -22.70 | $-1.00$ | 44.09 | 1381 | 7.36
PG 1012+008 |  | 2013-05-03 | F475W/F105W | 360/147 | 0.185 | 926.7 | 0.093 | -24.79 | $-0.30$ | 44.87 | 2615 | 8.33
PG 1022+519 | Mrk 142 | 2013-03-04 | F438W/F814W | 300/210 | 0.045 | 205.8 | 0.044 | -21.40 | $-0.64$ | 43.56 | 1566 | 7.19
PG 1049$-$005 |  | 2014-03-05 | F555W/F125W | 450/147 | 0.357 | 1958.1 | 0.155 | -25.93 | $-0.60$ | 45.59 | 5351 | 9.33
PG 1100+772 | 3C 249.1 | 2013-05-02 | F555W/F110W | 390/147 | 0.313 | 1680.6 | 0.096 | -25.86 | $2.51$ | 45.72 | 6151 | 9.52
PG 1103$-$006 | PKS 1103-0006 | 2013-04-12 | F555W/F125W | 540/147 | 0.425 | 2403.6 | 0.116 | -25.96 | $2.43$ | 45.73 | 6183 | 9.53
PG 1114+445 |  | 2014-03-25 | F475W/F105W | 540/147 | 0.144 | 703.6 | 0.042 | -24.01 | $-0.89$ | 44.68 | 4554 | 8.71
PG 1119+120 | Mrk 734 | 2013-05-22 | F438W/F814W | 300/180 | 0.049 | 224.7 | 0.089 | -22.49 | $-0.82$ | 43.95 | 1773 | 7.50
PG 1149$-$110 |  | 2014-02-12 | F438W/F814W | 300/210 | 0.049 | 224.7 | 0.114 | -21.90 | $-0.06$ | 42.72 | 3032 | 7.31
PG 1202+281**Observed in GYRO mode. | GQ Comae | 2013-10-28 | F475W/F105W | 360/147 | 0.165 | 816.7 | 0.056 | -23.75 | $-0.72$ | 44.29 | 5036 | 8.59
PG 1216+069 |  | 2013-01-14 | F555W/F110W | 390/147 | 0.334 | 1811.9 | 0.059 | -26.33 | $0.22$ | 45.81 | 5180 | 9.42
PG 1226+023**Observed in GYRO mode. | 3C 273 | 2013-03-17 | F547M/F105W | 360/147 | 0.158 | 778.7 | 0.055 | -27.15 | $3.06$ | 46.20 | 3500 | 9.29
PG 1244+026**Observed in GYRO mode. |  | 2014-02-16 | F438W/F814W | 360/270 | 0.048 | 219.9 | 0.068 | -21.77 | $-0.28$ | 43.42 | 721 | 6.43
PG 1259+593 |  | 2012-11-05 | F606W/F125W | 360/147 | 0.472 | 2722.7 | 0.022 | -26.82 | $-1.00$ | 46.16 | 3377 | 9.24
PG 1302$-$102 | PKS 1302-102 | 2013-04-23 | F555W/F110W | 240/147 | 0.286 | 1514.8 | 0.117 | -26.60 | $2.27$ | 45.76 | 3383 | 9.03
PG 1322+659 |  | 2013-01-05 | F475W/F105W | 540/147 | 0.168 | 833.1 | 0.051 | -24.23 | $-0.92$ | 44.81 | 2765 | 8.34
PG 1351+640 |  | 2013-01-06 | F438W/F105W | 300/147 | 0.087 | 409.6 | 0.055 | -24.08 | $0.64$ | 44.45 | 5646 | 8.77
PG 1354+213 |  | 2013-03-09 | F555W/F110W | 480/147 | 0.300 | 1600.3 | 0.067 | -24.34 | $-1.10$ | 44.87 | 4127 | 8.73
PG 1444+407 |  | 2013-04-28 | F555W/F110W | 540/147 | 0.267 | 1400.2 | 0.037 | -25.18 | $-1.10$ | 45.49 | 2457 | 8.61
PG 1448+273 |  | 2014-07-19 | F438W/F105W | 540/147 | 0.065 | 301.4 | 0.079 | -23.30 | $-0.60$ | 44.14 | 815 | 6.93
PG 1534+580 | Mrk 290 | 2013-09-20 | F438W/F814W | 225/300 | 0.030 | 135.7 | 0.038 | -21.44 | $-0.15$ | 43.65 | 5324 | 8.29
PG 1535+547 | Mrk 486 | 2013-04-22 | F438W/F814W | 180/270 | 0.038 | 172.9 | 0.038 | -22.15 | $-0.85$ | 43.72 | 1420 | 7.18
PG 1613+658 | Mrk 876 | 2013-03-25 | F475W/F105W | 450/270 | 0.129 | 624.4 | 0.071 | -24.22 | $0.00$ | 44.96 | 8441 | 9.40
PG 1626+554 |  | 2013-07-09 | F475W/F105W | 450/270 | 0.133 | 645.4 | 0.015 | -23.54 | $-0.96$ | 44.60 | 4474 | 8.65
PG 1700+518 |  | 2012-12-15 | F555W/F110W | 240/147 | 0.282 | 1490.5 | 0.089 | -26.44 | $0.37$ | 45.80 | 2185 | 8.67
PG 1704+608 | 3C 351 | 2012-12-03 | F555W/F125W | 240/147 | 0.371 | 2048.2 | 0.060 | -26.38 | $2.81$ | 45.90 | 6552 | 9.67
PG 2112+059 |  | 2013-05-18 | F606W/F125W | 300/147 | 0.466 | 2681.5 | 0.239 | -27.26 | $-0.49$ | 46.11 | 3176 | 9.16
PG 2214+139 | Mrk 304 | 2012-11-12 | F438W/F105W | 390/147 | 0.067 | 311.2 | 0.194 | -23.39 | $-1.30$ | 44.37 | 4532 | 8.54
PG 2251+113 | PKS 2251+11 | 2013-06-10 | F555W/F110W | 480/147 | 0.323 | 1742.9 | 0.229 | -26.24 | $2.56$ | 45.60 | 4147 | 9.12

Note. — Column (1): Object name. Column (2): Alternative name of the object.
Column (3): Date of observations. Column (4): Filters for rest-frame $B$ and
$I$ band. Column (5): Total long exposure time for rest-frame $B$ and $I$ band
image. Column (6): Redshift from Boroson & Green (1992). Column (7):
Luminosity distance. Column (8): Galactic extinction in the $V$ band,
estimated from Schlafly & Finkbeiner (2011). Column (9): $V$-band absolute
magnitude from Boroson & Green (1992). Column (10): Ratio of radio-to-optical
luminosity from Kellermann et al. (1989). Column (11): AGN 5100 Å continuum
luminosity converted from $B$-band nuclear emission derived from our GALFIT
decomposition. Column (12): FWHM of broad H$\beta$ line from Boroson & Green
(1992). Column (13) BH mass calculated in Section 4.3.

Table 2: Fitted Parameters

| Nucleus |  | Bulge |  | Total |  |
---|---|---|---|---|---|---|---
Name | $m_{I}$ | $m_{B}$ |  | $m_{I}$ | $m_{B}$ | $R_{e}$ | $n$ | $\mu_{e,I}$ | $\mu_{e,B}$ |  | Parametric $m_{I}$ | Nonparametric $m_{I}$ | Parametric $m_{B}$ | Nonparametric $m_{B}$ | $B/T$ |  | $A_{L}$
| (mag) | (mag) |  | (mag) | (mag) | (″) |  | (mag arcsec-2) | (mag arcsec-2) |  | (mag) | (mag) | (mag) | (mag) |  |  |
(1) | (2) | (3) |  | (4) | (5) | (6) | (7) | (8) | (9) |  | (17) | (18) | (19) | (20) | (21) |  | (22)
PG 0050$+$124 | $13.60\pm 0.05$ | $15.17\pm 0.05$ |  | $14.08\pm 0.33$ | $17.24\pm 0.32$ | $1.33\pm 0.05$ | $1.69\pm 0.05$ | $17.64\pm 0.05$ | $20.80\pm 0.05$ |  | $13.36\pm 0.32$ | $13.36\pm 0.05$ | $16.23\pm 0.35$ | $16.22\pm 0.09$ | $0.52\pm 0.05$ |  | $0.15\pm 0.05$
PG 0052$+$251 | $14.80\pm 0.06$ | $15.92\pm 0.06$ |  | $17.31\pm 0.47$ | $19.33\pm 0.51$ | $1.92\pm 0.06$ | $1$ | $21.42\pm 0.13$ | $23.37\pm 0.26$ |  | $15.75\pm 0.44$ | $15.70\pm 0.06$ | $17.66\pm 0.48$ | $17.22\pm 0.10$ | $0.24\pm 0.05$ |  | $0.03\pm 0.05$
PG 0804$+$761 | $13.56\pm 0.06$ | $14.55\pm 0.06$ |  | $17.20\pm 0.36$ | $19.15\pm 0.35$ | $0.48\pm 0.05$ | $0.58\pm 0.07$ | $18.06\pm 0.13$ | $20.01\pm 0.15$ |  | $15.33\pm 0.35$ | $15.29\pm 0.11$ | $17.77\pm 0.35$ |  $\cdots$ | $0.18\pm 0.05$ |  | $0.02\pm 0.05$
PG 0923$+$129 | $15.20\pm 0.05$ | $15.74\pm 0.05$ |  | $14.53\pm 0.32$ | $17.04\pm 0.32$ | $1.01\pm 0.05$ | $1.04\pm 0.05$ | $17.26\pm 0.05$ | $19.77\pm 0.05$ |  | $13.44\pm 0.35$ | $13.43\pm 0.05$ | $15.60\pm 0.34$ | $15.54\pm 0.05$ | $0.37\pm 0.05$ |  | $0.02\pm 0.05$
PG 0934$+$013 | $15.63\pm 0.05$ | $16.51\pm 0.05$ |  | $18.50\pm 0.50$ | $21.25\pm 0.44$ | $0.57\pm 0.09$ | $0.47\pm 0.08$ | $19.65\pm 0.11$ | $22.40\pm 0.05$ |  | $15.19\pm 0.55$ | $15.20\pm 0.05$ | $17.10\pm 0.45$ | $16.98\pm 0.05$ | $0.05\pm 0.05$ |  | $0.00\pm 0.05$
PG 0953$+$414 | $14.42\pm 0.06$ | $15.04\pm 0.05$ |  | $18.12\pm 0.59$ | $19.68\pm 0.51$ | $3.34\pm 0.83$ | $1$ | $23.43\pm 0.06$ | $24.99\pm 0.33$ |  | $18.12\pm 0.59$ | $17.77\pm 0.43$ | $19.68\pm 0.51$ |  $\cdots$ | 1 |  | $0.17\pm 0.05$
PG 1011$-$040 | $14.95\pm 0.05$ | $15.50\pm 0.05$ |  | $16.68\pm 0.45$ | $19.44\pm 0.44$ | $0.84\pm 0.07$ | $2.84\pm 0.66$ | $19.51\pm 0.11$ | $22.27\pm 0.05$ |  | $14.52\pm 0.52$ | $14.51\pm 0.05$ | $16.71\pm 0.46$ | $16.55\pm 0.05$ | $0.14\pm 0.05$ |  | $0.10\pm 0.05$
PG 1012$+$008 | $15.77\pm 0.06$ | $16.26\pm 0.06$ |  | $16.23\pm 0.50$ | $18.73\pm 0.44$ | $1.32\pm 0.74$ | $7.00\pm 2.82$ | $20.52\pm 0.96$ | $23.02\pm 0.23$ |  | $15.11\pm 0.45$ | $15.12\pm 0.05$ | $17.71\pm 0.45$ | $17.56\pm 0.12$ |  $\cdots$ |  | $0.51\pm 0.05$
PG 1022$+$519 | $15.68\pm 0.05$ | $16.33\pm 0.05$ |  | $16.63\pm 0.46$ | $18.51\pm 0.44$ | $0.43\pm 0.05$ | $1.06\pm 0.15$ | $17.52\pm 0.09$ | $19.40\pm 0.05$ |  | $14.99\pm 0.46$ | $14.98\pm 0.05$ | $17.03\pm 0.46$ | $16.93\pm 0.05$ | $0.22\pm 0.05$ |  | $0.05\pm 0.05$
PG 1049$-$005 | $14.88\pm 0.05$ | $16.02\pm 0.05$ |  | $18.13\pm 0.36$ | $20.21\pm 0.30$ | $0.87\pm 0.06$ | $0.34\pm 0.07$ | $20.08\pm 0.17$ | $22.16\pm 0.05$ |  | $16.89\pm 0.62$ | $16.66\pm 0.05$ | $18.79\pm 0.41$ |  $\cdots$ |  $\cdots$ |  | $0.54\pm 0.31$
PG 1100$+$772 | $14.60\pm 0.05$ | $15.44\pm 0.05$ |  | $16.92\pm 0.32$ | $18.48\pm 0.31$ | $2.09\pm 0.17$ | $1.08\pm 0.19$ | $21.25\pm 0.15$ | $22.81\pm 0.06$ |  | $16.92\pm 0.32$ | $16.68\pm 0.05$ | $18.48\pm 0.31$ |  $\cdots$ | 1 |  | $0.03\pm 0.05$
PG 1103$-$006 | $15.43\pm 0.05$ | $16.12\pm 0.05$ |  | $17.06\pm 0.34$ | $19.54\pm 0.30$ | $2.03\pm 0.38$ | $1.81\pm 0.45$ | $21.58\pm 0.33$ | $24.06\pm 0.05$ |  | $17.06\pm 0.34$ | $16.87\pm 0.05$ | $19.54\pm 0.30$ |  $\cdots$ | 1 |  | $0.04\pm 0.05$
PG 1114$+$445 | $15.26\pm 0.06$ | $16.23\pm 0.06$ |  | $16.28\pm 0.38$ | $18.88\pm 0.32$ | $0.97\pm 0.25$ | $2.40\pm 0.88$ | $19.34\pm 0.47$ | $21.94\pm 0.05$ |  | $15.87\pm 0.33$ | $15.82\pm 0.05$ | $18.34\pm 0.42$ |  $\cdots$ | $0.69\pm 0.08$ |  | $0.03\pm 0.05$
PG 1119$+$120 | $14.76\pm 0.05$ | $15.48\pm 0.05$ |  | $15.79\pm 0.36$ | $18.05\pm 0.32$ | $0.60\pm 0.07$ | $1$ | $17.37\pm 0.10$ | $19.63\pm 0.05$ |  | $14.62\pm 0.38$ | $14.49\pm 0.05$ | $16.82\pm 0.33$ | $16.53\pm 0.05$ | $0.34\pm 0.05$ |  | $0.06\pm 0.05$
PG 1149$-$110 | $17.60\pm 0.05$ | $18.54\pm 0.05$ |  | $14.40\pm 0.33$ | $17.25\pm 0.32$ | $2.25\pm 0.16$ | $2.28\pm 0.10$ | $19.26\pm 0.11$ | $22.11\pm 0.05$ |  | $13.87\pm 0.39$ | $13.93\pm 0.05$ | $16.27\pm 0.35$ | $16.33\pm 0.05$ | $0.61\pm 0.08$ |  | $0.12\pm 0.05$
PG 1202$+$281 | $16.37\pm 0.06$ | $17.48\pm 0.06$ |  | $16.28\pm 0.32$ | $18.55\pm 0.30$ | $1.18\pm 0.17$ | $4.50\pm 0.79$ | $20.09\pm 0.29$ | $22.36\pm 0.05$ |  | $16.28\pm 0.32$ | $16.29\pm 0.05$ | $18.55\pm 0.30$ | $18.52\pm 0.08$ | 1 |  | $0.04\pm 0.05$
PG 1216$+$069 | $14.52\pm 0.05$ | $15.40\pm 0.05$ |  | $17.70\pm 0.62$ | $18.65\pm 0.35$ | $5.35\pm 1.59$ | $1.17\pm 0.62$ | $24.11\pm 0.62$ | $25.06\pm 0.18$ |  | $17.70\pm 0.62$ | $17.52\pm 0.10$ | $18.65\pm 0.35$ |  $\cdots$ | 1 |  | $0.02\pm 0.05$
PG 1226$+$023 | $12.12\pm 0.06$ | $12.71\pm 0.08$ |  | $14.87\pm 0.68$ | $16.43\pm 0.40$ | $11.61\pm 4.13$ | $4$ | $23.58\pm 0.15$ | $25.14\pm 0.05$ |  | $14.87\pm 0.68$ | $14.56\pm 0.27$ | $16.43\pm 0.40$ | $16.27\pm 0.90$ | 1 |  | $0.09\pm 0.07$
PG 1244$+$026 | $16.11\pm 0.05$ | $16.81\pm 0.05$ |  | $18.12\pm 0.44$ | $18.72\pm 0.44$ | $0.10\pm 0.05$ | $1$ | $15.81\pm 0.05$ | $16.41\pm 0.05$ |  | $15.54\pm 0.46$ | $15.54\pm 0.05$ | $17.35\pm 0.44$ | $17.23\pm 0.05$ | $0.09\pm 0.05$ |  | $0.04\pm 0.05$
PG 1259$+$593 | $14.68\pm 0.05$ | $15.47\pm 0.05$ |  | $16.72\pm 0.54$ | $18.63\pm 0.60$ | $11.07\pm 3.33$ | $4$ | $25.33\pm 0.33$ | $27.24\pm 0.45$ |  | $16.72\pm 0.54$ | $17.55\pm 0.05$ | $18.63\pm 0.60$ | $18.45\pm 0.09$ | 1 |  | $0.09\pm 0.05$
PG 1302$-$102 | $14.45\pm 0.05$ | $15.10\pm 0.06$ |  | $17.12\pm 0.44$ | $18.92\pm 0.53$ | $1.60\pm 0.05$ | $0.87\pm 0.05$ | $20.77\pm 0.05$ | $22.55\pm 0.30$ |  | $16.12\pm 0.44$ | $15.92\pm 0.25$ | $18.07\pm 0.54$ |  $\cdots$ |  $\cdots$ |  | $0.18\pm 0.05$
PG 1322$+$659 | $15.56\pm 0.05$ | $16.24\pm 0.06$ |  | $18.08\pm 0.33$ | $20.80\pm 0.55$ | $0.58\pm 0.05$ | $1$ | $19.59\pm 0.13$ | $22.31\pm 0.45$ |  | $16.79\pm 0.32$ | $16.85\pm 0.07$ | $18.91\pm 0.38$ |  $\cdots$ | $0.30\pm 0.05$ |  | $0.03\pm 0.05$
PG 1351$+$640 | $14.71\pm 0.06$ | $15.57\pm 0.05$ |  | $14.67\pm 0.31$ | $17.23\pm 0.30$ | $1.53\pm 0.24$ | $4.98\pm 0.79$ | $19.10\pm 0.32$ | $21.66\pm 0.05$ |  | $14.67\pm 0.31$ | $14.71\pm 0.05$ | $17.23\pm 0.30$ | $16.96\pm 0.05$ | 1 |  | $0.02\pm 0.05$
PG 1354$+$213 | $16.34\pm 0.05$ | $17.49\pm 0.05$ |  | $18.92\pm 0.36$ | $20.82\pm 0.32$ | $0.55\pm 0.05$ | $0.19\pm 0.10$ | $19.71\pm 0.25$ | $21.61\pm 0.05$ |  | $17.10\pm 0.33$ | $17.11\pm 0.05$ | $19.51\pm 0.33$ | $19.21\pm 0.05$ | $0.19\pm 0.05$ |  | $0.03\pm 0.05$
PG 1444$+$407 | $14.93\pm 0.05$ | $15.69\pm 0.05$ |  | $17.99\pm 0.35$ | $19.82\pm 0.32$ | $0.95\pm 0.05$ | $0.23\pm 0.07$ | $20.02\pm 0.16$ | $21.85\pm 0.05$ |  | $16.88\pm 0.33$ | $16.82\pm 0.05$ | $18.66\pm 0.32$ |  $\cdots$ | $0.36\pm 0.05$ |  | $0.00\pm 0.05$
PG 1448$+$273 | $14.57\pm 0.06$ | $15.66\pm 0.05$ |  | $14.94\pm 0.44$ | $17.21\pm 0.44$ | $1.54\pm 0.05$ | $2.43\pm 0.15$ | $19.01\pm 0.05$ | $21.28\pm 0.05$ |  | $14.53\pm 0.44$ | $14.51\pm 0.05$ | $16.76\pm 0.45$ | $16.60\pm 0.05$ |  $\cdots$ |  | $0.62\pm 0.05$
PG 1534$+$580 | $14.58\pm 0.05$ | $15.23\pm 0.05$ |  | $15.57\pm 0.40$ | $17.87\pm 0.30$ | $0.73\pm 0.18$ | $1.61\pm 0.48$ | $17.81\pm 0.35$ | $20.11\pm 0.05$ |  | $14.30\pm 0.34$ | $14.28\pm 0.05$ | $16.37\pm 0.31$ | $16.23\pm 0.05$ | $0.31\pm 0.05$ |  | $0.03\pm 0.05$
PG 1535$+$547 | $14.41\pm 0.05$ | $15.58\pm 0.05$ |  | $16.89\pm 0.53$ | $19.60\pm 0.35$ | $0.99\pm 0.28$ | $1$ | $19.56\pm 0.11$ | $22.27\pm 0.15$ |  | $15.47\pm 0.35$ | $15.28\pm 0.05$ | $17.74\pm 0.33$ | $17.50\pm 0.05$ | $0.27\pm 0.08$ |  | $0.02\pm 0.05$
PG 1613$+$658 | $14.46\pm 0.05$ | $15.22\pm 0.05$ |  | $14.05\pm 0.36$ | $16.64\pm 0.33$ | $5.99\pm 1.91$ | $6.11\pm 0.99$ | $21.55\pm 0.56$ | $24.14\pm 0.06$ |  | $13.71\pm 0.34$ | $13.82\pm 0.05$ | $16.27\pm 0.34$ | $16.29\pm 0.06$ |  $\cdots$ |  | $0.30\pm 0.05$
PG 1626$+$554 | $15.37\pm 0.06$ | $16.26\pm 0.06$ |  | $18.16\pm 0.34$ | $20.70\pm 0.59$ | $0.71\pm 0.05$ | $1$ | $20.11\pm 0.11$ | $22.65\pm 0.49$ |  | $15.93\pm 0.32$ | $15.91\pm 0.05$ | $18.46\pm 0.44$ | $18.08\pm 0.18$ | $0.13\pm 0.05$ |  | $0.01\pm 0.05$
PG 1700$+$518 | $13.99\pm 0.05$ | $14.99\pm 0.05$ |  | $17.46\pm 0.36$ | $18.63\pm 0.30$ | $2.09\pm 0.88$ | $4$ | $22.45\pm 0.93$ | $23.62\pm 0.05$ |  | $16.94\pm 0.32$ | $16.73\pm 0.05$ | $18.45\pm 0.30$ |  $\cdots$ |  $\cdots$ |  | $0.19\pm 0.05$
PG 1704$+$608 | $14.17\pm 0.05$ | $15.44\pm 0.05$ |  | $16.42\pm 0.37$ | $17.93\pm 0.34$ | $6.38\pm 1.81$ | $4$ | $23.83\pm 0.38$ | $25.34\pm 0.15$ |  | $16.42\pm 0.37$ | $16.57\pm 0.05$ | $17.93\pm 0.34$ |  $\cdots$ | 1 |  | $0.12\pm 0.05$
PG 2112$+$059 | $14.31\pm 0.05$ | $15.34\pm 0.05$ |  | $18.12\pm 0.53$ | $19.71\pm 0.98$ | $5.35\pm 2.05$ | $4$ | $25.15\pm 0.80$ | $26.74\pm 0.89$ |  | $18.12\pm 0.53$ | $17.96\pm 0.05$ | $19.71\pm 0.98$ |  $\cdots$ | 1 |  | $0.01\pm 0.18$
PG 2214$+$139 | $13.75\pm 0.06$ | $14.99\pm 0.05$ |  | $13.77\pm 0.31$ | $16.88\pm 0.31$ | $3.44\pm 0.32$ | $2.42\pm 0.13$ | $19.58\pm 0.17$ | $22.69\pm 0.06$ |  | $13.77\pm 0.31$ | $13.77\pm 0.05$ | $16.88\pm 0.31$ |  $\cdots$ | 1 |  | $0.01\pm 0.05$
PG 2251$+$113 | $14.53\pm 0.05$ | $15.66\pm 0.05$ |  | $16.53\pm 0.34$ | $18.26\pm 0.32$ | $4.96\pm 1.93$ | $4$ | $23.40\pm 0.63$ | $25.13\pm 0.10$ |  | $16.53\pm 0.34$ | $16.82\pm 0.05$ | $18.26\pm 0.32$ |  $\cdots$ | 1 |  | $0.07\pm 0.05$

Note. — Column (1): Object name. Column (2): Apparent nuclear magnitude in
restframe $I$ band. Column (3): Apparent nuclear magnitude in restframe $B$
band. Column (4): Apparent bulge magnitude in restframe $I$ band. Column (5):
Apparent bulge magnitude in restframe $B$ band. Column (6): Effective radius
of bulge in restframe $I$ band. Column (7): Sérsic index of bulge component.
Column (8): Surface brightness of bulge at the effective radius in restframe
$I$ band. Column (9): Surface brightness of bulge at the effective radius in
restframe $B$ band. Column (10): Flags on disk (d), bar (b), merger companions
(c) or tidal features (f). Column (11): Apparent magnitude of disk, bar,
merger companions or tidal features for restframe $I$ band image. Column (12):
Apparent magnitude of disk, bar, merger companions or tidal features for
restframe $B$ band image. Column (13): Effective radius of disk, bar, merger
companions or tidal features in restframe $I$ band. Column (14): Sérsic index
of disk, bar, merger companions or tidal features. Column (15): Surface
brightness of disk, bar, merger companions or tidal features at the effective
radius in restframe $I$ band. Column (16): Surface brightness of disk, bar,
merger companions or tidal features at the effective radius in restframe $B$
band. Column (17): Parametric host magnitude of restframe $I$ band image from
combining the luminosity of the bulge and disk, bar, merger companions or
tidal features. Column (18): Host magnitude measured nonparametrically after
subtracting the best-fit nucleus in restframe $I$ band . Column (19):
Parametric host magnitude of restframe $B$ band image from combining the
luminosity of the bulge and disk, bar, merger companions or tidal features.
Column (20): Host magnitude measured nonparametrically after subtracting the
best-fit nucleus in restframe $B$ band . Column (21): Bulge-to-total ratio.
Column (22): Absolute amplitude of the first Fourier mode, which measures the
lopsidedness of the galaxy. Columns (10) – (16) only appear in the machine-
readable version of the table.

Table 3: Physical Parameters

|  | Bulge |  | Total
---|---|---|---|---
Name | Morphology | $M_{B}$ | $M_{I}$ | $(B-I)$ | $\log\ M_{\rm bulge}$ |  | $M_{B}$ | $M_{I}$ | $(B-I)$ | $\log\ M_{\rm total}$
|  | (mag) | (mag) | (mag) | ($M_{\odot}$) |  | (mag) | (mag) | (mag) | ($M_{\odot}$)
(1) | (2) | (3) | (4) | (5) | (6) |  | (7) | (8) | (9) | (10)
PG 0050$+$124 | D | $-20.72\pm 0.36$ | $-22.69\pm 0.34$ | $1.97\pm 0.50$ | $10.96\pm 0.49$ |  | $-21.72\pm 0.38$ | $-23.42\pm 0.33$ | $1.70\pm 0.50$ | $11.13\pm 0.49$
PG 0052$+$251 | B | $-20.44\pm 0.51$ | $-21.77\pm 0.48$ | $1.33\pm 0.70$ | $10.31\pm 0.63$ |  | $-22.11\pm 0.49$ | $-23.33\pm 0.45$ | $1.22\pm 0.67$ | $10.89\pm 0.60$
PG 0804$+$761 | B | $-19.64\pm 0.41$ | $-20.90\pm 0.37$ | $1.26\pm 0.55$ | $9.93\pm 0.53$ |  | $-21.23\pm 0.41$ | $-22.67\pm 0.36$ | $1.44\pm 0.55$ | $10.72\pm 0.52$
PG 0923$+$129 | D | $-19.00\pm 0.32$ | $-21.16\pm 0.32$ | $2.16\pm 0.45$ | $10.43\pm 0.47$ |  | $-20.39\pm 0.35$ | $-22.25\pm 0.35$ | $1.86\pm 0.49$ | $10.74\pm 0.50$
PG 0934$+$013 | B | $-16.29\pm 0.49$ | $-18.49\pm 0.50$ | $2.20\pm 0.70$ | $9.38\pm 0.64$ |  | $-20.21\pm 0.49$ | $-21.77\pm 0.55$ | $1.56\pm 0.74$ | $10.41\pm 0.67$
PG 0953$+$414 | U | $-21.14\pm 0.55$ | $-22.09\pm 0.59$ | $0.95\pm 0.81$ | $10.27\pm 0.72$ |  | $-21.14\pm 0.55$ | $-22.09\pm 0.59$ | $0.95\pm 0.81$ | $10.27\pm 0.72$
PG 1011$-$040 | B | $-18.39\pm 0.45$ | $-20.61\pm 0.45$ | $2.22\pm 0.64$ | $10.24\pm 0.59$ |  | $-21.01\pm 0.48$ | $-22.76\pm 0.52$ | $1.75\pm 0.71$ | $10.89\pm 0.65$
PG 1012$+$008 | M | $-22.69\pm 0.15$ | $-24.40\pm 0.07$ | $1.71\pm 0.17$ | $11.53\pm 0.32$ |  | $-22.69\pm 0.15$ | $-24.40\pm 0.07$ | $1.71\pm 0.17$ | $11.53\pm 0.32$
PG 1022$+$519 | B | $-18.36\pm 0.47$ | $-20.01\pm 0.46$ | $1.65\pm 0.66$ | $9.75\pm 0.60$ |  | $-19.97\pm 0.48$ | $-21.67\pm 0.46$ | $1.70\pm 0.66$ | $10.43\pm 0.61$
PG 1049$-$005 | M | $-23.06\pm 0.44$ | $-24.06\pm 0.63$ | $1.00\pm 0.77$ | $11.08\pm 0.71$ |  | $-23.06\pm 0.44$ | $-24.06\pm 0.63$ | $1.00\pm 0.77$ | $11.08\pm 0.71$
PG 1100$+$772 | E | $-22.89\pm 0.33$ | $-23.86\pm 0.32$ | $0.97\pm 0.46$ | $10.99\pm 0.47$ |  | $-22.89\pm 0.33$ | $-23.86\pm 0.32$ | $0.97\pm 0.46$ | $10.99\pm 0.47$
PG 1103$-$006 | E | $-23.07\pm 0.38$ | $-24.29\pm 0.34$ | $1.22\pm 0.51$ | $11.27\pm 0.50$ |  | $-23.07\pm 0.38$ | $-24.29\pm 0.34$ | $1.22\pm 0.51$ | $11.27\pm 0.50$
PG 1114$+$445 | D | $-20.62\pm 0.32$ | $-22.53\pm 0.39$ | $1.91\pm 0.50$ | $10.87\pm 0.51$ |  | $-21.16\pm 0.42$ | $-22.94\pm 0.34$ | $1.78\pm 0.54$ | $10.98\pm 0.51$
PG 1119$+$120 | D | $-19.22\pm 0.35$ | $-21.09\pm 0.36$ | $1.87\pm 0.50$ | $10.28\pm 0.50$ |  | $-20.45\pm 0.35$ | $-22.26\pm 0.38$ | $1.81\pm 0.52$ | $10.72\pm 0.51$
PG 1149$-$110 | D | $-20.15\pm 0.33$ | $-22.50\pm 0.33$ | $2.35\pm 0.47$ | $11.05\pm 0.48$ |  | $-21.05\pm 0.38$ | $-23.02\pm 0.39$ | $1.97\pm 0.54$ | $11.09\pm 0.53$
PG 1202$+$281 | E | $-21.32\pm 0.31$ | $-22.93\pm 0.32$ | $1.61\pm 0.45$ | $10.90\pm 0.47$ |  | $-21.32\pm 0.31$ | $-22.93\pm 0.32$ | $1.61\pm 0.45$ | $10.90\pm 0.47$
PG 1216$+$069 | E | $-22.88\pm 0.35$ | $-23.23\pm 0.62$ | $0.35\pm 0.71$ | $10.47\pm 0.68$ |  | $-22.88\pm 0.35$ | $-23.23\pm 0.62$ | $0.35\pm 0.71$ | $10.47\pm 0.68$
PG 1226$+$023 | E | $-22.73\pm 0.44$ | $-24.24\pm 0.69$ | $1.51\pm 0.82$ | $11.38\pm 0.76$ |  | $-22.73\pm 0.44$ | $-24.24\pm 0.69$ | $1.51\pm 0.82$ | $11.38\pm 0.76$
PG 1244$+$026 | B | $-18.20\pm 0.44$ | $-18.65\pm 0.44$ | $0.45\pm 0.62$ | $8.68\pm 0.58$ |  | $-19.71\pm 0.47$ | $-21.26\pm 0.46$ | $1.55\pm 0.66$ | $10.20\pm 0.60$
PG 1259$+$593 | E | $-23.73\pm 0.61$ | $-24.90\pm 0.54$ | $1.17\pm 0.81$ | $11.49\pm 0.71$ |  | $-23.73\pm 0.61$ | $-24.90\pm 0.54$ | $1.17\pm 0.81$ | $11.49\pm 0.71$
PG 1302$-$102 | M | $-23.11\pm 0.54$ | $-24.36\pm 0.44$ | $1.25\pm 0.70$ | $11.31\pm 0.62$ |  | $-23.11\pm 0.54$ | $-24.36\pm 0.44$ | $1.25\pm 0.70$ | $11.31\pm 0.62$
PG 1322$+$659 | D | $-19.19\pm 0.56$ | $-21.12\pm 0.33$ | $1.93\pm 0.65$ | $10.32\pm 0.56$ |  | $-21.01\pm 0.39$ | $-22.47\pm 0.33$ | $1.46\pm 0.51$ | $10.65\pm 0.50$
PG 1351$+$640 | E | $-21.33\pm 0.35$ | $-22.98\pm 0.32$ | $1.65\pm 0.47$ | $10.94\pm 0.48$ |  | $-21.33\pm 0.35$ | $-22.98\pm 0.32$ | $1.65\pm 0.47$ | $10.94\pm 0.48$
PG 1354$+$213 | D | $-20.47\pm 0.33$ | $-21.67\pm 0.36$ | $1.20\pm 0.49$ | $10.21\pm 0.50$ |  | $-21.78\pm 0.34$ | $-23.49\pm 0.33$ | $1.71\pm 0.47$ | $11.17\pm 0.48$
PG 1444$+$407 | D | $-21.05\pm 0.32$ | $-22.28\pm 0.36$ | $1.23\pm 0.48$ | $10.47\pm 0.49$ |  | $-22.21\pm 0.32$ | $-23.39\pm 0.34$ | $1.18\pm 0.47$ | $10.89\pm 0.48$
PG 1448$+$273 | M | $-21.24\pm 0.05$ | $-22.48\pm 0.05$ | $1.24\pm 0.07$ | $10.56\pm 0.31$ |  | $-21.24\pm 0.05$ | $-22.48\pm 0.05$ | $1.24\pm 0.07$ | $10.56\pm 0.31$
PG 1534$+$580 | D | $-18.12\pm 0.32$ | $-20.15\pm 0.40$ | $2.03\pm 0.51$ | $9.97\pm 0.52$ |  | $-19.60\pm 0.32$ | $-21.42\pm 0.34$ | $1.82\pm 0.47$ | $10.39\pm 0.48$
PG 1535$+$547 | D | $-17.03\pm 0.36$ | $-19.38\pm 0.53$ | $2.35\pm 0.64$ | $9.80\pm 0.62$ |  | $-18.83\pm 0.35$ | $-20.79\pm 0.35$ | $1.96\pm 0.49$ | $10.20\pm 0.50$
PG 1613$+$658 | M | $-22.92\pm 0.06$ | $-24.72\pm 0.05$ | $1.80\pm 0.08$ | $11.70\pm 0.31$ |  | $-22.92\pm 0.06$ | $-24.72\pm 0.05$ | $1.80\pm 0.08$ | $11.70\pm 0.31$
PG 1626$+$554 | D | $-18.53\pm 0.59$ | $-20.43\pm 0.35$ | $1.90\pm 0.69$ | $10.03\pm 0.59$ |  | $-20.77\pm 0.44$ | $-22.66\pm 0.33$ | $1.89\pm 0.55$ | $10.91\pm 0.52$
PG 1700$+$518 | M | $-22.59\pm 0.31$ | $-23.56\pm 0.33$ | $0.97\pm 0.45$ | $10.87\pm 0.47$ |  | $-22.59\pm 0.31$ | $-23.56\pm 0.33$ | $0.97\pm 0.45$ | $10.87\pm 0.47$
PG 1704$+$608 | E | $-23.95\pm 0.34$ | $-24.60\pm 0.37$ | $0.65\pm 0.50$ | $11.14\pm 0.51$ |  | $-23.95\pm 0.34$ | $-24.60\pm 0.37$ | $0.65\pm 0.50$ | $11.14\pm 0.51$
PG 2112$+$059 | E | $-22.70\pm 0.99$ | $-23.57\pm 0.53$ | $0.87\pm 1.12$ | $10.83\pm 0.87$ |  | $-22.70\pm 0.99$ | $-23.57\pm 0.53$ | $0.87\pm 1.12$ | $10.83\pm 0.87$
PG 2214$+$139 | E | $-21.36\pm 0.35$ | $-23.24\pm 0.32$ | $1.88\pm 0.47$ | $11.14\pm 0.48$ |  | $-21.36\pm 0.35$ | $-23.24\pm 0.32$ | $1.88\pm 0.47$ | $11.14\pm 0.48$
PG 2251$+$113 | E | $-23.35\pm 0.34$ | $-24.37\pm 0.34$ | $1.02\pm 0.48$ | $11.22\pm 0.49$ |  | $-23.35\pm 0.34$ | $-24.37\pm 0.34$ | $1.02\pm 0.48$ | $11.22\pm 0.49$

Note. — Column (1): Object name. Column (2): Galaxy morphology: D = unbarred
disk galaxy; B = barred disk galaxy; E = elliptical galaxy; M = merger; U =
uncertain. Column (3): $B$-band absolute magnitude of the bulge. Column (4):
$I$-band absolute magnitude of the bulge. Column (5): $B-I$ color of the
bulge. Column (6): Mass of the bulge. Column (7): $B$-band absolute magnitude
of the host galaxy. Column (8): $I$-band absolute magnitude of the host
galaxy. Column (9): $B-I$ color of the host galaxy. Column (10): Total mass of
the host galaxy.