11institutetext: Institut für Astrophysik, Universität Göttingen, Friedrich-
Hund Platz 1, D-37077 Göttingen, Germany
11email<EMAIL_ADDRESS>22institutetext:
Astronomisches Institut, Ruhr-Universität Bochum, Universitätsstrasse 150,
44801 Bochum, Germany 33institutetext: School of Physics & Astronomy and the
Wise Observatory, The Raymond and Beverly Sackler Faculty of Exact Sciences
Tel-Aviv University, Tel-Aviv 69978, Israel 44institutetext: Technion,
Institute of Technology, Haifa 3200003, Israel 55institutetext: Physics
Department and the Haifa Research Center for Theoretical Physics and
Astrophysics, University of Haifa, Haifa 3498838, Israel 66institutetext: SUPA
School of Physics and Astronomy, University of St. Andrews, St. Andrews, Fife,
KY16 9SS, UK 77institutetext: Max-Planck-Institut für Sonnensystemforschung,
Justus-von-Liebig-Weg 3, D-37077 Göttingen, Germany 88institutetext: Instituto
de Astronomía, Universidad Católica del Norte, Avenida Angamos 0610,
Antofagasta, Chile

# The Great Slump: Mrk 926 reveals discrete and varying Balmer line satellite
components during a drastic phase of decline ††thanks: Based on observations
obtained with the Hobby-Eberly Telescope. ††thanks: The reduced spectra and
the photometric light curve are only available in electronic form at the CDS
via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via
http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/.

Wolfram Kollatschny 11 Martin W. Ochmann 1122 Shai Kaspi 33 Claas Schumacher
11 Ehud Behar 44 Doron Chelouche 55 Keith Horne 66 Björn Müller 1177 Stephen
E. Rafter 44 Rolf Chini 2288 Martin Haas and Malte A. Probst 2211

(Received August 12, 2021 ; Accepted October 11, 2021)

###### Abstract

Aims. Mrk 926 is known to be a highly variable active galactic nucleus.
Furthermore, it is known to show very broad-line profiles. We intended to
study the continuum and line profile variations of this object with high
temporal resolution in order to determine its broad-line region (BLR)
structure and to derive its black hole mass.

Methods. We carried out a high-cadence spectroscopic variability campaign of
Mrk 926 with the 10m HET telescope, aided by photometric V-band data taken
with the C18 telescope at the Wise Observatory, over a period of about five
months. We extracted spectroscopic continuum and line light curves, and
computed cross-correlation functions (CCFs) as well as velocity-resolved CCFs
with respect to the combined spectroscopic and photometric V-band light curve.

Results. The continuum luminosity of Mrk 926 showed a drastic decrease during
our campaign. The luminosity dropped to less than 50% of its original
luminosity within only 2.5 months. Furthermore, the spectra of Mrk 926 show
complex and very broad Balmer line profiles, including outer Balmer satellites
ranging from $\pm 5000$ to $\pm 13\,000$ km s-1. The integrated H$\alpha$,
H$\beta$, and He i $\lambda 5876$ line light curves are delayed relative to
the continuum light curve. The H$\alpha$ and H$\beta$ lines show two velocity-
delay structures in the central part of their line profile (within $\pm 5000$
km s-1), at $\sim 10$ and $\sim 57$ light-days and at $\sim 5$ and $\sim 48$
light-days, respectively. These structures might be interpreted as the upper
and lower halves of an ellipse in the velocity-delay plane, which might be the
signature of a line-emitting ring, inclined by $\sim 50^{\circ}$ to the line
of sight and orbiting the black hole at radii, $R$, of 33.5 and 26.5 light-
days. We determined continuum luminosities, $\log(\lambda\,L_{\lambda}/{\rm
erg\leavevmode\nobreak\ s}^{-1}),$ of 43.68 to 44.13, which are in good
agreement with the established $R_{\text{BLR}}-L_{\text{AGN}}$ relation.
Adopting delays of 33.5 and 26.5 days for H$\alpha$ and H$\beta$,
respectively, we derive a black hole mass of $(1.1\pm 0.2)\times
10^{8}M_{\odot}$; this indicates a low Eddington ratio, which decreased from 8
to 3 percent during our campaign. The Balmer satellite components show a
higher correlation coefficient with respect to the continuum than the central
line profile, and their response to the continuum variations is on the order
of only $3-5$ days. We attribute this to the central line segment and the
Balmer satellites having different, spatially distinct regions of origin.

###### Key Words.:

Galaxies: active – Galaxies: Seyfert – Galaxies: nuclei – Galaxies:
individual: Mrk926 – (Galaxies:) quasars: emission lines

##  1 Introduction

Active galactic nuclei (AGN) are variable in all frequency bands on timescales
of days to years. It is generally accepted that the variability of the central
ionizing continuum source causes line intensity variations in the broad
emission lines. Furthermore, changes in the kinematics and structure in the
broad-line region (BLR) as well as obscuration can induce variations in the
line profiles. However, many details of the innermost BLR, such as size,
structure, and kinematics, are poorly understood. In addition, most of the
variability campaigns carried out in the past were devoted to Seyfert galaxies
that show broad emission lines with full widths at half maximum (FWHMs) of
1000 to 4000 km s-1, and only very few reverberation campaigns have been
carried out for AGN with very broad lines (FWHM $\gg$ 4000 km s-1; e.g.,
3C390.3: Sergeev et al., 2002; Gezari et al., 2007, and other AGN therein).
Broad-line reverberation studies exist for more than 50 Seyfert galaxies
(e.g., Peterson et al., 2002; Kaspi et al., 2005; Denney et al., 2009; Barth
et al., 2011; Bentz et al., 2013; Shapovalova et al., 2013; Shen et al., 2016;
Du et al., 2018, and references therein). However, detailed velocity-resolved
reverberation mapping studies have been carried out for only about a dozen
Seyfert galaxies (e.g., Mrk 110: Kollatschny (2003); Arp 151: Bentz et al.
(2008); Mrk 335, Mrk 1501, 3C 120, PG 2130: Grier et al. (2013b); 3C120:
Kollatschny et al. (2014); Mrk 335, Mrk 1044, Mrk 142: Xiao et al. (2018);
HE1136-2304: Kollatschny et al. (2018); NGC 5548: Horne et al. (2021) and
references therein).

Mrk 926 can be classified as either a strong Seyfert 1 galaxy or a quasar
based on its luminosity during the campaign presented here (see Sect. 4.4).
Its Balmer and He i lines have widths (FWHM) of more than 5000 km s-1. In a
study on the long-term variability of very broad-line AGN (Kollatschny et al.,
2006), Mrk 926 additionally exhibited strong H$\beta$ line variability
amplitudes. The nuclear activity in Mrk 926 was detected by Ward et al. (1978)
when they took a spectrum of the optical counterpart of a strong X-ray source.
Mrk 926 is a compact source on direct images in the optical (Garnier et al.,
1996). We show an SDSS (DR16: Ahumada et al., 2020) image of Mrk 926 in Fig.
3. Based on SDSS r-band images, it has an isophotal diameter (major axis) of
57.12 arcsec. This corresponds to 53 kpc. The deVaucouleurs diameter amounts
to 6.00 arcsec, corresponding to 5.56 kpc111https://ned.ipac.caltech.edu/. Mrk
926 is the brighter galaxy member (R.A. = 23:04:43.5, Dec. = $-$08:41:09
(2000), $z=0.04701$) in a double system with PGC 1000273. This close starburst
companion is located one arcmin to the south.

Figure 1: Composite SDSS image of Mrk 926 taken from the SDSS
SkyServer333https://skyserver.sdss.org. The crosshairs are centered on the
nucleus.

Osterbrock & Shuder (1982) published the first emission-line profiles of Mrk
926. They reported that the lines did not vary between July 1978 and December
1979. We were able to show that Mrk 926 varied by a factor of 2 in the
continuum as well as by a factor of 2.5 in the Balmer lines over a period of
eight years between 1990 and 1997 (Kollatschny et al., 2006). In a follow-up
paper (Kollatschny & Zetzl, 2010), we focused on line profile variations in
Mrk 926 as the line profiles of AGN and their variations can give us
information about the structure of the line-emitting region. We intensively
studied the variations in individual line segments of the Balmer lines and the
changes in the profiles during a variability campaign in 2005 and another
shorter campaign conducted in 2004.

Here, we selected Mrk 926 (MCG-2-58-22) as a target for an even more detailed,
combined spectroscopic and photometric variability study of an AGN that shows
very broad emission lines. The paper is structured as follows: In Sect. 2 we
describe the observations and data reduction. The results of the variability
and reverberation mapping study are presented in Sect. 3. We discuss the
results in Sect. 4 and give a summary in Sect. 5. Throughout this paper, we
assume a $\Lambda$ cold dark matter cosmology with a Hubble constant of
$H_{0}$ = 73 km s-1 Mpc-1, $\Omega_{\Lambda}$ = 0.73, and $\Omega_{\rm M}$ =
0.27. With a redshift of $z=0.04701,$ this results in a luminosity distance of
$D_{\rm L}$ = 200.2 Mpc = 6.18 $\times 10^{26}$ cm using the cosmology
calculator developed by Wright (2006).

##  2 Observations and data reduction

###  2.1 Spectroscopic observations with HET

We took optical spectra of Mrk 926 with the 9.2 m Hobby-Eberly Telescope (HET)
located at the McDonald Observatory at 31 epochs between 2010 July 19 and
November 19. Thus, we obtained 31 spectra over a period of 122.7 days. The
date of the campaign was fixed by the allocated observing time slots. The
average interval between the observations was 4.1 days. Excluding the first
two observations, the average time interval between observations reduces to
2.7 days with a majority of consecutive observations only 1-2 days apart. The
log of our spectroscopic observations is given in Table 1. All spectroscopic
observations were taken with identical instrumental setups. We used the
Marcario Low Resolution Spectrograph (LRS) attached at the prime focus of HET.
The detector was a 3072x1024 15 $\mu$m pixel Ford Aerospace charge-coupled
device (CCD) with 2x2 binning. The spectra cover the wavelength range from
4160 Å to 6930 Å (LRS grism 2 configuration) in the rest frame of the galaxy
(and 4350 Å to 7250 Å observed frame) with a resolving power of 650 at 5000 Å
(7.7 Å FWHM). Nearly all observations were taken with an exposure time of 20
minutes, which in most cases yielded a S/N of $>100$ per resolution element in
the continuum. The slit width was fixed to 2″. 0 projected on the sky at an
optimized position angle to minimize differential refraction. Furthermore, all
observations were taken at the same airmass owing to the particular design of
the HET. Typical seeing values ranged from 1″. 4 to 1″. 9.

In addition to the spectra of Mrk 926, necessary calibration images including
bias, flat, and arc frames (CdHgAr and Xe), as well as spectrophotometric
standard stars (Hiltner 600) for flux calibration were taken. We reduced the
spectra (bias subtraction, cosmic ray correction, flat-field correction,
2D-wavelength calibration, night sky subtraction, flux calibration, etc.) in a
homogeneous way with standard IRAF444https://iraf.net/ reduction packages
(Tody, 1986, 1993) like we did in previous publications (e.g., Kollatschny et
al., 2001). In addition, we corrected for atmospheric absorption in the oxygen
B band by fitting the continuum of the standard star Hiltner 600 and using
this fit to normalize its spectrum. In this way, we created a template of the
oxygen absorption that we scaled to each object spectrum individually. The
spatial resolution per binned pixel was 0″. 472. We extracted four columns in
our object spectra. This conforms to the slit width. Great care was taken to
ensure high-quality intensity and wavelength calibrations to keep the
intrinsic measurement errors very low (e.g., Kollatschny et al., 2001;
Kollatschny & Zetzl, 2010). The spectra were not always taken under
photometric conditions. Therefore, we calibrated all spectra to the same
absolute forbidden narrow-line fluxes under the assumption of a non-varying
narrow-line flux on timescales of years. We chose [O iii] $\lambda$4959 and [O
iii] $\lambda$5007, as well as [O i] $\lambda$6300/63, [N ii]
$\lambda$6548/84, and [S ii] $\lambda$6717/31 as calibration lines for the
blue and red section of the spectrum, respectively. The blue part of the
spectrum was scaled to an [O iii] $\lambda$5007 flux of $3.14\times
10^{-13}\rm erg\,s^{-1}\,cm^{-2}$. The same flux value has been used before
for Mrk 926 in Kollatschny et al. (2006) and Kollatschny & Zetzl (2010).
Durret & Bergeron (1988) and Morris & Ward (1988) derived similar [O iii]
$\lambda$5007 fluxes from their Mrk 926 spectra: $2.05\times 10^{-13}\rm
erg\,s^{-1}\,cm^{-2}$ and $3.31\times 10^{-13}\rm erg\,s^{-1}\,cm^{-2}$,
respectively. We note that hereafter all continuum flux densities and line
fluxes at wavelengths $>\leavevmode\nobreak\ 6500$ Å (observed frame) refer to
the intercalibration with respect to [O i] $\lambda$6300/63, [N ii]
$\lambda$6548/84, and [S ii] $\lambda$6717/31. The accuracy of the flux
calibration was tested for all forbidden emission lines in the spectra. We
calculated difference spectra for all epochs with respect to the mean spectrum
of our variability campaign. Corrections for small spectral shifts ($<$ 0.5 Å)
and small scaling factors were implemented by minimizing the residuals of the
narrow emission lines in the difference spectra. A relative flux accuracy of 1
– 2% was obtained for most of our spectra.

Table 1: Log of spectroscopic HET observations of Mrk 926. Mod. Julian Date |  UT Date |  Exp. time
---|---|---
|  |  [sec.]
55396.41 |  2010-07-19 |  1200
55397.42 |  2010-07-20 |  1200
55441.31 |  2010-09-02 |  1200
55443.30 |  2010-09-04 |  1200
55444.30 |  2010-09-05 |  1200
55445.28 |  2010-09-06 |  1200
55446.29 |  2010-09-07 |  1200
55451.27 |  2010-09-12 |  1200
55452.26 |  2010-09-13 |  1200
55453.26 |  2010-09-14 |  1200
55467.23 |  2010-09-28 |  1800
55470.20 |  2010-10-01 |  1200
55471.21 |  2010-10-02 |  1200
55472.22 |  2010-10-03 |  1200
55477.20 |  2010-10-08 |  1200
55480.20 |  2010-10-11 |  1800
55481.20 |  2010-10-12 |  951
55483.20 |  2010-10-14 |  1300
55494.16 |  2010-10-25 |  1200
55496.15 |  2010-10-27 |  1200
55497.16 |  2010-10-28 |  1200
55498.17 |  2010-10-29 |  1200
55499.16 |  2010-10-30 |  1200
55501.14 |  2010-11-01 |  1200
55505.12 |  2010-11-05 |  1200
55506.12 |  2010-11-06 |  1200
55507.11 |  2010-11-07 |  1200
55508.11 |  2010-11-08 |  1200
55514.12 |  2010-11-14 |  1200
55515.10 |  2010-11-15 |  1200
55519.10 |  2010-11-19 |  1200

###  2.2 Photometric observations at the Wise Observatory

Mrk 926 was observed with a Bessell V-band filter (Bessell, 1990) at the C18
telescope (Brosch et al., 2008) of the Wise Observatory in Israel, from 2010
August 28 to December 10. Observations were carried out daily with gaps due to
bad weather on some nights. The SBIG STL-6303 CCD was used with
3072$\times$2048 9 $\mu$m pixel, and plate-scale of 1.47 arcsec pix-1, which
gives a field of view of 75$\times$50 arcmin. Each night several exposures of
5 min each were obtained. The V-band filter has a bandwidth ranging from 4700
Å to 6500 Å and an effective central wavelength of 5415Å. The contribution of
the total H$\beta$ flux to the V band can be estimated to be $<10$%. Since the
variable part of H$\beta$ only accounts for 30 – 40 % of the total H$\beta$
flux during our campaign (see Table 5) the variable H$\beta$ contribution to
the V band is $<4$% and is therefore negligible.

Data reduction was carried out in the standard way for bias, dark, and flat-
field, using standard IRAF routines. In order to measure the instrumental
magnitude of the objects in each image we used the DAOPHOT package (Stetson,
1987) as implemented in IRAF. We derived the point-spread function for each
image and measured all available objects in the image. We then used the
DAOSTAT program (Netzer et al., 1996) to choose the non-variable stars in the
field and use them to inter-calibrate by differential photometry the
instrumental magnitude of Mrk 926 between the images and to get its light
curve. Measurements from the same night were averaged to obtain a single
measurement for each night in order to decrease the uncertainty on each point
in the light curve.

Figure 2: Reduced and [O iii] $\lambda$5007-intercalibrated optical spectra of
Mrk 926 taken with the HET telescope during our variability campaign in 2010.
The continuum bands and the broad-line integration intervals are marked
according to Table 2. The spectra shown here are not corrected for Galactic
foreground extinction.

Figure 3: Combined mean spectrum (black) and rms spectrum (red; see Sect. 3.1
for details) of Mrk 926 from our variability campaign. The rms spectrum was
scaled to allow for a direct comparison. The central wavelengths of the
pseudo-continua used for linear continuum subtraction are denoted by the
vertical dashed lines. The line integration limits are marked by the shaded
areas.

##  3 Results

###  3.1 Continuum and spectral line variations

We present all the reduced [O iii] $\lambda$5007-intercalibrated optical
spectra of our variability campaign from Mrk 926 in Fig. 3. They clearly show
variations in the continuum intensities. The mean and rms spectra of Mrk 926
are shown in Fig. 3.555The presented mean and rms spectra are a combination of
the blue and red intercalibrations. We merged them at 6345 Å and used a small
scaling factor to adapt the redward calibration to the intercalibration with
respect to [O iii] $\lambda$5007. The rms spectrum indicates the variable part
of the emission lines. It was scaled to allow for a better comparison with the
mean spectrum and to enhance weaker line structures. The continuum regions and
integration limits of the broad emission lines used in our present study are
given in Table 2 and shown at the bottom of the spectra in Fig. 3. We selected
the continuum regions by inspecting the mean and rms spectra for regions that
are free of both strong emission and absorption lines.

A continuum region at 5100 Å (rest frame) is often used in studies of the
variable continuum flux in AGN. Normally, this region is free of strong
emission lines and close to the [O iii] $\lambda$5007 flux calibration line.
However, in our case this region is contaminated by H$\beta$ because of the
very broad line profile widths in Mrk 926. Thus, we settled for a continuum
range slightly more shifted to the red part of the spectrum (5180 Å in the
rest frame). Finally, we determined the continuum intensities at seven
wavelength ranges (at 4425, 4755, 5420, 5485, 5755, 6530 and 7255 Å in the
observed frame). We integrated the broad emission-line intensities of the
Balmer and He i lines between the wavelength boundaries given in Table 2. The
red wing of the He ii $\lambda 4686$ line and the blue wing of H$\beta$ are
superimposed because of their very broad line widths. There is no clear
minimum between the two lines. Therefore, we assigned the emission-line flux
shortward of 4960 Å to a mix of the He ii $\lambda 4686$ line and blue
H$\beta$ wing, and the flux longward of 4960 Å to H$\beta$. Before integrating
the emission-line flux, we subtracted a linear pseudo-continuum defined by the
continua given in Table 2 (Col. 3). The results of the continuum and the
integrated line intensity measurements – including the narrow-line components
– are given in Tables 9 and 10. The light curves of the integrated emission-
line fluxes of H$\alpha$, H$\beta$, H$\gamma$, H$\beta$${}_{\textrm{blue}}$+He
ii $\lambda 4686$, and He i $\lambda 5876_{\textrm{center}}$ are given in Fig.
6. For further analysis, we also measured the narrow-line components of each
broad line in the mean spectrum. The de-blending process of the narrow and
broad components in the Balmer and He i $\lambda$ 5876 lines is described in
more detail in Sect. 3.4. The resulting narrow-line fluxes are presented in
Table 3.

Table 2: Continuum boundaries, line integration limits, and pseudo-continua.

Cont./Line | Wavelength Range | Pseudo-Continuum
---|---|---
(1) | (2) | (3)
Cont. 4425 (4225) | 4420 – 4430 |
Cont. 4755 (4540) | 4742 – 4770 |
Cont. 5420 (5180) | 5415 – 5430 |
Cont. 5485 (5240) | 5460 – 5510 |
Cont. 5755 (5500) | 5735 – 5775 |
Cont. 6530 (6240) | 6525 – 6540 |
Cont. 7255 (6930) | 7245 – 7265 |
H$\gamma$ | 4430 – 4740 | 4425 – 4755
H$\beta$ | 4810 – 5410 | 4755 – 5420
H$\beta_{\textrm{blue}}$+He ii $\lambda 4686$ | 4810 – 4960 | 4755 – 5420
H$\beta_{\textrm{center}}$ | 4960 – 5175 | 4755 – 5420
H$\beta_{\textrm{red}}$ | 5175 – 5410 | 4755 – 5420
He i $\lambda 5876_{\textrm{center}}$ | 5980 – 6270 | 5755 – 6530
H$\alpha$ | 6540 – 7245 | 6530 – 7255
H$\alpha_{\textrm{blue}}$ | 6540 – 6740 | 6530 – 7255
H$\alpha_{\textrm{center}}$ | 6740 – 7000 | 6530 – 7255
H$\alpha_{\textrm{red}}$ | 7000 – 7245 | 6530 – 7255
666The continuum wavelength values in brackets correspond to the rest frame of
Mrk 926. All wavelengths are given in Å.

Table 3: Observed narrow-emission-line fluxes as measured in the mean spectrum
shown in Fig. 3.

Line | Flux | Line | Flux
---|---|---|---
H$\gamma_{\rm narrow}$ | 7.4 $\pm$ 1.0 |  $\left[{O\textsc{i}}\right]\,\lambda$6363 | 8.0 $\pm$ 1.5
He ii $\lambda 4686$ | 6.3 $\pm$ 0.9 |  $\left[{N\textsc{ii}}\right]\,\lambda$6548 | 18.8 $\pm$ 3.0
H$\beta_{\textrm{narrow}}$ | 15.7 $\pm$ 1.8 | H$\alpha_{\textrm{narrow}}$ | 42.4 $\pm$ 3.4
$\left[{O\textsc{iii}}\right]\,\lambda$4959 | 102\. $\pm$ 2.1 |  $\left[{N\textsc{ii}}\right]\,\lambda$6584 | 54.0 $\pm$ 3.4
$\left[{O\textsc{iii}}\right]\,\lambda$5007 | 314\. $\pm$ 3.3 |  $\left[{S\textsc{ii}}\right]\,\lambda$6717 | 41.1 $\pm$ 2.1
He i $\lambda 5876$ | 3.1 $\pm$ 0.7 |  $\left[{S\textsc{ii}}\right]\,\lambda$6731 | 37.0 $\pm$ 2.0
$\left[{O\textsc{i}}\right]\,\lambda$6300 | 33.8 $\pm$ 2.7 |  |
777In units of 10-15 erg cm-2 s-1.

####  3.1.1 Intercalibration of HET and Wise light curves

We present the light curves of the continuum flux densities at 4425, 4755,
5420, 5485, 5744, and 6530 Å (observed frame) in Fig. 4. In order to increase
the temporal resolution, we created a combined continuum light curve based on
the HET continuum flux densities at 5420 Å and on the photometric V-band
continuum light curve obtained at the Wise Observatory. We followed the
procedure described in Grier et al. (2012) and Zetzl et al. (2018) by applying
a multiplicative scale factor and an additional flux adjustment component to
put the Wise Observatory instrumental magnitude light curve on the same scale
as the HET continuum light curve in order to correct for different
instrumentations, differences in the aperture sizes and therefore different
host galaxy contributions, and constant contributions from emission lines
within the bandwidth of the V-band filter. The optimal scale factor and flux
adjustment component were determined by finding the set of parameters for
which the cumulative difference of (nearly) simultaneous spectroscopic and
photometric continuum measurements reached its minimum. The resulting combined
light curve is given in Table 12 and shown in Fig. 5. Excluding the first two
data points, the time interval between observations amounts to only 1.1 days.
For further analysis, we corrected this light curve for Galactic foreground
extinction ($A_{\rm V}=0.113$; Schlafly & Finkbeiner 2011) and subtracted the
V-band host galaxy contribution ($1.35\times 10^{-15}$ erg cm-2 s-1 Å-1; see
Sect. 3.2). For comparison, we also show the combined continuum light curve in
Figs. 6, 13, and 14.

Figure 4: Light curves of the continuum flux densities at 4425, 4755, 5420,
5485, 5755, and 6530 Å (observed frame; in units of 10-15 erg cm-2 s-1 Å-1) of
our HET variability campaign from 2010 July 19 to November 19. Figure 5:
Combined spectroscopic (HET continuum flux densities at 5420 Å, or 5180 Å in
rest frame; blue) and photometric (V-band continuum data obtained at the Wise
Observatory; black) continuum light curve of our campaign (see Sect. 3.1 for
details). Figure 6: Light curves of the combined continuum flux densities at
5420 Å (observed frame; in units of 10-15 erg cm-2 s-1 Å-1) as well as of the
integrated emission-line fluxes of H$\alpha$, H$\beta$, H$\gamma$, He ii
$\lambda 4686$ \+ H$\beta$blue, and He i $\lambda 5876$ for our HET
variability campaign from 2010 July 19 to November 19.

###  3.2 Host galaxy contribution to the optical continuum flux

The observed flux of the variable AGN component is contaminated by the flux
contribution of the host galaxy. This contribution is nearly constant for
observations taken with identical aperture, and only minor deviations arise
due to slight variations from the exact pointing as well as different seeing
conditions. However, these effects are taken into account by our error
estimates. The overall host galaxy contribution varies from band to band as
the stellar component of the host galaxy has a different flux distribution
than the AGN component. To estimate the relative contribution of the host
galaxy flux, we used the flux variation gradient (FVG) method (Choloniewski,
1981; Winkler et al., 1992; Haas et al., 2011; Zetzl et al., 2018). This
method allows the varying AGN flux in our aperture to be disentangled from the
host galaxy contribution. We took the continuum flux densities at 4425, 5420,
and 6530 Å (observed frame) as proxies for B-, V-, and R-band fluxes,
respectively, since these continuum ranges are close to the maxima of the B-,
V-, and R-filter curves (B-band filter: 4300 Å, V-band filter: 5400 Å, R-band
filter: 6100 Å). In this way, we also excluded the contribution of emission
lines (see Fig. 3). Figure 7 shows the B versus V and B versus R fluxes (black
solid circles) of Mrk 926 based on the HET spectra (aperture: 2 x 2 arcsec).
The B, V, and R values in this figure have been corrected for Galactic
foreground extinction (AB = 0.149, AV = 0.113, AR = 0.089; Schlafly &
Finkbeiner (2011)) and are given in units of mJy. The converted light curves
before and after Galactic foreground extinction are given in Table 11. The
blue dashed line in Fig. 7 gives the best linear fit to the B versus V and B
versus R fluxes. The red shaded area marks the range of host slopes for nearby
AGN as determined by Sakata et al. (2010). The intersection point between the
AGN and host galaxy slopes allows the contribution of the host galaxy fluxes
in the B, V, and R bands to be determined; the resulting host galaxy
contributions in the individual bands are indicated by the dashed black lines.

We derived a B-band host galaxy contribution of 0.76 mJy (mean of 0.56 mJy and
0.96 mJy). The corresponding values are 1.32 mJy for the V band and 2.82 mJy
for the R band. By subtracting the host galaxy contribution for each
individual band, we isolate the AGN flux in all bands. All flux values
(minimal and maximal host+AGN, host contribution as well as true minimal and
maximal AGN flux) are listed in Table 4. Based on the HET spectra, the host
galaxy contribution corresponds to 12 – 26 %, 20 – 42 %, and 37 – 66 % for the
B, V, and R bands, respectively. Altogether, the spectra of Mrk 926 are
dominated by emission from the nucleus. The weak host contribution in the B
and V bands and, especially, the very strong and broad emission lines prevent
us from subtracting a reliable host galaxy template since the stellar
signature of the host galaxy is suppressed.

Table 4: Extinction-corrected B, V, and R values for the combined host galaxy plus AGN fluxes as well as for the host galaxy and AGN fluxes alone as determined by the FVG method. When a range is given, it corresponds to the minimum and maximum flux. Flux Component | B band | V band | R band
---|---|---|---
| [mJy]
(1) | (2) | (3) | (4)
Host+AGN (BvsR) | 2.94 – 6.54 |  | 4.26 – 7.60
Host+AGN (BvsV) | 2.94 – 6.54 | 3.15 – 6.51 |
Host (BvsR) | 0.96 |  | 2.82
Host (BvsV) | 0.56 | 1.32 |
AGN (BvsR) | 1.98 – 5.58 |  | 1.44 – 4.78
AGN (BvsV) | 2.39 – 5.98 | 1.83 – 5.19 |
| [10-15 erg s-1 cm-2 Å-1]
(1) | (2) | (3) | (4)
Host+AGN (BvsR) | 4.51 – 10.01 |  | 3.00 – 5.34
Host+AGN (BvsV) | 4.51 – 10.01 | 3.22 – 6.65 |
Host (BvsR) | 1.47 |  | 1.98
Host (BvsV) | 0.86 | 1.35 |
AGN (BvsR) | 3.03 – 8.54 |  | 1.01 – 3.36
AGN (BvsV) | 3.66 – 9.16 | 1.87 – 5.30 |

Figure 7: Extinction-corrected B versus V (upper panel) and B versus R (lower
panel) flux variations of Mrk 926. The dashed blue lines (AGN slope) represent
the best linear fit to the B versus V and B versus R fluxes, respectively. The
red shaded area shows the range of host slopes. The dashed black lines
indicate the B and V as well as the B and R values of the host galaxy.

###  3.3 Variability statistics

The following statistics are based on the spectroscopic and photometric
observations from HET and the Wise Observatory, respectively. The continua
were corrected for Galactic foreground extinction (see Sect. 3.1), and the
lines were corrected for contribution from the narrow lines (see Table 3). In
Table 5, we present the minimum and maximum fluxes, $F_{\text{min}}$ and
$F_{\text{max}}$, peak-to-peak amplitudes, $R_{\text{max}}$ =
$F_{\text{max}}$/$F_{\text{min}}$, the mean flux over the period of
observations, $<F>$, the standard deviation, $\sigma_{F}$, and the fractional
variation,

$F_{\rm var}=\sqrt{{\sigma_{F}}^{2}-\Delta^{2}}/<F>,$ (1)

as defined by Rodríguez-Pascual et al. (1997). The quantity $\Delta^{2}$ is
the mean square value of the uncertainties $\Delta_{\text{i}}$ associated with
the fluxes $F_{\text{i}}$. The $F_{\rm var}$ uncertainties are defined in
Edelson et al. (2002). The peak-to-peak amplitudes and the fractional
variations in the continuum decrease as a function of wavelength.

Table 5: Variability statistics ($F_{\text{min}}$, $F_{\text{max}}$,
$R_{\text{max}}$, $<F>$, $\sigma_{\text{F}}$, and F${}_{\text{var}}$) for the
continuum (after correction for Galactic foreground extinction) and for the
broad emission lines (after correction for the contribution of narrow-line
components; see Table 3).

Cont./Line | $F_{\text{min}}$ | $F_{\text{max}}$ | $R_{\text{max}}$ | $<F>$ | $\sigma_{\text{F}}$ | $F_{\text{var}}$
---|---|---|---|---|---|---
Comb. V band | 3.22 | 6.65 | 2.07 | 4.41 | 0.97 | 0.21
Cont. 4425 (4225) | 4.51 | 10.01 | 2.22 | 6.22 | 1.77 | 0.28
Cont. 4755 (4540) | 4.04 | 8.87 | 2.20 | 5.67 | 1.52 | 0.27
Cont. 5420 (5180) | 3.22 | 6.65 | 2.07 | 4.50 | 1.01 | 0.22
Cont. 5485 (5240) | 3.28 | 6.59 | 2.01 | 4.54 | 0.99 | 0.22
Cont. 5755 (5500) | 3.02 | 6.04 | 2.00 | 4.23 | 0.90 | 0.21
Cont. 6530 (6240) | 3.00 | 5.34 | 1.78 | 3.85 | 0.68 | 0.17
Cont. 7255 (6930) | 3.20 | 5.03 | 1.57 | 3.93 | 0.54 | 0.14
H$\gamma$ | 122.2 | 184.7 | 1.51 | 154.6 | 17.6 | 0.11
H$\beta$ | 542. | 754. | 1.39 | 670. | 46. | 0.07
H$\beta_{\rm blue}$+He ii $\lambda$4686 | 33.2 | 76.80 | 2.31 | 57.8 | 9.6 | 0.17
H$\beta_{\rm center}$ | 419. | 556. | 1.33 | 502. | 30. | 0.06
H$\beta_{\rm red}$ | 89.2 | 126.8 | 1.42 | 109.5 | 8.5 | 0.08
HeI $\lambda 5876_{\rm center}$ | 63.7 | 129.1 | 2.03 | 92.7 | 13.2 | 0.14
H$\alpha$ | 2293. | 2813. | 1.23 | 2562. | 118. | 0.04
H$\alpha_{\rm blue}$ | 131.3 | 192.4 | 1.47 | 159.2 | 14.7 | 0.09
H$\alpha_{\rm center}$ | 2082. | 2456. | 1.18 | 2286. | 93.3 | 0.04
H$\alpha_{\rm red}$ | 79.6 | 166.9 | 2.10 | 116.1 | 23.9 | 0.21
888In units of 10-15 erg s-1 cm-2 Å-1 for the continuum and in units of 10-15
erg s-1 cm-2 for the broad emission lines. The fluxes in this table are
observed-frame fluxes. The continuum wavelength values in brackets correspond
to the rest frame of Mrk 926.

###  3.4 Mean and rms line profiles

Figure 8: Normalized mean and rms profiles of the Balmer lines and HeI
$\lambda 5876$ in velocity space. The dashed part of the H$\gamma$ line
profile denotes the part that is heavily affected by uncertainties in the flux
calibration due to boundary effects at the blue end of the spectrum.

We determined normalized mean and rms profiles of the Balmer and He I lines in
Mrk 926 after subtracting a linear pseudo-continuum from each mean and rms
broad-line profile. In order to obtain the true mean broad-line profiles, we
subtracted the narrow Balmer and He I line components in velocity-space by
means of a scaled [O iii] $\lambda$5007 template based on the observed [O iii]
$\lambda$5007 line profile in our mean spectrum. Furthermore, we subtracted
the $\left[\text{NII}\right]$6548/6584 emission lines that are superimposed on
the mean H$\alpha$ profile using the same template and sticking to a line-
ratio of 1:3. The rms profiles illustrate the line profile variations during
our campaign. Therefore, the constant narrow components disappear in these rms
profiles. Finally, we normalized both the mean and rms profiles such that the
maximum line flux was set to unity. The final normalized mean and rms profiles
of the Balmer lines H$\alpha$, H$\beta$, H$\gamma$, as well as of the He i
$\lambda 5876$ line in velocity space are shown in Fig. 8. The subtraction of
all narrow Balmer and He I components as well as of
$\left[\text{NII}\right]$6548/6584 and the normalization of the true broad-
line profiles allows the mean and rms profiles to be compared with each other
in a more accurate way.

The Balmer lines exhibit – besides to their central component – an additional
inner red component at 1200 $\pm$ 300 km s-1 in the mean as well as rms
profiles (see Fig. 8). Furthermore, the Balmer lines show additional outer
broad-line components in their rms profiles (Fig. 8) at $\pm 5000$ to $\pm
13\,000$ km s-1. These Balmer satellites are clearly separated from the
central component ($\pm$ 5000 km s-1) and exhibit an amplitude that amounts to
25 – 40% of that of the maximum rms amplitude in H$\alpha$ and H$\beta$,
respectively. These components are not recognizable in the mean profiles. The
H$\gamma$ profile is contaminated by the [OIII] $\lambda$4363 line and a
strong residual is present in the rms profile. This is due to H$\gamma$ being
at the blue end of the spectrum where the flux calibration for the first few
pixels is not optimal. Furthermore, it is not possible to fix a line-free
continuum region on the blue side. These boundary effects heavily affect the
blue part of the H$\gamma$ line profile. Therefore, the H$\gamma$ profile can
only be used to derive some general trends. Despite the boundary effects
mentioned before, the rms profile of H$\gamma$ confirms the presence of Balmer
satellites at $\pm 5000$ to $\pm 13\,000$ km s-1. The red wing of He i
$\lambda 5876$ might also indicate the presence of such a line satellite,
however, a clear detection cannot be stated due to the interference by
absorption in the line wings. In addition, the central He i $\lambda 5876$
profile is contaminated by NaD absorption.

We determined the line widths (FWHM) of the mean and rms line profiles of all
the Balmer lines and He i $\lambda 5876$. Additionally, we parameterized the
line widths of the mean profiles by their line dispersion
$\sigma_{\text{line}}$ (Fromerth & Melia, 2000; Peterson et al., 2004).
However, due to the very broad, partially overlapping and complex rms line
profiles with additional outer components, it was not possible do determine
reliable $\sigma_{\text{line}}$ (rms) values. Likewise, by cause of H$\gamma$
being affected by boundary effects, the overlapping He ii $\lambda$4686 and
H$\beta_{\rm blue}$ profiles, as well as strong absorption in He i $\lambda
5876$, a determination of $\sigma_{\text{line}}$ (mean) for these lines was
deemed unreliable. Therefore, we determined the line dispersion
$\sigma_{\text{line}}$ (mean) solely for the H$\alpha$ mean profile over a
velocity interval from $-15\,000$ to $+15\,000$ km s-1.

The resulting widths (FWHM) for all emission lines, and the line dispersion of
H$\alpha$ ($\sim 3500$ km s-1) are given in Table 6. The H$\alpha$ and
H$\beta$ lines show similar mean profiles with line widths (FWHM) of $\sim
5200$ km s-1as well as similar rms profiles with line widths (FWHM) of $\sim
4000$ km s-1. The H$\gamma$ line exhibits a profile similar to that of
H$\alpha$ and H$\beta$, but the line width and line dispersion are smaller
than for the other Balmer lines. We note, however, that these results are less
reliable due to the aforementioned boundary effects. The He i $\lambda 5876$
line (mean and rms) is by a factor of about 1.5 broader than the Balmer lines.

Table 6: FWHM of the mean and rms line profiles as well as the line dispersion, $\sigma_{\text{line}}$, of the mean profiles. Line | FWHM (mean) | FWHM (rms) |  $\sigma_{\text{line}}$ (mean)
---|---|---|---
| [km s-1] | [km s-1] | [km s-1]
H$\alpha$ | 5370 $\pm$ 400 | 3890 $\pm$ 500 | 3460 $\pm$ 1000
H$\beta$ | 5110 $\pm$ 400 | 4230 $\pm$ 400 | –
H$\gamma$ | 4300 $\pm$ 1000 | 3260 $\pm$ 800 | –
He i $\lambda 5876$ | 8500 $\pm$ 1000 | 7650 $\pm$ 800 | –

###  3.5 CCF analysis of the broad emission lines

The mean distances of the broad emission-line regions to the central ionizing
source can be determined by correlating the broad emission-line light curves
with the light curve of the ionizing continuum flux. Normally, an optical
continuum light curve is used as surrogate for the ionizing light curve. For
this study, we correlated the integrated light curves as well as the segment
light curves (center, blue, red) of the Balmer lines H$\alpha$ and H$\beta$,
and the He I $\lambda 5876_{\rm center}$ line with the combined 5180 Å (rest
frame) and V-band continuum light curve (see Fig. 5 in Sect. 3.1). For
H$\gamma$, the boundary effects mentioned in Sect. 3.4 prevented a reliable
light curve extraction, and therefore no cross-correlation function (CCF) was
calculated. The correlation technique we used is a variant of the discrete
correlation function (DCF; Edelson & Krolik, 1988). Instead of classically
binning the data points (i.e., using a rectangular weighting function), we
performed a weighted averaging of the data points using a Gaussian kernel
$b_{\tau}(t)=(1/\sqrt{2\pi h})\exp(-(\tau-t)^{2}/2h^{2})$ as a smooth density
function (Rehfeld et al., 2011, and references therein). The width $h$ was
chosen such that it was on the order of the mean sampling rate, which in our
case translates to $\sim 2$ days. The Pearson correlation coefficient,
$\rho_{x,y}=\text{COV}(x,y)/(\sigma_{x}\sigma_{y}),$ of two time series, $x$
and $y$, with time steps $t^{x}$ and $t^{y}$ then transforms to

$\text{CCF}(\tau)=\frac{\sum^{N_{x}}_{i=1}(x_{i}-\mu_{x})\cdot(z^{\tau}_{i}-\mu_{z^{\tau}})}{N_{x}\sigma_{x}\sigma_{z}},$
(2)

with $z^{\tau}_{i}=\sum^{N_{y}}_{j=1}y_{j}\cdot
b_{\tau}(t^{y}_{j}-t^{x}_{i})$. The normalization of the CCF ensures that
CCF$(\tau)\leq 1$ for all $\tau$. We note that the light curves were not
detrended as the timescale of the campaign does not suffice to remove long-
term trends. The derived cross-correlation functions, CCF$\left(\tau\right),$
are shown in Fig. 9.

(a)

(b)

(c)

(d)

(e)

(f)

(g)

(h)

Figure 9: Left panel: CCFs of the integrated H$\alpha$ (a) and H$\beta$ (c)
line, of their segments (center, blue, red), and of the central He I line (e)
with respect to the combined continuum at 5180 Å (rest frame). The time lag,
$\tau_{\rm cent}$, is denoted by a dashed line, with the shaded area
corresponding to a $\pm 1\sigma$ interval. The CCF of the spectroscopic
continuum light curve with respect to the combined (photometric plus
spectroscopic) continuum reference light curve is shown in (f). Right panel:
CCCDs of the integrated H$\alpha$ (b) and H$\beta$ (d) line, of their segments
(center, blue, red), and of the central He I line (g) with respect to the
combined continuum at 5180 Å (rest frame). The boundary between CCCDs of
individual peaks is denoted by a dashed gray line. The CCCD of the
spectroscopic continuum light curve with respect to the combined continuum
reference light curve is shown in (h).

We determined the centroids $\tau_{\text{cent}}$ of the CCFs by using only
those parts of the CCFs above 80% of the peak value. A threshold value of 0.8
$r_{\text{max}}$ is generally a good choice as has been shown before (Peterson
et al., 2004). We derive the uncertainties of the time lags,
$\tau_{\text{cent}}$, by calculating the cross-correlation lags a large number
of times using a model-independent Monte Carlo method known as flux
randomization/random subset selection (FR/RSS). This method has been described
by Peterson et al. (1998). The resulting cross-correlation centroid
distributions (CCCDs) for the integrated H$\alpha$, H$\beta$, and He I lines
as well as for their line segments (center, blue, red) with the combined 5180
Å (rest frame) and V-band continuum are presented in Fig. 9. Each CCCD was
determined after $2\times 10^{4}$ independent runs (i.e., $2\times 10^{4}$
independent subsamples) for each light curve. For the CCFs of the central
Balmer line segments showing two distinct peaks, the CCCD was calculated for
each peak individually. The final centroid time lags are given in Table 7. The
error intervals correspond to 68% ($\pm 1\sigma$) confidence level.

The resulting CCFs and CCCDs are presented in Fig. 9. The CCFs of the
integrated Balmer line light curves and of the central Balmer line component
light curves (within $\pm 5000$ km s-1) exhibit two discrete peaks. In
contrast, the blue and red Balmer segment light curves (i.e., the light curves
of the Balmer satellites) are single-peaked. The double-peaked CCFs of the
integrated Balmer line light curves show delays of $8^{+2}_{-5}$ and
$56^{+3}_{-6}$ days for H$\alpha$, and $5^{+2}_{-2}$ and $48^{+9}_{-6}$ days
for H$\beta$. The CCFs of the central Balmer line component light curves show
delays of $10^{+3}_{-2}$ and $57^{+3}_{-3}$ days for H$\alpha$, and
$5^{+2}_{-2}$ and $48^{+7}_{-6}$ days for H$\beta$. The interpretation of
these double-peaked CCFs is explained in more detail in Sect. 3.6. The CCFs of
the Balmer satellite light curves are single-peaked and show, in comparison to
the CCFs of the integrated and central component light curves, much shorter
delays of only $3-5$ light-days (i.e., the Balmer satellites react promptly to
the variations in the continuum). There are indications for a secondary peak
in the red Balmer satellite CCF of H$\beta$ (see Fig. 9(c)); however, the
light curve of this wing is relatively noisy due to the strong underlying [O
III] lines (see Sect. 4.5.2 for more information). The CCF of the blue
H$\beta$ Balmer satellite shows no indication of a secondary peak.

We also calculated the CCFs and CCCDs for the weak HeI $\lambda 5876$ line.
The outer wings are heavily contaminated by absorption. Therefore, out of
precaution we excluded the blue and red wing from the integrated line light
curve, leaving only the central profile light curve. In contrast to the Balmer
lines, the CCF of the central profile light curve only shows a singular peak
at a delay of $7^{+4}_{-2}$ days and no evidence for a secondary peak.
Instead, the singular peak is very broad and high correlation values are found
for delays up to $\sim 40$ days.

The overall lags determined by the CCF analysis do not depend significantly on
the use of either the purely spectroscopic or the combined spectroscopic and
photometric continuum light curve. Both driving light curves recover the same
responses of the emission lines with only minor variations of $\pm 2$ days at
maximum. In comparison to the spectroscopic driving light curve, the use of
the combined continuum light curve generates more stable CCFs and narrower,
more confined CCCDs.

Table 7: Cross-correlation lags of the integrated H$\alpha$ and H$\beta$ line light curves as well as of their line segments (center, blue, red) with respect to the combined 5180 Å (rest frame) and V-band continuum light curve. | Line
---|---
| H$\alpha$ | H$\beta$ | He I
Segment |  $\tau$ [days] |  $\tau$ [days] |  $\tau$ [days]
inner | outer | inner | outer | inner | outer
total | $8^{+2}_{-5}$ | $56^{+3}_{-6}$ | $5^{+2}_{-2}$ | $48^{+9}_{-5}$ | - | -
center | $10^{+3}_{-2}$ | $57^{+3}_{-3}$ | $5^{+2}_{-2}$ | $48^{+7}_{-6}$ | $7^{+4}_{-2}$ | -
blue | $4^{+1}_{-3}$ | - | $5^{+1}_{-3}$ | - | - | -
red | $3^{+2}_{-2}$ | - | $6^{+2}_{-5}$ | - | - | -

###  3.6 Velocity-resolved CCFs of the H$\alpha$ and H$\beta$ line

In Sect. 3.5 we calculated the time lags of the integrated H$\alpha$,
H$\beta$, and HeI lines as well as of the central and outer wing regions with
respect to the combined V-band continuum at 5180 Å (rest frame). Now, we
investigate the profile variations in the H$\alpha$ and H$\beta$ lines in more
detail by calculating the lags of individual line segments. For the weak HeI
$\lambda 5876$ line, the absorption in the central line segment and in the
line wings prohibits a clear determination of a velocity-resolved CCF. The way
we proceed has been described before in our studies of line profile variations
in Mrk 110 (Kollatschny & Bischoff, 2002; Kollatschny, 2003), Mrk 926
(Kollatschny & Zetzl, 2010), 3C120 (Kollatschny et al., 2014) and HE 1136-2304
(Kollatschny et al., 2018).

We sliced the velocity profiles of the continuum-subtracted Balmer lines into
velocity segments with a width of $\Delta\leavevmode\nobreak\ v=400$ km s-1.
This value of 400 km s-1 corresponds to the spectral resolution of our
observations. A central line-segment was integrated from $v=-200$ to $+200$ km
s-1. Afterward, we measured the intensities of all subsequent line-of-sight
velocity segments from $v=-15\,800$ to $+15\,800$ km s-1 and compiled their
light curves. We computed the CCF$\left(\tau\right)$ for each individual line-
segment light curve ($\Delta{}v=400$ km s-1) of H$\alpha$ and H$\beta$ with
respect to the combined 5180 Å (rest frame) continuum light curve. In this
way, we computed velocity-resolved CCFs of H$\alpha$ and H$\beta$ as a
function of distance to the line center (blue scale). These velocity-resolved
CCFs are shown in Figs. 10(a) and 10(b). The green lines in these figures
delineate the contour lines of the correlation coefficient at different levels
(0.2, 0.4, 0.5, 0.6, 0.7, 0.75, 0.8, 0.825). The red line shows the rms line
profile for comparison.

(a)

(b)

Figure 10: Velocity-resolved CCFs($\tau$,$v$) showing the correlation
coefficient of the H$\alpha$ (a) and H$\beta$ (b) line segment light curves
($\Delta v\sim 400$ km s-1) with respect to the combined continuum light curve
as a function of velocity and time delay (blue scale). Contours of the
correlation coefficients are plotted at levels 0.2, 0.4, 0.5, 0.6, 0.7, 0.75,
0.8, and 0.825 (green lines). In each plot, the rms profile of the emission
line is shown in red. The dashed curve shows the escape velocities for a
Keplerian disk inclined by $\sim 50^{\circ}$ with a central mass of 1.1
$\times 10^{8}M_{\odot}$. The dashed ellipse corresponds to a circular
Keplerian orbit at $R/c=33.5$ d inclined by $\sim 50^{\circ}$ to the line of
sight mapped onto the velocity-delay plane. The line center ($v=0$ km s-1) and
the time delay of $\tau=0$ days are indicated by a vertical and a horizontal
black line, respectively. The delays computed for the red wing of H$\beta$ are
disturbed by the [O III] $\lambda$ 4959, 5007 lines, which results in higher-
delay residuals not confined within the viral envelope.

The H$\alpha$ line shows two velocity-delay structures in the central segment
(within $\pm 5000$ km s-1) at a distance of $57^{+3}_{-3}$ and $10^{+3}_{-2}$
light-days (see Sect. 3.5). Adopting the interpretation of Horne et al. (2021)
for the velocity-delay maps of the NGC 5548 STORM campaign, this structure
might be interpreted as the upper and lower half of an ellipse in the
velocity-delay plane that might be the signature of a line-emitting ring
orbiting the black hole (BH) at a radius $R=33.5$ light-days. The stronger far
side of the annulus extends to $\tau=(R/c)\,(1+\sin i)\approx 57$ days, and
the weaker near side of the annulus has a delay of $\tau=(R/c)\,(1-\sin i)$
$\approx 10$ days. This gives $\sin i\sim 0.71$ and thus $i\sim 45^{\circ}$ as
the inclination angle of an assumed thin disk.

The H$\beta$ line also shows two velocity-delay structures in the central
segment at $48^{+7}_{-6}$ and $5^{+2}_{-2}$ light-days (see Sect. 3.5). A
possible line-emitting ring orbiting the BH therefore has a radius $R=26.5$
light-days. The weaker far side of the annulus then extends to
$\tau=(R/c)\,(1+\sin i)\approx 48$ days, and the stronger near side of the
annulus has a delay at $\tau=(R/c)\,(1-\sin i)\approx 5$ days. This gives
$\sin i\sim 0.81$ and thus $i\sim 54^{\circ}$ of the disk. Therefore, we adopt
$i\sim 50^{\circ}$ as the mean disk inclination angle.

###  3.7 Central black hole mass

The masses of the central BHs in AGN can be estimated – based on the
assumption that the gas dynamics are dominated by the central massive object –
by evaluating

$M=f\,c\,\tau\,\Delta\,v^{2}\,G^{-1}.$ (3)

The characteristic velocity $\Delta v$ of the emission-line region can be
estimated from the FWHM of the rms profiles or from the line dispersions
$\sigma_{\text{line}}$. Furthermore, the distance $c\,\tau$ of the line-
emitting regions from the ionizing source can be estimated from the delayed
response of the line light curves to the continuum variations. Characteristic
distances of the individual line-emitting regions are given by the centroid
$\tau_{\text{cent}}$ of the individual CCFs of the emission-line variations
relative to the continuum variations (e.g., Koratkar & Gaskell, 1991;
Kollatschny & Dietrich, 1997).

The scaling factor $f$ in Eq. 3 is on the order of unity and depends on the
kinematics, structure, and orientation of the BLR. This scaling factor differs
by $\sim 0.4$ dex rms from galaxy to galaxy, for example, depending on whether
we see the central accretion disk including the BLR from the edge or face-on.
Typically, the scaling factor $f$ is discussed with respect to the line
dispersion $\sigma_{\text{line}}$ (rms). Empirically found $f$ values are
$f=5.5$ (e.g., Onken et al., 2004), $f=4.31$ (e.g., Grier et al., 2013a), or
$f=3.6$ (Graham et al., 2011). The very broad lines in Mrk 926 (see Sect. 3.4)
and the high disk inclination angle of $i\sim 50^{\circ}$ (see Sect. 3.6)
indicate that Mrk 926’s orientation toward us is rather edge-on in comparison
to other AGN. Therefore, we adopt an $f$ value of $f=3.6$ in order to
calculate the BH mass. We note that the virial factor, $f$, for
$\sigma_{\text{line}}$ and FWHM differs since FWHM/$\sigma_{\text{line}}$
typically takes values of $\sim 2$ (Peterson et al., 2004). We account for
that by following the procedure in Kollatschny & Zetzl (2011, 2013) and, for
example, Kollatschny et al. (2018): adopting a correction factor in order to
obtain the true rotational velocity, $v$, from the observed FWHM for each
line. This reduces the effective $f$ factor of FWHM to $\sim 1.2$. We use FWHM
(rms), which generally gives more reliable results than FWHM (mean) due to the
fact that the rms profiles only shows the varying fraction of the BLR line-
emitting gas. Based on the delays of the central Balmer line regions (see
Table 7 and Sect. 3.6) and on the line widths of the rms profiles (FWHM) (see
Table 6), we derive a weighted mean BH mass (see Table 8) of

$M=(1.1\pm 0.2)\times 10^{8}M_{\odot}\ \ \ \text{or}\ \ \
\log(M/M_{\odot})=8.04\pm 0.1.$ (4)

The BH mass based on the HeI $\lambda 5876$ line has a larger error due to
larger uncertainties in their FWHM. However, it confirms the BH mass derived
from the Balmer lines. The BH mass estimate based on the H$\alpha$ line
dispersion $\sigma_{\text{line}}$ (mean) is by a factor of $\sim 2$ larger
than the BH estimate from the FWHM (rms). Considering the larger uncertainties
of $\sigma_{\text{line}}$, it is still in agreement with the BH mass estimates
from FWHM (rms).

Table 8:  Black hole masses based on the FWHM of the line profiles in the rms spectra and the line dispersion $\sigma_{\text{line}}$ of the mean H$\alpha$ profile. The masses were calculated assuming an $f$ value of 3.6. Line |  $M_{\text{BH, FWHM}}$ |  $M_{\text{BH},\sigma_{\text{line}}}$
---|---|---
|  [$10^{8}M_{\odot}$] |  [$10^{8}M_{\odot}$]
H$\alpha$ |  1.2 $\pm{}$ 0.4 |  2.8 $\pm{}$ 1.7
H$\beta$ |  1.1 $\pm{}$ 0.3 |  –
weigh. mean |  1.1 $\pm{}$ 0.2 |  –
HeI $\lambda 5876$ |  1.0 $\pm{}$ 0.6 |  –

##  4 Discussion

###  4.1 Optical variability

####  4.1.1 Optical variability amplitudes

It is known that Mrk 926 varied at least since the beginning of the 1990s
(Kollatschny et al., 2006). A first variability campaign was carried out
during the years 2004-2005 (Kollatschny & Zetzl, 2010). Back then, the optical
continuum flux density at 5180 Å varied between 2.4 and 3.6 $\times 10^{-15}$
erg cm-2 s-1 Å-1. In July 2010, the continuum flux density was significantly
higher and amounted to a value of 6 $\times 10^{-15}$ erg cm-2 s-1 Å-1 at 5180
Å (rest frame) at the beginning of the variability campaign presented here.
Immediately afterward, the optical continuum started to drop, finally reaching
less than 50% of its original intensity (see Fig. 5 and Table 9) within only
2.5 months. Even more remarkably, the true non-stellar continuum in Mrk 926
declines to only 30 – 40% continuum intensity between the beginning and end of
the variability campaign taking into account the constant host galaxy
contribution to the flux as determined by means of the FVG method (see Table
4).

The variability amplitude in the blue spectral range is higher than in the red
(see Table 5 and the rms spectrum in Fig. 3). Interestingly, the continuum
variation amplitudes in Mrk 926 are much higher than those observed in other
highly variable and even changing-look AGN. For example, during a post-
outburst campaign in 2014-2015, the changing-look AGN HE 1136-2304 exhibited a
fractional variation of $F_{\text{V,var}}$ = 0.11 and $F_{\text{V,var}}$ =
0.25 with and without host, respectively (Zetzl et al., 2018). Mrk 926
exhibits fractional variation of $F_{\text{V,var}}$ = 0.22 and
$F_{\text{V,var}}$ = 0.31 with and without host, respectively. The fractional
variation depends on the duration of the monitoring campaign, on the examined
wavelength, and on the (accurate) subtraction of the host galaxy contribution.
A typical value for the fractional variation $F_{\text{V,var}}$ of the
continuum at roughly 5100 Å is 0.05 to 0.15 for variability periods of 6 to 12
months. For longer campaigns, typical $F_{\text{V,var}}$ values are higher and
range from 0.1 to 0.25 (Zetzl et al., 2018, and references therein). On
timescales of years, a decline of the continuum to only 20% of its original
value has been observed for some sources (e.g., Fairall 9; Kollatschny &
Fricke, 1985). However, Mrk 926 was observed for a much shorter period and
showed a similarly drastic decline. To summarize, it can be stated that the
optical continuum variations in Mrk 926 are very strong.

####  4.1.2 Signature of an accretion disk in Mrk 926

We now intend to test whether the observed continuum spectral index in the
optical is consistent with theoretical predictions from the Shakura & Sunyaev
accretion disk model (SS73; Shakura & Sunyaev, 1973) . Specifically, the
Shakura & Sunyaev model predicts $f_{\nu}\propto\nu^{1/3}$ over the self-
similar part of the spectrum, that is, far away from the spectral bands
affected by the inner (UV to extreme UV) and outer (far-infrared) disk radii.
This holds for Mrk 926 over the optical range. The spectral index of
$\alpha={1/3}$ for $f_{\nu}$ then translates to $\beta={-7/3}$ for
$f_{\lambda}$ as $f_{\lambda}=f_{\nu}d\nu/d\lambda\sim
v^{1/3}v^{2}\sim\lambda^{-7/3}$.

Figure 11 shows the observed rms spectrum corrected for Galactic foreground
extinction. A power-law model $F_{\lambda}\propto\lambda^{-\beta}$ was fitted
to the continuum (blue shaded areas), giving a spectral index $\beta=2.33\pm
0.02$. The rms spectrum of the continuum is thus consistent with a power- law
with a spectral index of $\beta=7/3$ as predicted by the Shakura & Sunyaev
model.

Figure 11: RMS spectrum of the HET campaign corrected for Galactic foreground
extinction. A power-law model was fitted to the continuum, giving a spectral
index $\beta=2.33\pm 0.02$.

####  4.1.3 Balmer decrement variability

The Balmer decrement $F({\rm H}\alpha_{\rm narrow})/F({\rm H}\beta{}_{\rm
narrow})$ of the narrow-line components in Mrk 926 amounts to 2.70 (see Table
3). This corresponds to the expected theoretical line ratio (Case B) without
any reddening. The lack of reddening toward the narrow-line region (NLR) is
consistent with the insignificant reddening implied by the RMS fit to a
Shakura & Sunyaev model (see Sect. 4.1.2). In contrast to the constant Balmer
decrement of the narrow lines, the Balmer decrement of the broad-line
components takes values of 3.4 to 4.2. Figure 12 shows the Balmer decrement
$F({\rm H}\alpha_{\rm broad})/F({\rm H}\beta_{\rm broad})$ versus broad-line
H$\beta$ flux.

Figure 12: Balmer decrement, $F({\rm H}\alpha)/F({\rm H}\beta)$, of the broad-
line components versus broad-line H$\beta$ intensity. The dashed line on the
graph represents the linear regression.

The Balmer decrement varies as a function of the broad H$\beta$ flux. More
precisely, the Balmer decrement shows a roughly linear dependence on the
H$\beta$ flux, and increases with decreasing H$\beta$ broad-line intensity.
Similar, although even stronger variations in the Balmer decrement have been
found before in highly variable Seyfert galaxies as, for example, NGC 7603
(Kollatschny et al. 2000) or HE 1136-2304 (Kollatschny et al. 2018). The
broad-line Balmer decrement in these galaxies increased up to values of 7.3 in
combination with decreasing H$\beta$ line intensities. These observations
might be explained by optical depth effects in the BLR. This is in accordance
with the finding that H$\alpha$ often originates at larger distances than
H$\beta$. It has been discussed by Korista & Goad (2004) that the radial
stratification in the BLR is a result of optical-depth effects of the Balmer
lines: the broad-line Balmer decrement decreases in high continuum states and
increases in low states.

###  4.2 H$\beta$ lag versus optical continuum luminosity

Now, we test whether Mrk 926 follows the general trend of the H$\beta$-lag and
optical continuum luminosity relationship
(R${}_{\text{BLR}}$-L${}_{\text{AGN}}$) for AGN (e.g., Kaspi et al., 2000;
Bentz et al., 2013; Kilerci Eser et al., 2015; Grier et al., 2017; Kollatschny
et al., 2018, and references therein). We determined continuum luminosities
$\log(\lambda\,L_{\lambda}/{\rm erg\leavevmode\nobreak\ s}^{-1})$ of 43.68 to
44.13 (1.87 to 5.30 $\times 10^{-15}$ erg s-1 cm-2 Å-1) in the optical at 5180
Å (rest frame) after correction for the contribution of the host galaxy.
Furthermore, we derived a mean distance of 26.5 light-days for the H$\beta$
line-emitting region based on the delay of the H$\beta$ line variability with
respect to the optical continuum light curve. This derived distance of 26.5
light-days for the H$\beta$ line-emitting region and the observed continuum
luminosities closely follows the general$R_{\text{BLR}}$-$L_{\text{AGN}}$
relation with a slope of $\alpha\sim 0.53$ (Bentz et al., 2013).

###  4.3 Line profiles

####  4.3.1 Mean and rms line profiles of the 2010 campaign

The mean and rms profiles of the Balmer and He i $\lambda 5876$ lines in
velocity space are presented in Fig. 8. The mean and rms profiles of H$\alpha$
and H$\beta$ are very similar. They exhibit line widths (FWHM) of $\sim$ 5250
km s-1 (mean profiles) and line widths of $\sim$ 4000 km s-1 (rms profiles)
(see Table 6). These values are larger than the average. The H$\alpha$ FWHM
distribution for broad-line quasars – based on many thousand emission-line
galaxies from the Sloan Digital Sky Survey – shows a clear maximum at about
2000 km s-1 and falls steadily to FWHM values of about 9000 km s-1 (Hao et
al., 2005). There is an additional inner red emission component to be seen at
$\sim$ +1500 km s-1. Such inner red components have been seen before in other
variable AGN, for example Ark 120 (Kollatschny et al., 1981), NGC4593
(Kollatschny & Dietrich, 1997), NGC7603 (Kollatschny et al., 2000), or HE
1136-2304 (Kollatschny et al., 2018). Possible explanations for these red
components might be disk components, additional radial motions, or opacity
components. More comments regarding inner red line components can be found,
for example, in Gaskell & Harrington (2018) or Du et al. (2018).

The Balmer lines show symmetric additional outer-line components between $\pm
5000$ and $\pm 13\,000$ km s-1, which we term Balmer satellites. We note that
the red wing of He i $\lambda 5876$ might also indicate the presence of such a
line satellite; however, due to the interference by absorption, a clear
detection cannot be claimed. In particular, the Balmer satellites components
show up in the rms profiles and their relative variability amplitudes are
stronger than that of the central components between $-5000$ and $+5000$ km
s-1. The fractional variations in the individual line components are given in
Table 5. Five years before the variability campaign presented here, when Mrk
926 was by 50% fainter in the continuum compared to the beginning of our
campaign, these outer components in the rms profiles were also present, but
are now slightly more pronounced (Kollatschny & Zetzl, 2010). Back in 2005,
the central line components were noticeably broader when the source was in a
lower state. The central Balmer line components exhibited FWHM rms of roughly
8500 km s-1 in 2005 in comparison to 4000 km s-1 in 2010. The uncertainties in
the line dispersion $\sigma_{\text{line}}$ of the broad emission lines
determined in Sect. 3.4 are relatively high. Generally, the line dispersion
$\sigma_{\text{line}}$ is very sensitive to the line widths used for
integration. However, in Mrk 926, the uncertainties in the line widths – full
width at zero intensity (FWZI) – are high because of the large line widths and
the additional outer line-components. Therefore, we did not derive a
FWHM/$\sigma$ to determine the height/radius ratio of the BLR as presented in
other galaxies (e.g., Kollatschny et al., 2018). Typical values for
FWHM/$\sigma$ in other galaxies are on the order of two or more. Only in very
rare cases the $\sigma_{\text{line}}$ values are equal or larger than the FWHM
values (Peterson et al., 2004; Kollatschny & Zetzl, 2011, 2013).

####  4.3.2 Comparison of the 2010 and 2004-2005 campaigns

We now compare the Balmer line profiles observed in 2010 with the Balmer line
profiles observed during the campaign in 2004-2005. First, the continuum as
well as broad-line flux in 2005 was by a factor of about two smaller than in
2010 (Kollatschny & Zetzl, 2010). Second, the comparison of the Balmer line
profiles also reveals differences in the mean and rms profiles. A strong blue
component was present in the mean profiles in 2005, but nearly disappeared in
2010. A small residual of this blue component appears to be left in the rms
profiles at around $-1500$ km s-1. The narrow-line subtraction for the
campaign of 2010 revealed a red component in the mean profile of H$\alpha$ at
around $\pm 1200$ km s-1 (see Sect. 3.1) that is also discernible in the
H$\beta$ profile (see above). Moreover, the rms profiles show a variable red
component with the same velocity. This inner red component is in agreement
with a red and slightly stronger component already present in the mean and rms
profiles of the campaign in 2004/05. In summary, both the mean and rms
profiles of the campaigns in 2004/05 and 2010 show blue and red components at
mirrored velocities. These components were stronger in 2004/05. Presumably,
these symmetric (with respect to line-of-sight velocity) components are the
signature of an accretion disk as accretion disks are expected to result in
double-peaked profiles (e.g., Horne & Marsh, 1986; Eracleous & Halpern, 2003;
Gezari et al., 2007; Shapovalova et al., 2013; Storchi-Bergmann et al., 2017,
and references therein).

The highly variable outer rms Balmer satellite components were also present in
2005, but are more pronounced in 2010. Such variable outer line-components are
very rare and have been seen only in the rms profiles of a few broad-line
radio galaxies, for example, 3C390.3, Arp102B, 3C382, or 3C332 (Gezari et al.,
2007). Mrk 926 shows a radio flux of $18\pm 5$ mJy at 5 Ghz (Bicay et al.,
1995). The optical B-band flux is 6.6 mJy (McAlary et al., 1983). Therefore,
the radio-to-optical luminosity R (F(5 Ghz)/F(B-band)) has a value of 3. This
is above the definition for a radio-quiet AGN (R: $0.1-1$) and below the value
for a radio-loud AGN (R: 10 - 1000) (Kellermann et al., 1989). However, the
highly variable outer components in the rms profiles shown by Gezari et al.
(2007) are almost exclusively connected to the flanks of the double-peaked
mean profiles, which are assumed to originate from an accretion disk. This is
not the case for Mrk 926. Instead, the rms components are further out and
their position overlaps with the outer line-wings, but not the line-flanks
(see Fig. 8).

###  4.4 Central black hole mass and Eddington ratio, L/Ledd

In Sect. 3.7, we derived a central BH mass of $1.1\times 10^{8}M_{\odot}$.
This is a normal BH mass in the typical range of $10^{6}-10^{9}M_{\odot}$
(e.g., Woo & Urry, 2002; Peterson et al., 2004). In the following, we derive
the Eddington ratios, $L/L_{\text{edd}}$, for the low and high intensity state
of Mrk 926 during our variability campaign. The optical luminosity of Mrk 926
during the low state in November 2010 was by a factor of 2.8 lower with
respect to the high state in August 2010. We derive bolometric luminosities of
$L_{\text{bol}}(\text{max})=1.2\times 10^{45}$erg s-1 and
$L_{\text{bol}}(\text{min})=4.3\times 10^{44}$erg s-1 for the high and low
state, respectively, based on the monochromatic luminosity L5100 and on an
average bolometric correction factor of $f_{\text{bol}}$ =
$L_{\text{bol}}$/$L_{5100}$ = 10 (Kaspi et al., 2000; Castelló-Mor et al.,
2016). Although the boundary between Seyfert galaxies and the higher
luminosity quasars is not well defined, a bolometric luminosity of
$L_{\text{bol}}=10^{45}$erg s-1 of the central source is often used as a
dividing line to distinguish between these two types (Netzer, 2013). This
places Mrk 926 directly at the dividing line, hence being, depending on the
time of observation, either a very strong Seyfert 1 galaxy or a quasar.

The BH mass of $1.1\times 10^{8}M_{\odot}$ corresponds to an Eddington
luminosity of $L_{\text{edd}}=1.4\times 10^{46}$ erg s-1. Therefore, the
Eddington ratios, $L/L_{\text{edd}}$, for the high and low state have
comparatively low values, 8 and 3 percent only. Such low Eddington ratios of a
few percent are typical for highly variable objects. For example, the
changing-look AGN IRAS 23226-3843 showed a low Eddington ratio of only one
percent (Kollatschny et al., 2020). This low Eddington ratio is consistent
with investigations of Noda & Done (2018) and MacLeod et al. (2019). In
particular, Noda & Done (2018) suggested that all changing-look AGN are
associated with state transitions at Eddington ratios of a few percent.

###  4.5 Structure and kinematics of the BLR

####  4.5.1 CCFs of the central, blue, and red line segments

The CCFs of the integrated Balmer line light curves and of the central Balmer
line component light curves (within $\pm 5000$ km s-1) exhibit two discrete
peaks, one at a delay of $\sim 10$ days (H$\alpha$) and $\sim 5$ days
(H$\beta$), and the outer at a delay of $\sim 56$ days (H$\alpha$) and $\sim
48$ days (H$\beta$). The long-delay peaks of H$\alpha$ and H$\beta$ indicate a
stratification of the BLR as the CCFs peak at $57^{+3}_{-3}$ and
$48^{+7}_{-6}$ days, respectively. The blue and red Balmer segment light
curves (i.e., the light curves of the Balmer satellite components) are single-
peaked, with delays of $\sim 4$ days for H$\alpha$ and $\sim 5$ days for
H$\beta$. This means that the outer wings respond promptly to the variations
in the continuum.

The HeI $\lambda 5876$ line shows no double-peaked CCF. Instead, the CCF of
the central line profile light curve shows a very broad peak with high
correlations up to $\sim 40$ days. This broad peak might be caused by two
merging peaks (i.e., by two independently responding components exactly as
observed for H$\alpha$ and H$\beta$) but with relatively broad transfer
functions that our campaign was not able to resolve.

####  4.5.2 Velocity-resolved CCFs of the Balmer lines

The velocity-resolved CCFs of H$\alpha$ and H$\beta$ (Figs. 10(a) and 10(b))
are, in essence, similar to each other. The delays of the central profiles and
of the outer wings are all confined within the virial envelope for a Keplerian
disk inclined by $\sim 50^{\circ}$ with a central mass of $1.1\times
10^{8}M_{\odot}$. Signatures outside of the virial envelopes can be attributed
to relatively strong narrow-line residuals (e.g., in the red wing of
H$\beta$), to overlapping velocity-resolved line signatures (e.g., of He II
$\lambda$ 4686 and H$\beta$), or to strong absorption disturbing the segment
light curves (e.g., in the outermost blue wing of H$\alpha$). Both velocity-
resolved CCFs show a broad inner delay region with small time lags of $\sim
3-10$ days. The correlation is strongest in the line wings beyond $\pm 5000$
km s-1, exactly where we observed the maximum of the additional, varying
Balmer satellites in the rms profiles. These Balmer satellites respond
promptly to the continuum variations with delays of only $\sim 3-5$ days.

The H$\alpha$ line shows two velocity-delay structures in the central segment
(within $\pm 5000$ km s-1) at $10^{+3}_{-2}$ and $57^{+3}_{-3}$ light-days.
This might be interpreted as the upper and lower half of an ellipse in the
velocity-delay plane, which might be the signature of a line-emitting ring
orbiting the BH at a radius $R=33.5$ light-days. This interpretation of the
two velocity-delay structures being the signature of a line-emitting ring is
analogous to that of the velocity-delay maps of NGC 5548 based on the STORM
campaign (Horne et al., 2021), and is in turn supported by our finding that
the line profiles in Mrk 926 show the signature of an accretion disk (see
Sect. 4.3). Likewise, the H$\beta$ line shows two velocity-delay structures in
the central region at $5^{+2}_{-2}$ and $48^{+7}_{-6}$ light-days, although
less pronounced. The possible line-emitting ring orbiting the BH has a radius
of $R=26.5$ light-days. However, the relative strength of the responses of the
near and far side of the ring is different in H$\alpha$ and H$\beta$. This
might be caused by optical depths effects. Moreover, H$\alpha$ shows a strong
discontinuity between the signature of the central line profile and the Balmer
satellites, which is less pronounced in the velocity-resolved CCF of H$\beta$.
This might be due to interference of H$\beta$ with the He ii $\lambda 4686$
line.

Assuming the two Balmer velocity-delay structures to be the signature of an
inclined accretion disk, we determined a mean inclination angle of $i\sim
50^{\circ}$. This rather high inclination angle is supported by the remarkably
broad emission lines ($\sim 30\,000$ km s-1 FWZI) in Mrk 926.

####  4.5.3 Origin of the Balmer satellites

The existence of the additional, fast-response Balmer satellite components in
the rms line profiles and in the velocity-resolved CCFs is an indication for
an additional line-emitting component in the nuclear region that might not be
directly connected to the accretion disk. Because of the clear separation
between the central and outer components, their differing variability
behavior, the clearly distinct time lags, as well as the outer rms component
not being connected to the line-flanks, we propose that the Balmer satellite
components may originate in a region that is spatially distinct from the rest
of the BLR. For example, the outer wing components might be connected with a
different spatial region like a hollow outflow cone with a high tilt angle or
small-scale central radio jets. Notably, Ulvestad & Wilson (1984) found a
slightly resolved radio source for Mrk 926 with the VLA at 6 cm. Mundell et
al. (2000) suggested additional extended radio emission on scales between
about 2 and 260 mas based on observations with the VLBA and VLA and the
comparison of their radio beams.

##  5 Summary

We present results of a spectroscopic and photometric monitoring campaign of
the very broad-line AGN Mrk 926, which was carried out with the 10m HET
telescope and the Wise Observatory in 2010. Our findings can be summarized as
follows:

  1. (1)

Mrk 926 is a highly variable AGN. The continuum luminosity showed a drastic
decrease during our campaign. It dropped to less than 50% of its original
luminosity within only 2.5 months.

  2. (2)

Mrk 926 shows very broad H$\alpha$ and H$\beta$ line profiles, with additional
outer Balmer satellite components ranging from $\pm 5000$ to $\pm 13\,000$ km
s-1 that are clearly discernible in the rms spectra.

  3. (3)

The Balmer lines show two velocity-delay structures in their central line
component (within $\pm 5000$ km s-1), at $\sim 10$ and $\sim 57$ light-days
(H$\alpha$) and at $\sim 5$ and $\sim 48$ light-days (H$\beta$). These
structures might be interpreted as the upper and lower halves of an ellipse in
the velocity-delay plane, which might be the signature of a line-emitting ring
orbiting the BH at radii, $R$, of 33.5 and 26.5 light-days.

  4. (4)

The continuum luminosities $\log(\lambda\,L_{\lambda}/{\rm
erg\leavevmode\nobreak\ s}^{-1})$ at 5180 Å (rest frame) of 43.68 to 44.13 are
in good agreement with the established $R_{\text{BLR}}-L_{\text{AGN}}$
relation.

  5. (5)

The derived BH mass of $1.1\times 10^{8}M_{\odot}$ indicates a low Eddington
ratio, which decreased from 8 to 3 percent within only 2.5 months.

  6. (6)

Based on the comparison of the variability behavior of the central line
component and the outer line wings, we speculate that the outer emission
components (the Balmer satellites) originate in a different, spatially
distinct region.

Further densely sampled spectroscopic and photometric studies of this highly
variable AGN are desirable.

###### Acknowledgements.

This work has been supported by the DFG grants KO857/35-1 and CH71/34-3. KH
acknowledges support from STFC grant ST/M001296/1. DC acknowledges support
from ISF grant 2398/19. This paper is based on observations obtained with the
Hobby-Eberly Telescope, which is a joint project of the University of Texas at
Austin, the Pennsylvania State University, Ludwig-Maximilians-Universität
München, and Georg-August-Universität Göttingen.

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##  Appendix A Additional figures

Figure 13: Light curves of the combined continuum flux densities at 5420 Å
(observed frame; in units of 10-15 erg cm-2 s-1 Å-1) as well as of the
H$\alpha$ segment (center, blue, red) light curves.

Figure 14: Light curves of the combined continuum flux densities at 5420 Å
(observed frame; in units of 10-15 erg cm-2 s-1 Å-1) as well as of the
H$\beta$ segment (center, blue, red) light curves.

##  Appendix B Additional tables

Table 9: HET continuum flux densities (observed frame).

Mod. Julian Date | Cont. 4425 Å | Cont. 4755 Å | Cont. 5420 Å | Cont. 5485 Å | Cont. 5755 Å | Cont. 6530 Å | Cont. 7255 Å
---|---|---|---|---|---|---|---
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8)
55396.41 | 8.12 $\pm$ 0.16 | 7.07 $\pm$ 0.14 | 5.40 $\pm$ 0.11 | 5.39 $\pm$ 0.11 | 4.95 $\pm$ 0.10 | 4.65 $\pm$ 0.09 | 4.22 $\pm$ 0.08
55397.42 | 8.73 $\pm$ 0.26 | 7.73 $\pm$ 0.23 | 5.99 $\pm$ 0.18 | 5.94 $\pm$ 0.18 | 5.44 $\pm$ 0.16 | 4.92 $\pm$ 0.15 | 4.48 $\pm$ 0.13
55441.31 | 7.49 $\pm$ 0.11 | 6.61 $\pm$ 0.10 | 5.02 $\pm$ 0.08 | 4.99 $\pm$ 0.07 | 4.60 $\pm$ 0.07 | 4.30 $\pm$ 0.06 | 4.27 $\pm$ 0.06
55443.30 | 7.44 $\pm$ 0.11 | 6.85 $\pm$ 0.10 | 5.39 $\pm$ 0.08 | 5.41 $\pm$ 0.08 | 4.96 $\pm$ 0.07 | 4.53 $\pm$ 0.07 | 4.29 $\pm$ 0.06
55444.30 | 7.73 $\pm$ 0.15 | 7.04 $\pm$ 0.14 | 5.61 $\pm$ 0.11 | 5.66 $\pm$ 0.11 | 5.28 $\pm$ 0.11 | 4.43 $\pm$ 0.09 | 4.40 $\pm$ 0.09
55445.28 | 7.76 $\pm$ 0.12 | 6.99 $\pm$ 0.10 | 5.58 $\pm$ 0.08 | 5.62 $\pm$ 0.08 | 5.27 $\pm$ 0.08 | 4.29 $\pm$ 0.06 | 4.63 $\pm$ 0.07
55446.29 | 7.43 $\pm$ 0.15 | 6.59 $\pm$ 0.13 | 5.00 $\pm$ 0.10 | 5.00 $\pm$ 0.10 | 4.57 $\pm$ 0.09 | 4.30 $\pm$ 0.09 | 4.31 $\pm$ 0.09
55451.27 | 7.09 $\pm$ 0.11 | 6.26 $\pm$ 0.09 | 4.82 $\pm$ 0.07 | 4.84 $\pm$ 0.07 | 4.43 $\pm$ 0.07 | 4.03 $\pm$ 0.06 | 4.03 $\pm$ 0.06
55452.26 | 6.25 $\pm$ 0.09 | 5.95 $\pm$ 0.09 | 5.00 $\pm$ 0.08 | 5.03 $\pm$ 0.08 | 4.77 $\pm$ 0.07 | 4.19 $\pm$ 0.06 | 3.91 $\pm$ 0.06
55453.26 | 7.04 $\pm$ 0.11 | 6.19 $\pm$ 0.09 | 4.90 $\pm$ 0.07 | 4.90 $\pm$ 0.07 | 4.57 $\pm$ 0.07 | 4.07 $\pm$ 0.06 | 4.01 $\pm$ 0.06
55467.23 | 5.41 $\pm$ 0.08 | 5.03 $\pm$ 0.08 | 4.30 $\pm$ 0.07 | 4.37 $\pm$ 0.07 | 4.09 $\pm$ 0.06 | 3.51 $\pm$ 0.05 | 4.03 $\pm$ 0.06
55470.21 | 5.32 $\pm$ 0.08 | 4.69 $\pm$ 0.07 | 3.70 $\pm$ 0.07 | 3.73 $\pm$ 0.06 | 3.47 $\pm$ 0.05 | 3.31 $\pm$ 0.05 | 3.47 $\pm$ 0.05
55471.21 | 4.98 $\pm$ 0.15 | 4.48 $\pm$ 0.13 | 3.61 $\pm$ 0.11 | 3.65 $\pm$ 0.11 | 3.40 $\pm$ 0.10 | 3.33 $\pm$ 0.10 | 3.31 $\pm$ 0.10
55472.22 | 4.57 $\pm$ 0.07 | 4.34 $\pm$ 0.07 | 3.70 $\pm$ 0.06 | 3.77 $\pm$ 0.06 | 3.54 $\pm$ 0.05 | 3.31 $\pm$ 0.05 | 3.34 $\pm$ 0.05
55477.20 | 4.83 $\pm$ 0.14 | 4.46 $\pm$ 0.13 | 3.62 $\pm$ 0.11 | 3.67 $\pm$ 0.11 | 3.40 $\pm$ 0.10 | 3.36 $\pm$ 0.10 | 3.45 $\pm$ 0.10
55480.20 | 4.40 $\pm$ 0.07 | 4.23 $\pm$ 0.06 | 3.65 $\pm$ 0.05 | 3.74 $\pm$ 0.06 | 3.53 $\pm$ 0.05 | 3.44 $\pm$ 0.05 | 3.45 $\pm$ 0.05
55481.20 | 4.59 $\pm$ 0.09 | 4.21 $\pm$ 0.08 | 3.53 $\pm$ 0.07 | 3.56 $\pm$ 0.07 | 3.32 $\pm$ 0.07 | 3.26 $\pm$ 0.07 | 3.32 $\pm$ 0.07
55483.20 | 4.76 $\pm$ 0.07 | 4.23 $\pm$ 0.06 | 3.42 $\pm$ 0.05 | 3.45 $\pm$ 0.05 | 3.24 $\pm$ 0.05 | 3.04 $\pm$ 0.05 | 3.31 $\pm$ 0.05
55494.16 | 4.12 $\pm$ 0.12 | 3.78 $\pm$ 0.11 | 3.30 $\pm$ 0.10 | 3.35 $\pm$ 0.10 | 3.15 $\pm$ 0.09 | 3.10 $\pm$ 0.09 | 3.18 $\pm$ 0.10
55496.15 | 4.23 $\pm$ 0.06 | 3.98 $\pm$ 0.06 | 3.46 $\pm$ 0.05 | 3.51 $\pm$ 0.05 | 3.26 $\pm$ 0.05 | 3.13 $\pm$ 0.05 | 3.39 $\pm$ 0.05
55497.16 | 4.25 $\pm$ 0.09 | 3.95 $\pm$ 0.08 | 3.46 $\pm$ 0.07 | 3.50 $\pm$ 0.07 | 3.29 $\pm$ 0.07 | 3.26 $\pm$ 0.06 | 3.31 $\pm$ 0.07
55498.17 | 4.28 $\pm$ 0.06 | 3.88 $\pm$ 0.06 | 3.29 $\pm$ 0.05 | 3.35 $\pm$ 0.05 | 3.16 $\pm$ 0.05 | 3.03 $\pm$ 0.05 | 3.31 $\pm$ 0.05
55499.16 | 4.36 $\pm$ 0.09 | 4.18 $\pm$ 0.08 | 3.73 $\pm$ 0.07 | 3.84 $\pm$ 0.08 | 3.62 $\pm$ 0.07 | 3.42 $\pm$ 0.07 | 3.48 $\pm$ 0.07
55501.14 | 4.07 $\pm$ 0.06 | 3.89 $\pm$ 0.06 | 3.52 $\pm$ 0.05 | 3.60 $\pm$ 0.05 | 3.45 $\pm$ 0.05 | 3.02 $\pm$ 0.05 | 3.46 $\pm$ 0.05
55505.12 | 4.26 $\pm$ 0.09 | 4.03 $\pm$ 0.08 | 3.53 $\pm$ 0.07 | 3.57 $\pm$ 0.07 | 3.41 $\pm$ 0.07 | 3.12 $\pm$ 0.06 | 3.34 $\pm$ 0.07
55506.12 | 4.05 $\pm$ 0.06 | 3.72 $\pm$ 0.06 | 3.16 $\pm$ 0.05 | 3.18 $\pm$ 0.05 | 2.96 $\pm$ 0.04 | 3.07 $\pm$ 0.05 | 3.10 $\pm$ 0.05
55507.11 | 4.27 $\pm$ 0.06 | 3.81 $\pm$ 0.06 | 3.09 $\pm$ 0.05 | 3.14 $\pm$ 0.05 | 2.92 $\pm$ 0.04 | 2.86 $\pm$ 0.04 | 3.04 $\pm$ 0.05
55508.11 | 4.13 $\pm$ 0.06 | 3.92 $\pm$ 0.06 | 3.38 $\pm$ 0.05 | 3.46 $\pm$ 0.05 | 3.25 $\pm$ 0.05 | 3.13 $\pm$ 0.05 | 3.31 $\pm$ 0.05
55514.12 | 4.26 $\pm$ 0.06 | 3.93 $\pm$ 0.06 | 3.44 $\pm$ 0.05 | 3.52 $\pm$ 0.05 | 3.34 $\pm$ 0.05 | 2.86 $\pm$ 0.04 | 3.32 $\pm$ 0.05
55515.12 | 4.01 $\pm$ 0.12 | 3.57 $\pm$ 0.11 | 3.07 $\pm$ 0.09 | 3.12 $\pm$ 0.09 | 2.91 $\pm$ 0.09 | 2.89 $\pm$ 0.09 | 2.95 $\pm$ 0.09
55519.10 | 3.93 $\pm$ 0.12 | 3.52 $\pm$ 0.11 | 2.90 $\pm$ 0.09 | 2.96 $\pm$ 0.09 | 2.72 $\pm$ 0.08 | 2.76 $\pm$ 0.08 | 2.96 $\pm$ 0.09

999Continuum flux densities in units of 10-15 erg s-1 cm-2 Å-1.

Table 10:  Integrated broad-line fluxes – including the narrow components – of
the Balmer and helium lines for different epochs.

Mod. Julian Date | H$\alpha$ | H$\beta$ | H$\gamma$ | He I${}_{\textrm{center}}$ | He II \+ H$\beta$blue
---|---|---|---|---|---
(1) | (2) | (3) | (4) | (5) | (6)
55396.41 | 2765\. $\pm$ 55. | 1072\. $\pm$ 21. | 152.7 $\pm$ 3.1 | 85.4 $\pm$ 2.2 | 57.9 $\pm$ 1.2
55397.42 | 2965\. $\pm$ 89. | 1143\. $\pm$ 34. | 163.6 $\pm$ 4.9 | 93.0 $\pm$ 1.4 | 61.8 $\pm$ 1.9
55441.31 | 2939\. $\pm$ 44. | 1087\. $\pm$ 17. | 136.6 $\pm$ 2.1 | 96.6 $\pm$ 1.5 | 69.1 $\pm$ 1.0
55443.30 | 2777\. $\pm$ 42. | 1160\. $\pm$ 17. | 129.6 $\pm$ 1.9 | 105.0 $\pm$ 2.7 | 75.0 $\pm$ 1.1
55444.30 | 2784\. $\pm$ 84. | 1192\. $\pm$ 36. | 129.7 $\pm$ 3.9 | 115.8 $\pm$ 1.8 | 77.0 $\pm$ 2.3
55445.28 | 2808\. $\pm$ 42. | 1174\. $\pm$ 18. | 147.3 $\pm$ 2.2 | 109.1 $\pm$ 2.8 | 73.1 $\pm$ 1.1
55446.29 | 2867\. $\pm$ 57. | 1145\. $\pm$ 23. | 146.6 $\pm$ 2.9 | 93.7 $\pm$ 1.5 | 78.7 $\pm$ 1.6
55451.27 | 2832\. $\pm$ 42. | 1144\. $\pm$ 17. | 157.6 $\pm$ 2.4 | 98.8 $\pm$ 2.0 | 76.8 $\pm$ 1.2
55452.26 | 3032\. $\pm$ 45. | 1166\. $\pm$ 18. | 138.2 $\pm$ 2.1 | 132.2 $\pm$ 2.0 | 74.7 $\pm$ 1.1
55453.26 | 3048\. $\pm$ 46. | 1173\. $\pm$ 18. | 160.6 $\pm$ 2.4 | 121.5 $\pm$ 1.9 | 83.1 $\pm$ 1.3
55467.23 | 2759\. $\pm$ 41. | 1148\. $\pm$ 17. | 155.5 $\pm$ 2.3 | 104.6 $\pm$ 1.6 | 70.0 $\pm$ 1.1
55470.21 | 2702\. $\pm$ 41. | 1077\. $\pm$ 16. | 152.4 $\pm$ 2.3 | 92.4 $\pm$ 1.4 | 63.9 $\pm$ 1.0
55471.21 | 2637\. $\pm$ 53. | 1080\. $\pm$ 22. | 156.6 $\pm$ 3.1 | 99.4 $\pm$ 1.5 | 65.5 $\pm$ 1.3
55472.22 | 2775\. $\pm$ 42. | 1097\. $\pm$ 17. | 155.0 $\pm$ 2.3 | 106.2 $\pm$ 1.6 | 65.6 $\pm$ 1.0
55477.20 | 2759\. $\pm$ 55. | 1104\. $\pm$ 22. | 163.5 $\pm$ 3.3 | 98.8 $\pm$ 2.0 | 62.7 $\pm$ 1.3
55480.20 | 2763\. $\pm$ 41. | 1115\. $\pm$ 17. | 164.4 $\pm$ 2.5 | 105.5 $\pm$ 1.6 | 68.1 $\pm$ 1.0
55481.20 | 2834\. $\pm$ 57. | 1108\. $\pm$ 22. | 173.4 $\pm$ 3.5 | 101.3 $\pm$ 1.6 | 65.1 $\pm$ 1.3
55483.20 | 2735\. $\pm$ 41. | 1088\. $\pm$ 16. | 178.4 $\pm$ 2.7 | 93.7 $\pm$ 1.5 | 61.5 $\pm$ 0.9
55494.16 | 2694\. $\pm$ 54. | 1084\. $\pm$ 22. | 164.6 $\pm$ 3.3 | 89.8 $\pm$ 1.8 | 60.2 $\pm$ 1.2
55496.15 | 2960\. $\pm$ 44. | 1137\. $\pm$ 17. | 183.5 $\pm$ 2.8 | 92.2 $\pm$ 1.4 | 64.8 $\pm$ 1.0
55497.16 | 2878\. $\pm$ 58. | 1115\. $\pm$ 22. | 184.7 $\pm$ 3.7 | 92.7 $\pm$ 1.9 | 64.5 $\pm$ 1.3
55498.17 | 2802\. $\pm$ 42. | 1093\. $\pm$ 16. | 187.8 $\pm$ 2.8 | 85.0 $\pm$ 1.7 | 62.1 $\pm$ 0.9
55499.16 | 2725\. $\pm$ 55. | 1148\. $\pm$ 23. | 192.1 $\pm$ 3.8 | 93.1 $\pm$ 2.4 | 67.8 $\pm$ 1.4
55501.14 | 2844\. $\pm$ 43. | 1114\. $\pm$ 17. | 181.2 $\pm$ 2.7 | 92.3 $\pm$ 1.4 | 59.9 $\pm$ 0.9
55505.12 | 2913\. $\pm$ 58. | 1108\. $\pm$ 22. | 190.3 $\pm$ 3.8 | 94.7 $\pm$ 2.4 | 56.3 $\pm$ 1.1
55506.12 | 2825\. $\pm$ 42. | 1058\. $\pm$ 16. | 166.1 $\pm$ 2.5 | 85.1 $\pm$ 1.3 | 53.2 $\pm$ 0.8
55507.11 | 2613\. $\pm$ 39. | 1056\. $\pm$ 16. | 180.3 $\pm$ 2.7 | 73.1 $\pm$ 1.1 | 54.2 $\pm$ 0.8
55508.11 | 2794\. $\pm$ 42. | 1090\. $\pm$ 16. | 177.6 $\pm$ 2.7 | 86.3 $\pm$ 1.3 | 55.1 $\pm$ 0.8
55514.12 | 2661\. $\pm$ 40. | 1067\. $\pm$ 16. | 159.9 $\pm$ 2.4 | 85.8 $\pm$ 1.3 | 49.3 $\pm$ 0.7
55515.12 | 2683\. $\pm$ 80. | 1035\. $\pm$ 31. | 148.3 $\pm$ 4.4 | 79.7 $\pm$ 1.2 | 50.8 $\pm$ 1.5
55519.10 | 2528\. $\pm$ 76. | 980\. $\pm$ 29. | 143.8 $\pm$ 4.3 | 66.8 $\pm$ 1.0 | 39.5 $\pm$ 1.2

101010 Line fluxes in units of 10-15 erg s-1 cm-2.

Table 11: HET continuum light curves used for the FVG method (observed frame).

Mod. | Cont. 4425 Å | Cont. 4425 Å | Cont. 4425 Å | Cont. 4425 Å | Cont. 5420 Å | Cont. 5420 Å | Cont. 5420 Å | Cont. 5420 Å | Cont. 6530 Å | Cont. 6530 Å | Cont. 6530 Å | Cont. 6530 Å
---|---|---|---|---|---|---|---|---|---|---|---|---
Julian Date |  |  | [mJy] | [mJy] |  |  | [mJy] | [mJy] |  |  | [mJy] | [mJy]
|  | ext. corr. |  | ext. corr. |  | ext. corr. |  | ext. corr. |  | ext. corr. |  | ext. corr.
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10) | (11) | (12) | (13)
55396.41 | 8.12 $\pm$ 0.16 | 9.31 $\pm$ 0.19 | 5.30 $\pm$ 0.11 | 6.08 $\pm$ 0.12 | 5.40 $\pm$ 0.11 | 5.99 $\pm$ 0.13 | 5.29 $\pm$ 0.11 | 5.87 $\pm$ 0.12 | 4.65 $\pm$ 0.09 | 5.05 $\pm$ 0.10 | 6.61 $\pm$ 0.13 | 7.18 $\pm$ 0.14
55397.42 | 8.73 $\pm$ 0.26 | 10.01 $\pm$ 0.30 | 5.70 $\pm$ 0.17 | 6.54 $\pm$ 0.20 | 5.99 $\pm$ 0.18 | 6.65 $\pm$ 0.20 | 5.87 $\pm$ 0.18 | 6.51 $\pm$ 0.20 | 4.92 $\pm$ 0.15 | 5.34 $\pm$ 0.17 | 7.00 $\pm$ 0.22 | 7.60 $\pm$ 0.24
55441.31 | 7.49 $\pm$ 0.11 | 8.59 $\pm$ 0.13 | 4.89 $\pm$ 0.08 | 5.61 $\pm$ 0.09 | 5.02 $\pm$ 0.08 | 5.57 $\pm$ 0.09 | 4.92 $\pm$ 0.08 | 5.46 $\pm$ 0.09 | 4.30 $\pm$ 0.06 | 4.67 $\pm$ 0.07 | 6.12 $\pm$ 0.09 | 6.64 $\pm$ 0.10
55443.30 | 7.44 $\pm$ 0.11 | 8.53 $\pm$ 0.13 | 4.86 $\pm$ 0.08 | 5.57 $\pm$ 0.09 | 5.39 $\pm$ 0.08 | 5.98 $\pm$ 0.09 | 5.28 $\pm$ 0.08 | 5.86 $\pm$ 0.09 | 4.53 $\pm$ 0.07 | 4.92 $\pm$ 0.08 | 6.44 $\pm$ 0.10 | 6.99 $\pm$ 0.11
55444.30 | 7.73 $\pm$ 0.15 | 8.87 $\pm$ 0.18 | 5.05 $\pm$ 0.10 | 5.79 $\pm$ 0.12 | 5.61 $\pm$ 0.11 | 6.23 $\pm$ 0.13 | 5.50 $\pm$ 0.11 | 6.10 $\pm$ 0.12 | 4.43 $\pm$ 0.09 | 4.81 $\pm$ 0.10 | 6.30 $\pm$ 0.13 | 6.84 $\pm$ 0.14
55445.28 | 7.76 $\pm$ 0.12 | 8.90 $\pm$ 0.14 | 5.07 $\pm$ 0.08 | 5.81 $\pm$ 0.09 | 5.58 $\pm$ 0.08 | 6.19 $\pm$ 0.09 | 5.47 $\pm$ 0.08 | 6.07 $\pm$ 0.09 | 4.29 $\pm$ 0.06 | 4.66 $\pm$ 0.07 | 6.10 $\pm$ 0.09 | 6.62 $\pm$ 0.10
55446.29 | 7.43 $\pm$ 0.15 | 8.52 $\pm$ 0.18 | 4.85 $\pm$ 0.10 | 5.57 $\pm$ 0.12 | 5.00 $\pm$ 0.10 | 5.55 $\pm$ 0.12 | 4.90 $\pm$ 0.10 | 5.44 $\pm$ 0.11 | 4.30 $\pm$ 0.09 | 4.67 $\pm$ 0.10 | 6.12 $\pm$ 0.13 | 6.64 $\pm$ 0.14
55451.27 | 7.09 $\pm$ 0.11 | 8.13 $\pm$ 0.13 | 4.63 $\pm$ 0.08 | 5.31 $\pm$ 0.09 | 4.82 $\pm$ 0.07 | 5.35 $\pm$ 0.08 | 4.72 $\pm$ 0.07 | 5.24 $\pm$ 0.08 | 4.03 $\pm$ 0.06 | 4.37 $\pm$ 0.07 | 5.73 $\pm$ 0.09 | 6.22 $\pm$ 0.10
55452.26 | 6.25 $\pm$ 0.09 | 7.17 $\pm$ 0.11 | 4.08 $\pm$ 0.06 | 4.68 $\pm$ 0.07 | 5.00 $\pm$ 0.08 | 5.55 $\pm$ 0.09 | 4.90 $\pm$ 0.08 | 5.44 $\pm$ 0.09 | 4.19 $\pm$ 0.06 | 4.55 $\pm$ 0.07 | 5.96 $\pm$ 0.09 | 6.47 $\pm$ 0.10
55453.26 | 7.04 $\pm$ 0.11 | 8.08 $\pm$ 0.13 | 4.60 $\pm$ 0.08 | 5.27 $\pm$ 0.09 | 4.90 $\pm$ 0.07 | 5.44 $\pm$ 0.08 | 4.80 $\pm$ 0.07 | 5.33 $\pm$ 0.08 | 4.07 $\pm$ 0.06 | 4.42 $\pm$ 0.07 | 5.79 $\pm$ 0.09 | 6.28 $\pm$ 0.10
55467.23 | 5.41 $\pm$ 0.08 | 6.21 $\pm$ 0.10 | 3.53 $\pm$ 0.06 | 4.05 $\pm$ 0.06 | 4.30 $\pm$ 0.07 | 4.77 $\pm$ 0.08 | 4.21 $\pm$ 0.07 | 4.68 $\pm$ 0.08 | 3.51 $\pm$ 0.05 | 3.81 $\pm$ 0.06 | 4.99 $\pm$ 0.08 | 5.42 $\pm$ 0.08
55470.21 | 5.32 $\pm$ 0.08 | 6.10 $\pm$ 0.10 | 3.47 $\pm$ 0.06 | 3.99 $\pm$ 0.06 | 3.70 $\pm$ 0.07 | 4.11 $\pm$ 0.08 | 3.63 $\pm$ 0.07 | 4.02 $\pm$ 0.08 | 3.31 $\pm$ 0.05 | 3.59 $\pm$ 0.06 | 4.71 $\pm$ 0.08 | 5.11 $\pm$ 0.08
55471.21 | 4.98 $\pm$ 0.15 | 5.71 $\pm$ 0.18 | 3.25 $\pm$ 0.10 | 3.73 $\pm$ 0.12 | 3.61 $\pm$ 0.11 | 4.01 $\pm$ 0.13 | 3.54 $\pm$ 0.11 | 3.93 $\pm$ 0.12 | 3.33 $\pm$ 0.10 | 3.61 $\pm$ 0.11 | 4.74 $\pm$ 0.15 | 5.14 $\pm$ 0.16
55472.22 | 4.57 $\pm$ 0.07 | 5.24 $\pm$ 0.09 | 2.98 $\pm$ 0.05 | 3.42 $\pm$ 0.06 | 3.70 $\pm$ 0.06 | 4.11 $\pm$ 0.07 | 3.63 $\pm$ 0.06 | 4.02 $\pm$ 0.07 | 3.31 $\pm$ 0.05 | 3.59 $\pm$ 0.06 | 4.71 $\pm$ 0.08 | 5.11 $\pm$ 0.08
55477.20 | 4.83 $\pm$ 0.14 | 5.54 $\pm$ 0.17 | 3.15 $\pm$ 0.10 | 3.62 $\pm$ 0.11 | 3.62 $\pm$ 0.11 | 4.02 $\pm$ 0.13 | 3.55 $\pm$ 0.11 | 3.94 $\pm$ 0.12 | 3.36 $\pm$ 0.10 | 3.65 $\pm$ 0.11 | 4.78 $\pm$ 0.15 | 5.19 $\pm$ 0.16
55480.20 | 4.40 $\pm$ 0.07 | 5.05 $\pm$ 0.09 | 2.87 $\pm$ 0.05 | 3.30 $\pm$ 0.06 | 3.65 $\pm$ 0.05 | 4.05 $\pm$ 0.06 | 3.58 $\pm$ 0.05 | 3.97 $\pm$ 0.06 | 3.44 $\pm$ 0.05 | 3.73 $\pm$ 0.06 | 4.89 $\pm$ 0.08 | 5.31 $\pm$ 0.08
55481.20 | 4.59 $\pm$ 0.09 | 5.27 $\pm$ 0.11 | 3.00 $\pm$ 0.06 | 3.44 $\pm$ 0.07 | 3.53 $\pm$ 0.07 | 3.92 $\pm$ 0.08 | 3.46 $\pm$ 0.07 | 3.84 $\pm$ 0.08 | 3.26 $\pm$ 0.07 | 3.54 $\pm$ 0.08 | 4.64 $\pm$ 0.10 | 5.03 $\pm$ 0.11
55483.20 | 4.76 $\pm$ 0.07 | 5.46 $\pm$ 0.09 | 3.11 $\pm$ 0.05 | 3.57 $\pm$ 0.06 | 3.42 $\pm$ 0.05 | 3.80 $\pm$ 0.06 | 3.35 $\pm$ 0.05 | 3.72 $\pm$ 0.06 | 3.04 $\pm$ 0.05 | 3.30 $\pm$ 0.06 | 4.32 $\pm$ 0.08 | 4.69 $\pm$ 0.08
55494.16 | 4.12 $\pm$ 0.12 | 4.73 $\pm$ 0.14 | 2.69 $\pm$ 0.08 | 3.09 $\pm$ 0.09 | 3.30 $\pm$ 0.10 | 3.66 $\pm$ 0.12 | 3.23 $\pm$ 0.10 | 3.59 $\pm$ 0.11 | 3.10 $\pm$ 0.09 | 3.36 $\pm$ 0.10 | 4.41 $\pm$ 0.13 | 4.79 $\pm$ 0.14
55496.15 | 4.23 $\pm$ 0.06 | 4.85 $\pm$ 0.07 | 2.76 $\pm$ 0.04 | 3.17 $\pm$ 0.05 | 3.46 $\pm$ 0.05 | 3.84 $\pm$ 0.06 | 3.39 $\pm$ 0.05 | 3.76 $\pm$ 0.06 | 3.13 $\pm$ 0.05 | 3.40 $\pm$ 0.06 | 4.45 $\pm$ 0.08 | 4.83 $\pm$ 0.08
55497.16 | 4.25 $\pm$ 0.09 | 4.88 $\pm$ 0.11 | 2.78 $\pm$ 0.06 | 3.18 $\pm$ 0.07 | 3.46 $\pm$ 0.07 | 3.84 $\pm$ 0.08 | 3.39 $\pm$ 0.07 | 3.76 $\pm$ 0.08 | 3.26 $\pm$ 0.06 | 3.54 $\pm$ 0.07 | 4.64 $\pm$ 0.09 | 5.03 $\pm$ 0.10
55498.17 | 4.28 $\pm$ 0.06 | 4.91 $\pm$ 0.07 | 2.80 $\pm$ 0.04 | 3.21 $\pm$ 0.05 | 3.29 $\pm$ 0.05 | 3.65 $\pm$ 0.06 | 3.22 $\pm$ 0.05 | 3.58 $\pm$ 0.06 | 3.03 $\pm$ 0.05 | 3.29 $\pm$ 0.06 | 4.31 $\pm$ 0.08 | 4.68 $\pm$ 0.08
55499.16 | 4.36 $\pm$ 0.09 | 5.00 $\pm$ 0.11 | 2.85 $\pm$ 0.06 | 3.27 $\pm$ 0.07 | 3.73 $\pm$ 0.07 | 4.14 $\pm$ 0.08 | 3.65 $\pm$ 0.07 | 4.06 $\pm$ 0.08 | 3.42 $\pm$ 0.07 | 3.71 $\pm$ 0.08 | 4.86 $\pm$ 0.10 | 5.28 $\pm$ 0.11
55501.14 | 4.07 $\pm$ 0.06 | 4.67 $\pm$ 0.07 | 2.66 $\pm$ 0.04 | 3.05 $\pm$ 0.05 | 3.52 $\pm$ 0.05 | 3.91 $\pm$ 0.06 | 3.45 $\pm$ 0.05 | 3.83 $\pm$ 0.06 | 3.02 $\pm$ 0.05 | 3.28 $\pm$ 0.06 | 4.30 $\pm$ 0.08 | 4.66 $\pm$ 0.08
55505.12 | 4.26 $\pm$ 0.09 | 4.89 $\pm$ 0.11 | 2.78 $\pm$ 0.06 | 3.19 $\pm$ 0.07 | 3.53 $\pm$ 0.07 | 3.92 $\pm$ 0.08 | 3.46 $\pm$ 0.07 | 3.84 $\pm$ 0.08 | 3.12 $\pm$ 0.06 | 3.39 $\pm$ 0.07 | 4.44 $\pm$ 0.09 | 4.82 $\pm$ 0.10
55506.12 | 4.05 $\pm$ 0.06 | 4.65 $\pm$ 0.07 | 2.65 $\pm$ 0.04 | 3.03 $\pm$ 0.05 | 3.16 $\pm$ 0.05 | 3.51 $\pm$ 0.06 | 3.10 $\pm$ 0.05 | 3.44 $\pm$ 0.06 | 3.07 $\pm$ 0.05 | 3.33 $\pm$ 0.06 | 4.37 $\pm$ 0.08 | 4.74 $\pm$ 0.08
55507.11 | 4.27 $\pm$ 0.06 | 4.90 $\pm$ 0.07 | 2.79 $\pm$ 0.04 | 3.20 $\pm$ 0.05 | 3.09 $\pm$ 0.05 | 3.43 $\pm$ 0.06 | 3.03 $\pm$ 0.05 | 3.36 $\pm$ 0.06 | 2.86 $\pm$ 0.04 | 3.10 $\pm$ 0.05 | 4.07 $\pm$ 0.06 | 4.42 $\pm$ 0.07
55508.11 | 4.13 $\pm$ 0.06 | 4.74 $\pm$ 0.07 | 2.70 $\pm$ 0.04 | 3.09 $\pm$ 0.05 | 3.38 $\pm$ 0.05 | 3.75 $\pm$ 0.06 | 3.31 $\pm$ 0.05 | 3.68 $\pm$ 0.06 | 3.13 $\pm$ 0.05 | 3.40 $\pm$ 0.06 | 4.45 $\pm$ 0.08 | 4.83 $\pm$ 0.08
55514.12 | 4.26 $\pm$ 0.06 | 4.89 $\pm$ 0.07 | 2.78 $\pm$ 0.04 | 3.19 $\pm$ 0.05 | 3.44 $\pm$ 0.05 | 3.82 $\pm$ 0.06 | 3.37 $\pm$ 0.05 | 3.74 $\pm$ 0.06 | 2.86 $\pm$ 0.04 | 3.10 $\pm$ 0.05 | 4.07 $\pm$ 0.06 | 4.42 $\pm$ 0.07
55515.12 | 4.01 $\pm$ 0.12 | 4.60 $\pm$ 0.14 | 2.62 $\pm$ 0.08 | 3.00 $\pm$ 0.09 | 3.07 $\pm$ 0.09 | 3.41 $\pm$ 0.10 | 3.01 $\pm$ 0.09 | 3.34 $\pm$ 0.10 | 2.89 $\pm$ 0.09 | 3.14 $\pm$ 0.10 | 4.11 $\pm$ 0.13 | 4.46 $\pm$ 0.14
55519.10 | 3.93 $\pm$ 0.12 | 4.51 $\pm$ 0.14 | 2.57 $\pm$ 0.08 | 2.94 $\pm$ 0.09 | 2.90 $\pm$ 0.09 | 3.22 $\pm$ 0.10 | 2.84 $\pm$ 0.09 | 3.15 $\pm$ 0.10 | 2.76 $\pm$ 0.08 | 3.00 $\pm$ 0.09 | 3.93 $\pm$ 0.12 | 4.26 $\pm$ 0.13

111111 Unless stated otherwise, the flux is given in 10-15 erg s-1 cm-2 Å-1.
In order to calculate the light curves corrected for Galactic extinction, we
adopted Galactic extinction values of $A_{\rm B}=0.149$, $A_{\rm V}=0.113$,
and $A_{\rm R}=0.089$ from Schlafly & Finkbeiner (2011) (NED).

Table 12: Combined HET 5420 Å (5180 Å in rest frame) and Wise Observatory
V-band light curve.

Mod. | Comb. V band | Telescope | Mod. | Comb. V band | Telescope
---|---|---|---|---|---
Julian Date | [10-15 erg s-1 cm-2 Å-1] |  | Julian Date | [10-15 erg s-1 cm-2 Å-1] |
(1) | (2) | (3) | (1) | (2) | (3)
55396.41 | 5.40 $\pm$ 0.11 | HET | 55481.80 | 3.59 $\pm$ 0.15 | Wise
55397.42 | 5.99 $\pm$ 0.18 | HET | 55482.85 | 4.01 $\pm$ 0.15 | Wise
55436.92 | 5.68 $\pm$ 0.14 | Wise | 55483.20 | 3.42 $\pm$ 0.05 | HET
55437.86 | 5.64 $\pm$ 0.08 | Wise | 55483.82 | 3.60 $\pm$ 0.13 | Wise
55438.80 | 5.71 $\pm$ 0.14 | Wise | 55484.80 | 3.44 $\pm$ 0.21 | Wise
55439.96 | 5.48 $\pm$ 0.13 | Wise | 55485.81 | 3.36 $\pm$ 0.11 | Wise
55440.90 | 5.28 $\pm$ 0.41 | Wise | 55489.79 | 3.18 $\pm$ 0.30 | Wise
55441.31 | 5.02 $\pm$ 0.08 | HET | 55491.66 | 3.47 $\pm$ 0.26 | Wise
55441.80 | 5.48 $\pm$ 0.14 | Wise | 55492.67 | 3.44 $\pm$ 0.19 | Wise
55442.88 | 5.36 $\pm$ 0.09 | Wise | 55494.16 | 3.30 $\pm$ 0.10 | HET
55443.30 | 5.39 $\pm$ 0.08 | HET | 55494.67 | 3.32 $\pm$ 0.26 | Wise
55443.78 | 5.28 $\pm$ 0.29 | Wise | 55495.68 | 3.22 $\pm$ 0.19 | Wise
55444.30 | 5.61 $\pm$ 0.11 | HET | 55496.15 | 3.46 $\pm$ 0.05 | HET
55445.28 | 5.58 $\pm$ 0.08 | HET | 55496.75 | 3.25 $\pm$ 0.15 | Wise
55445.92 | 5.13 $\pm$ 0.19 | Wise | 55497.16 | 3.46 $\pm$ 0.07 | HET
55446.29 | 5.00 $\pm$ 0.10 | HET | 55497.73 | 3.32 $\pm$ 0.20 | Wise
55446.84 | 5.14 $\pm$ 0.15 | Wise | 55498.17 | 3.29 $\pm$ 0.05 | HET
55447.81 | 5.15 $\pm$ 0.18 | Wise | 55499.16 | 3.73 $\pm$ 0.07 | HET
55448.81 | 5.17 $\pm$ 0.13 | Wise | 55501.14 | 3.52 $\pm$ 0.05 | HET
55449.78 | 5.16 $\pm$ 0.15 | Wise | 55501.69 | 3.11 $\pm$ 0.32 | Wise
55451.27 | 4.82 $\pm$ 0.07 | HET | 55502.71 | 3.04 $\pm$ 0.15 | Wise
55451.79 | 5.04 $\pm$ 0.20 | Wise | 55503.79 | 3.30 $\pm$ 0.11 | Wise
55452.26 | 5.00 $\pm$ 0.08 | HET | 55504.79 | 3.14 $\pm$ 0.17 | Wise
55453.26 | 4.90 $\pm$ 0.07 | HET | 55505.12 | 3.53 $\pm$ 0.07 | HET
55453.78 | 4.84 $\pm$ 0.11 | Wise | 55505.76 | 3.14 $\pm$ 0.14 | Wise
55457.80 | 4.60 $\pm$ 0.25 | Wise | 55506.12 | 3.16 $\pm$ 0.05 | HET
55464.79 | 4.15 $\pm$ 0.18 | Wise | 55506.80 | 3.19 $\pm$ 0.15 | Wise
55465.76 | 4.15 $\pm$ 0.13 | Wise | 55507.11 | 3.09 $\pm$ 0.05 | HET
55466.85 | 4.24 $\pm$ 0.08 | Wise | 55507.72 | 3.06 $\pm$ 0.43 | Wise
55467.23 | 4.30 $\pm$ 0.07 | HET | 55508.11 | 3.38 $\pm$ 0.05 | HET
55467.84 | 4.09 $\pm$ 0.10 | Wise | 55508.77 | 3.25 $\pm$ 0.13 | Wise
55468.83 | 4.09 $\pm$ 0.16 | Wise | 55510.66 | 2.93 $\pm$ 0.47 | Wise
55469.84 | 4.12 $\pm$ 0.11 | Wise | 55512.75 | 3.09 $\pm$ 0.16 | Wise
55470.21 | 3.70 $\pm$ 0.07 | HET | 55514.12 | 3.44 $\pm$ 0.05 | HET
55470.81 | 4.25 $\pm$ 0.22 | Wise | 55515.12 | 3.07 $\pm$ 0.09 | HET
55471.21 | 3.61 $\pm$ 0.11 | HET | 55519.10 | 2.90 $\pm$ 0.09 | HET
55471.84 | 3.95 $\pm$ 0.22 | Wise | 55519.69 | 3.09 $\pm$ 0.15 | Wise
55472.22 | 3.70 $\pm$ 0.06 | HET | 55520.65 | 3.01 $\pm$ 0.34 | Wise
55472.80 | 4.45 $\pm$ 0.16 | Wise | 55522.67 | 3.09 $\pm$ 0.45 | Wise
55473.92 | 3.84 $\pm$ 0.40 | Wise | 55524.79 | 3.09 $\pm$ 0.22 | Wise
55475.76 | 3.88 $\pm$ 0.18 | Wise | 55526.71 | 3.19 $\pm$ 0.23 | Wise
55476.85 | 3.88 $\pm$ 0.16 | Wise | 55527.76 | 3.23 $\pm$ 0.10 | Wise
55477.20 | 3.62 $\pm$ 0.11 | HET | 55530.71 | 3.40 $\pm$ 0.34 | Wise
55478.89 | 3.71 $\pm$ 0.21 | Wise | 55531.70 | 3.17 $\pm$ 0.42 | Wise
55479.88 | 3.34 $\pm$ 0.75 | Wise | 55533.75 | 3.37 $\pm$ 0.19 | Wise
55480.20 | 3.65 $\pm$ 0.05 | HET | 55537.67 | 3.11 $\pm$ 0.24 | Wise
55480.69 | 3.84 $\pm$ 0.28 | Wise | 55540.72 | 3.34 $\pm$ 0.17 | Wise
55481.20 | 3.53 $\pm$ 0.07 | HET |  |  |