Spectropolarimetry and spectral decomposition of high-accreting
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Astronomy & Astrophysics manuscript no. aa-main ©ESO 2022
March 1, 2022
Spectropolarimetry and spectral decomposition of high-accreting
Narrow Line Seyfert 1 galaxies?
Marzena Śniegowska??1, 2 , Swayamtrupta Panda1, 2, 3,??? , Bożena Czerny2 ,
Ðorge Savić4, 5 , Mary Loli Martínez-Aldama2, 6 , Paola Marziani7 , Jian-Min Wang8, 9, 10 , Pu Du8 , and Luka Č.
Popović4, 11
1
Nicolaus Copernicus Astronomical Center, Polish Academy of Sciences, ul. Bartycka 18, 00-716 Warsaw, Poland
2
Center for Theoretical Physics, Polish Academy of Sciences, Al. Lotników 32/46, 02-668 Warsaw, Poland
3
Laboratório Nacional de Astrofísica - MCTIC, R. dos Estados Unidos, 154 - Nações, Itajubá - MG, 37504-364, Brazil
arXiv:2202.13839v1 [astro-ph.GA] 28 Feb 2022
4
Astronomical Observatory Belgrade Volgina 7, P.O. Box 74 11060, Belgrade, 11060, Serbia
5
Université de Strasbourg, CNRS, Observatoire Astronomique de Strasbourg, UMR 7550, 11 rue de l’Université, F-67000 Stras-
bourg, France
6
Departamento de Astronomia, Universidad de Chile, Camino del Observatorio 1515, Santiago, Chile
7
Istituto Nazionale di Astrofisica (INAF), Osservatorio Astronomico di Padova, 35122 Padova, Italy
8
Key Laboratory for Particle Astrophysics, Institute of High Energy Physics, Chinese Academy of Sciences, 19B Yuquan Road,
Beijing 100049, People’s Republic of China
9
School of Astronomy and Space Sciences, University of Chinese Academy of Sciences, 19A Yuquan Road, Beijing 100049, China
10
National Astronomical Observatories of China, Chinese Academy of Sciences, 20A Datun Road, Beijing 100020, China
11
Department of Astronomy, Faculty of Mathematics, University of Belgrade, Studentski trg 16, 11000 Belgrade, Serbia
March 1, 2022
ABSTRACT
Context. Narrow Line Seyfert 1 (NLSy1) galaxies have been shown to have high Eddington ratios and relatively small black hole
masses. The measurement of the black hole masses is based on the virial relation which is dependent on the distribution of the line-
emitting gas and the viewing angle to the source. Spectropolarimetry enables us to probe the geometry of this line-emitting gas and
allows us to estimate independently the viewing angle of the source by comparing the spectrum viewed under natural light and in the
polarized light.
Aims. We aim: (i) to estimate the virial factor using the viewing angles inferred from the spectropolarimetric measurements for a
sample of NLSy1s which influences the measurement of the black hole masses. (ii) to model the natural and polarized spectra around
the Hα region using spectral decomposition and spectral fitting techniques. (iii) to infer the physical conditions of the broad line
region and the scattering medium responsible for the polarization of the Hα emission line (and continuum), and to model the Stokes
parameters, using the polarization radiative transfer code STOKES.
Methods. We performed spectropolarimetric observations of three NLSy1 - Mrk 1044, SDSS J080101.41+184840.7, and IRAS
04416+1215 using the European Southern Observatory’s Very Large Telescope. We use ESO Reflex workflow to perform standard
data reduction and extract the natural and polarized spectra. We then model the Hα region in the reduced spectra using IRAF spectral
fitting procedures and estimate the Stokes parameters and the viewing angles of the three sources. We model the Stokes parameters
and infer the properties of the scattering media - located in the equatorial and polar regions, and simulate the spectra observed both in
natural light and in polarized light using the polarization radiative transfer code STOKES.
Results. The viewing angles recovered for the three sources indicate that they occupy separate locations in the viewing angle plane,
from an almost face-on (IRAS 04416+1215), to an intermediate (SDSS J080101.41+184840.7), to a highly inclined (Mrk 1044)
orientation. Nevertheless, we confirm that all three sources are high Eddington ratio objects. We are successful in recovering the
observed Hα line profile both in the natural and polarized light using the STOKES modeling. We recover the polarization fractions
of the order of 0.2-0.5% for the three sources although the recovery of the phase angle is sub-optimal, mainly due to the noise in
the observed data. Our principal component analysis shows that the sample of the 25 sources including our sources, Fairall 9 from
(Jiang et al. 2021), and sources from (Capetti et al. 2021) are mainly driven by the black hole mass and Eddington ratio. We re-affirm
the connection of the strength of the optical FeII emission with the Eddington ratio. But, the dependence on the viewing angle is
moderate, more like a secondary effect.
Key words. Galaxies: active – Galaxies: Seyfert – quasars: emission lines – accretion, accretion disks – Techniques: spectroscopic
– Techniques: polarimetric – Radiative transfer – Methods: statistical
1. Introduction
?
Based on data collected under programme 098.B-0426(B) Narrow Line Seyfert 1 (NLSy1) galaxies and quasars with rel-
?? atively narrow permitted lines are generally believed to have
msniegowska@camk.edu.pl
???
CNPq Fellow relatively high Eddington ratios (e.g. Mathur 2000a; Wang &
Article number, page 1 of 20A&A proofs: manuscript no. aa-main
Netzer 2003; Grupe 2004). NLSy1s are contextualized as a sub- sion lines in polarized light are much broader than in the unpolar-
population within the larger Population of sources accreting at ized spectrum. Even simply comparing the kinematic line width
relatively high rates (λEdd > 0.2 Population A, following Sulen- in total and in polarized light can verify the statement about the
tic et al. 2000a). The question remains: how high these ratios actual value of the black hole mass and the Eddington ratio of
can be? This is an important question from the theoretical point a source (Baldi et al. 2016). A new independent method of the
of view: we have a basic understanding of the accretion when black hole mass measurement has been introduced, based on the
the Eddington ratio is moderate but the models are unreliable change of the polarization angle across the broad emission line.
in the extremely high accretion rates. Moreover, strong outflows This change depends on the rotational velocity in the BLR, and
might, in principle, prevent too high accretion rates to happen it allows to directly measure the black hole mass (Afanasiev &
in active galactic nuclei (AGN). Therefore, the study of sources Popović 2015). It was proposed by Afanasiev et al. (2014, 2015),
considered as highly super-Eddington is of extreme importance. and applied to mostly Hα line based observations (e.g. Afanasiev
The determination of the Eddington ratio requires both the & Popović 2015; Afanasiev et al. 2015, 2019). using the spec-
measurement of the bolometric luminosity of the source and tropolarimetry data from the VLT/FORS2. (Jiang et al. 2021)
its black hole mass. Currently, the reverberation mapping (RM) show that for Fairall 9 the mass measurements from Hα and Hβ
method is considered to be the most reliable method of the black lines are consistent.
hole mass measurement (Peterson 1993; Peterson et al. 2004; The first spectropolarimetric black hole mass measurement
Cackett et al. 2021). It measures the profiles of the broad emis- with MgII line, using the rotation of the polarization angle was
sion lines coming from the broad line region (BLR) and the de- recently performed by (Savić et al. 2020). Their result is in good
lay between the variable continuum and the line response, which agreement with other black hole mass estimation techniques (see
constrains the velocity and the BLR size. The method is based on Afanasiev & Popović 2015; Afanasiev et al. 2019; Savić et al.
assumption that the BLR is predominantly in Keplerian motion, 2020).
and the results generally agree with the results based on the black We present the spectropolarimetric measurements and spec-
hole mass - bulge velocity relation (see e.g., Ferrarese & Mer- tral decomposition of the Hα spectral region of 3 NLSy1
ritt 2000; Gültekin et al. 2009), which supports the underlying galaxies: Mrk 1044, SDSS J080101.41+184840.7 and IRAS
assumption. However, some objects with Hβ full width at half 04416+12151 performed with the VLT/FORS2 instrument. We
maximum (FWHM) of about 2000 km s−1 show Hβ lags much also show the spectral decomposition of the archival spectra for
shorter than the well-known radius-luminosity relation (Bentz these three sources in Hβ and Hα region. All of these sources
et al. 2013), when we compare objects with the same luminosity are part of the SEAMBH Project sample (Wang et al. 2014a),
(Du et al. 2014; Wang et al. 2014a; Du et al. 2015, 2016, 2018). selected for reverberation monitoring as candidates for super-
This implies a surprisingly small black hole mass and, for a mea- Eddington accreting sources. We briefly describe archival infor-
sured bolometric luminosity, a very high Eddington ratio. mation of 3 chosen sources. Mrk 1044 (z = 0.016451±0.000037;
Such high Eddington objects challenge the theoretical mod- RA 02:30:05.5, DEC -08:59:53 from NED2 ) is one of the near-
els of the accretion process, and they are important for under- est and the brightest Seyfert galaxies (Véron-Cetty & Véron
standing the rapid growth of the black hole mass at high red- 2006) (V mag = 14.5 in r band, from NED). The line width of
shifts. Theoretically, these objects should be modeled with slim the Balmer lines was already measured by Rafanelli & Schulz
accretion disks (Abramowicz et al. 1988). The innermost part of (1983). The authors gave the values of FWHM of 3600 km
such a disk is geometrically puffed up, a significant fraction of s−1 for the broad components of both Hα and Hβ line. They
the radiation is expected to be released in a collimating funnel. stressed the comparable intensity of the narrow components,
In such a picture, the irradiation of the outer parts of the disk, with the FWHM width of 1000 km s−1 , and the very low value of
where the BLR is located, is geometrically less efficient (Wang [OIII]/Hβ ratio (below 0.6). The source was later reclassified as
et al. 2014b). Thus, the BLR may be located elsewhere, e.g. in the NLSy1 galaxy (Osterbrock & Pogge 1985). This source was
the biconical flow (Corbett et al. 2000). If so, the BLR we ob- also observed in X-rays, the viewing angle obtained by Mallick
serve would not be in Keplerian motion and the black hole mass et al. (2018) is i = 47.2+1.0
−2.5 from the joint fitting of Swift, XMM-
measurement may be highly biased. In addition, in extreme cases Newton, and NuSTAR X-ray spectra.
of a source seen along the symmetry axis, the lines can be also
SDSS J080101.41+184840.7 (z = 0.13954 ± 0.00001; RA
much narrower due to purely geometrical factor (e.g. Baldi et al.
08:01:01.41, DEC +18:48:40.78. from NED) with V mag =
2016). A comparison with stellar dispersion measured black hole
16.88, from Véron-Cetty & Véron (2010), is the second object
masses may not give the final argument. NLSy1 are sometimes
in our sample. (Liu et al. 2021) model XMM-Newton spectrum
argued to be outliers from the black hole mass - stellar disper-
and obtain steep photon index (2.33 ±0.06) using a power-law
sion relation due to evolutionary effect. Their black holes are too
model modified by Galactic absorption. However, in the model
small for their bulge masses and they accrete vigorously, increas-
used by (Liu et al. 2021) the viewing angle was not one of model
ing mass and moving towards the standard relation appropriate
parameters.
for mature AGN (e.g. Mathur 2000b; Mathur et al. 2001).
Using spectropolarimetry, we obtain new and independent The detailed analysis of broad band observations of IRAS
insight into the geometry, full velocity field of the line emitting 04416+1215 (r mag = 16.24, z = 0.089 was recently reported
material as well as into the location of the fully ionized scattering by Tortosa et al. (2022). This source shows narrow Hβ line,
medium. Spectropolarimetry revealed the nature of type 2 AGN (FWHM = 1670 km s−1 , Moran et al. 1996) and very broad
as active galaxies with the BLR hidden by the dusty/molecular [OIII], (FWHM = 1150 km s−1 , Véron-Cetty et al. 2001) lines.
torus (Antonucci & Miller 1985), but later Smith et al. (2004) In their work, Tortosa et al. (2022) the authors showed that
and Smith et al. (2005) showed the advantage of using this tech- the best-fitting model in the X-ray band is composed of a soft-
nique to study type 1 AGN as well. The wavelength-dependent excess, three ionized outflows, neutral absorption, and a reflec-
polarization angle also indicates whether the scattering takes
place in the polar or the equatorial region, and measures the 1
also known as SDSS J044428.77+122111.7
2
viewing angle of the system: in low viewing angle sources, emis- https://ned.ipac.caltech.edu/
Article number, page 2 of 20Marzena Śniegowska et al.: Spectropolarimetry and spectral decomposition of high-accreting Narrow Line Seyfert 1 galaxies
tion component. Prominent soft X-ray excess supports the claim was different in both spectra. The same effect was observed in
for super-Eddington accretion rate in this source. the reduced standard stars obtained with ESO Reflex, where an
Du et al. (2016), through dedicated monitoring campaign expected blackbody continuum was not recovered. It suggests
using the Lijiang 2.4m telescope under the SEAMBH Project, that one correction was not applied by the pipeline. To correct the
report the measurements for the Hβ FWHMs for these three continuum form in the natural light spectrum, we used the IRAF
sources (see Table 5 in their paper) - Mrk 1044 (1178±22 routines standard and sensfunc, and the standard star HD 64299.
km s−1 ), SDSS J080101.41+184840.7 (1930±18 km s−1 ), and, After the correction, a good agreement was obtained between the
IRAS 04416+1215 (1522±44 km s−1 ). natural light spectrum and the one from the SDSS catalog.
The paper is organized as follow: Details on the performed
observations and data reduction are reported in Section 2. In
Section 3, we present the analysis of spectropolarimetric data 3. Data analysis
obtained using the Very Large Telescope and the archival spec- 3.1. Standard star
troscopic data for each source. In Section 4 we show the spec-
tral decomposition for natural and polarized light for our objects We present in Table 2 the measurements for standard stars HD
and spectral modeling. Next, we perform polarization radiative 64299 and Ve 6-23 used in our work. Cikota et al. (2017) per-
transfer modeling using STOKES and compare it with our ob- form measurements for Ve 6-234 , between 7.20% and 7.03%
served results in Section 5. We discuss the implications of these for I band. In our case we obtained 7.12 %, so comparing with
results, outline the physical parameters responsible for the trends Cikota et al. (2017) we assume our systematic error at the level
observed in our sample using principal component analysis, and of 0.03%. Sosa et al. (2019) obtained 6.23± 0.03% polariza-
summarize our findings from this study in Section 6. In this work tion level for this star, but observations were performed with a
we use ΛCDM cosmology (ΩΛ = 0.7, ΩM = 0.3, H0 = 70 km s−1 different instrument. The polarization angle for Ve 6-23 is 170,
Mpc−1 ). which is consistent with one calculated by Cikota et al. (2017)
(171.95±0.01). In the case of HD 64299, Sosa et al. (2019) per-
formed the measurements for this star and for I band obtained P
2. Observations and Data Reduction = 0.16 ± 0.01% (Table 3 therein). Our measurements for this star
P = 0.018%, differ from Sosa et al. (2019) which may be con-
Our spectropolarimetric observations were performed with the
nected with different order of steps in computing polarization
FORS2 instrument mounted at the UT1 telescope of the 8.2m
degree (or fraction). Taking the formula forpthe mean polariza-
ESO VLT. The observations have been taken using the GRISM-
300I in combination with the blocking filter OG590. The 0.7 arc- tion using the Stokes parameters Pmean = Q2mean + Umean 2 , we
sec wide slit was oriented along the parallactic angle, and the can square Q and U arrays and then calculate mean or firstly
multi-object spectropolarimetry observations were performed calculate the Qmean and Umean and then square the single value.
with a 2048 × 2048 pixel CCD with a spatial resolution of 0.126 We decided to choose the latter procedure, because squaring all
arcsec pixel−1 . values along the array may artificially increase the polarization
The spectropolarimetric observation of Mrk 1044 was ac- degree since we do not consider different signs in flux. In the
complished 2016 Oct 27 (JD 57688.1145), the source SDSS case of HD 64299 using another method we increase P from
J080101.41+184840.7 was observed on 2016 Dec 26, 2016 Dec 0.018% to 0.32%, but in the case of the polarized star, we obtain
30, and 2017 Jan 02. In the case of IRAS 04416+1215, it was the same value for both procedures.
2016 Nov 06. For all observational runs, we used the Patat &
Romaniello (2006) prescription of observing with four quarter- 3.2. Spectropolarimetry of NLSy1
wave plate angles (0, 22.5, 45, 67.5 deg.). As we show in Table
1: for Mrk 1044 at each angle, 2 exposures were taken (4 × 2 In this work for our sample, we present three types of polar-
spectra), each exposure lasting 206 s, for IRAS 04416+1215 3 ization measurements: mean polarization of each source, polar-
exposures lasting 455 s and for SDSS J080101.41+184840.7 3 ization of continuum, and polarization of the line corrected for
exposures 250 s on 2017 Jan 02, for other days 3 exposures 280 continuum polarization. By a mean polarization (Pmean ), we un-
s. derstand the polarization calculated from mean Q and U of the
Together with sources, two standard stars were observed, whole wavelength range of the spectra. As a continuum polariza-
one highly polarized and one unpolarized star (Ve6-23 and tion (Pcont ), we consider the polarization calculated from mean
HD62499). Observations were performed in the service mode Q and U at two windows surrounding Hα from its blue and red
(program ID: 098.B-0426(B), PI: B. Czerny), and observing side, each window has width ≈ 100Å and they are separated from
conditions were good, with an atmospheric seeing ∼ 0.8. center of Hα at least 300Å. The continuum polarization angle
To reduce data we use ESO Reflex3 (Hook et al. 2008). The (θcont ) is also calculated from mean Q and U of the same win-
workflow combines the biases into a master bias which is then dows.
subtracted from the science image. The same procedure is per- For the line polarization measurement (Pline ), we consider
formed with the lamp flats, and the science image is flat-fielded. the polarization of the region with the width ≈ 50Å surrounding
The workflow then removes the cosmic ray events and performs the center of the emission line.
the wavelength calibration using the standard He-Ar arc lamps. We decided to correct the line polarization measurement for
As final products, among others, we obtain the extracted 1D the contamination by the underlying continuum and we per-
spectra with the Stokes parameters (U and Q), total linear po- formed the following calculations. We calculate (not normalized
larization (L), and flux in ADU/s (I) as a function of wavelength to the intensity, I) Stokes parameters Q and U from the left and
with the associated uncertainties. right parts of the continuum (which are the same length). Then,
A comparison between the natural light spectrum and the one we calculate the average vector of Qcont · Icont and Ucont · Icont for
from the SDSS catalog showed that the AGN continuum form the continuum. Finally, we calculate mean value of Qcont · Icont
3 4
https://www.eso.org/sci/software/esoreflex/ Ve 6-23 in Cikota et al. (2017) is marked as Vela1 95
Article number, page 3 of 20A&A proofs: manuscript no. aa-main
Table 1. Observed NLSy1 for this work. Columns show name of the source, date of observation, number of exposures and the exposure time.
NAME DATE Number of exposures Exposure time [s]
Mrk 1044 2016-10-27 2 206
SDSS J080101.41+184840.7 2016-12-26 3 280
2016-12-30 3 280
2017-01-02 3 250
IRAS 04416+1215 2016-11-06 3 455
Table 2. Observed standard stars. Type indicates if the star is polarized (Pol.) or unpolarized (unPol.) after Cikota et al. (2017). Columns show the
name of the source, the coordinates (taken from the SIMBAD Astronomical Database), the averaged P, the averaged θ and the averaged Stokes
parameters: Q and U.
NAME RA DEC P θ Q U
[J2000] [J2000] [%] [◦ ] [%] [%]
HD64299 07:52:25.4 -23:17:47.5 0.018±0.019 89±66 -0.017±0.009 0.00638±0.00687
(unPol.)
Ve6-23 09:06:00.01 -47:18:58.2 7.122±0.034 170±1 6.838±0.0648 -1.990±0.441
(Pol.)
and Ucont · Icont along the wavelength, and this values we subtract 3.3. Archival spectra of NLSy1
from each point of vectors Qcont ·Icont and Ucont ·Icont , accordingly.
Only now we normalize Qcont · Icont and Ucont · Icont to the inten- We decided to explore the physical parameters of our sam-
sity, by dividing both vectors by Iline . The further procedure of ple in the optical band starting from archival, reduced data.
calculating the fraction of linearly polarized radiation is the same For Mrk1044 we used spectrum from Jones et al. (2009) and
as for the Pmean and Pcont . We show windows for continuum and for other two sources, SDSS J080101.41+184840.7 and IRAS
line polarization for SDSS J080101.41+184840.7 marked in red 04416+1215, data from SDSS (Blanton et al. 2017)). For those
in Figure 1. We use the same windows for each object. spectra, we perform spectral decomposition of Hβ and Hα re-
gion. Then, we use values obtained for Hα region from archival
spectra as initial values for decomposition of non-polarized spec-
tra of Hα from FORS2/PMOS.
To analyze all spectra we used the specfit task from IRAF
(Kriss 1994), including various components in two ranges: Hβ
SDSS J080101.41+184840.7 region (4400-5200Å) and Hα region (6200-6800Å). In the Hβ
300 range, we basically follow the method presented in Negrete et al.
200 (2018). In both regions, we modeled the accretion disk emission
and the starlight contribution to the spectrum assuming a power-
I
100 law shape without performing a more complicated procedure for
the starlight component, since ranges are relatively narrow. To
0 model it we used the continuum windows around 4430, 4760,
and 5100Å in case of Hβ (e.g. Francis et al. 1991), and in case
1.0 of Hα 5700-6000Å (e.g. Kuraszkiewicz et al. 2002). In both re-
gions, we consider the Fe ii template from Marziani et al. (2009).
P (%)
0.5 The [OIII] doublet lines in Hβ region were modeled with two
Gaussians which represent narrow and semi broad components.
0.0 We assume the theoretical ratio of 1 to 3 between the strengths of
200 the components (Dimitrijević et al. 2007). The [N II] and [S II]
doublet lines in Hα region were modeled with Gaussian shape,
with the assumption of the theoretical ratio 1 to 3 between the
100
(°)
strengths of them, and FWHM of lines fixed to [OIII] narrow
component. For each Balmer line we decided to use Lorentzian
0 shape for the broad component since it is considered to be more
6000 6200 6400 6600 6800 7000 suitable for the NLSy1 than the Gaussian (e.g. Laor et al. 1997).
Wavelength [Å]
We modeled the Balmer emission lines taking into account two
components: one Lorentzian component and one narrow Gaus-
Fig. 1. Polarization windows for SDSS J080101.41+184840.7. As red sian component (to reproduce the peak of the line). Moreover,
we mark windows in which we calculate polarization: for continuum we keep the FWHM of narrow Gaussian component fixed to
polarization, we consider two surrounding Hα from its blue and red FWHM of narrow component for all forbidden lines and relax
side, each window has width ≈ 100Å and they are distant from the cen- it just for the last iteration of the fitting. To estimate errors, we
ter of Hα at least 300Å. For polarization of the line, we consider the red use the maximum and minimum continuum levels, and for each
window mark on the line. Each central window has width ≈50Å. case, we perform the fit of the spectrum. We assume that the dis-
tribution of errors follows the triangular distribution (d’Agostini
2003). For each line measurement we calculate the variance us-
Article number, page 4 of 20Marzena Śniegowska et al.: Spectropolarimetry and spectral decomposition of high-accreting Narrow Line Seyfert 1 galaxies
ing the following formula for the triangular distribution: the mean of Stokes parameters Q and U of the whole wavelength
range of the spectra using the formula:
∆2 x+ + ∆2 x− + ∆x+ + ∆x−
σ2 (X) = (1) q
18 Pmean = Q2mean + Umean2 . (2)
where ∆x+ and ∆x− are differences between measurements with The polarization of the continuum region (column 4) and line re-
maximum and best continuum, and with best and minimum con- gion (column 6) is calculated similarly, including different wave-
tinuum respectively. This can be explained as the linear decrease length ranges of the spectra. Additionally, the polarization of the
in either side of the maximum of the distribution (which is the lines includes the ’continuum polarization correction’, which we
best fit) to the values obtained for maximum and minimum con- explained in Section 3.2. The ’continuum polarization correc-
tributions of the continuum. We find this analytical method easy tion’ does not change the physical interpretation of the results.
and sufficiently precise. We propagate uncertainties using stan- For the SDSS J080101.41+184840.7, we obtain Pline of 0.31%
dard formulas of error propagation for reported values in this with the correction, and Pline equal 0.42% without this correc-
work. tion. Both values are lower than Pmean and Pcont for this source,
but by including this correction, the fraction of linearly polar-
4. Results ized radiation is reduced. The polarization position angle (which
is the angle of maximum polarization) we calculate using the
4.1. NLSy1 spectra decomposition in natural light following formulae (eq. 8 in Bagnulo et al. 2009):
!
We fit each spectrum individually, and for each, we follow the 1 −1 U
same steps to keep the same accuracy. As a first step, we de- θ = tan (if Q > 0 and U ≥ 0);
2 Q
compose the Hβ and Hα emission line profiles in natural light !
from archival spectra. Then we fit Hα in natural light obtained 1 U
θ = tan−1 + 180◦ (if Q > 0 and U < 0); (3)
from our FORS2/VLT observations. To remain consistent, we 2 Q
keep the same values for Hα components obtained from archival 1 U
!
spectra as input, and then fit the Hα in natural light from VLT. θ = tan−1 + 90◦ (if Q < 0)
2 Q
We present our measurements of FWHM of Balmer lines’ com-
ponents in Table 3. Each line is modeled with 3 components: We use mean Q and U from the whole region (column 3), the
broad (Lorentzian), narrow (Gaussian), and blue (Gaussian). In continuum region (column 5), and the line region (column 7).
all cases FWHM of Hα broad component (BC) is slightly nar- We show in Figure 3 the intensity of the Hα, the polarization
rower than Hβ one. Whereas, the blue components are broader percentage, and the polarization angle for each object. In red we
ones. We present those fits and fits from archival spectra in Fig- mark the range which we used to compute the polarization de-
ure 2. We do not see any significant differences between Hα gree (or fraction) of the line.
fits from archival spectra and the VLT spectra. Thus, we treat Overall, the level of polarization in our three sources is very
the archival spectra of Hα just as input parameters. For fur- low, therefore the measurement of the change of the polarization
ther analysis, we use the archival spectra of Hβ (since we did angle across the broad emission line is hard to measure. There-
not perform observation for this line), and the VLT Hα spec- fore, we concentrate on the estimates of the viewing angle to
tra. Balmer lines are slightly asymmetric. In the case of IRAS improve the black hole mass measurement.
04416+1215 for both lines and for Mrk 1044 for Hβ blue compo- In addition, we notice that the polarized flux profile has
nents (which may indicate outflow) are visible in fits. However, a peaky appearance with a width not so much larger than
Hβ blue component in Mrk 1044 is almost negligible, For SDSS the one measured in natural light, at least for Mrk 1044 and
J080101.41+184840.7, blue components are the weakest, but SJ080101.41+184840.7 (the case of IRAS 04416+1215 is less
still needed to obtain the lowest residual values. All objects from clear due to pattern in polarization degree dominated by noise).
our sample show prominent Fe ii emission, but the [OIII] doublet This can be an indication of relatively high viewing angle (as
and Hβ narrow component (NC) differ considerably between ob- discussed below in § 4.4), or as evidence of a polar scatterer
jects. In the case of Mrk 1044, Hβ and Hα NC intensities are sim- (§5). The evidence of a polar scatterer is especially strong in the
ilar to BCs of those lines. We reproduced relatively well the pro- case of Mrk 1044, where the polarization degree is significantly
file of [OIII] doublet using just one component - NC. The notice- higher than the one of the adjacent continuum (an increase of
able asymmetry in [OIII]5007Å is nicely reproduced by fitting ≈ 0.5%!) and the polarization angle θ undergoes a change by
the Fe ii template. In the case of SDSS J080101.41+184840.7, ≈ 90%, a typical signature of polar scattering.
Hβ and Hα the NC of those lines is around 1/3 of the BC. The
intensity of the semi broad component of [OIII] is comparable to
the NC of this line. In IRAS 04416+1215, the NC of Hβ and Hα 4.3. NLSy1 spectra decomposition in polarized light
plays the least significant role in the modeling of the full profile. In Figure 4, we show spectral decomposition of VLT Hα emis-
The most prominent, in comparison to the other two sources, sion line profiles in the polarized light. The expected equato-
asymmetry in the red side of Hα is caused by the [NII] doublet, rial placed scatterer is modeled as a single broad component,
which intensity is the highest for this source. The semi broad broader than the Hα in natural light. We plot the broad compo-
component of [OIII] is dominant, and it is the only source from nent using a red line. The case in which just a single compo-
our sample with that feature. nent is enough to model the line profile is visible in the right
panel (IRAS 04416+1215). The case of Mrk 1044 and SDSS
4.2. Polarization measurements J080101.41+184840.7 are different. Modeling the full line pro-
file using just a broad component was not possible. Thus, we
We present the spectropolarimetric measurements in Table 4. For decided to add the narrow components marked in black. The nar-
each source Pmean is0.7 1.2 1
0.6 Mrk 1044 1 Mrk 1044 Mrk 1044
0.8
0.5 archival Hβ 0.8 archival Hα this work Hα
0.4 0.6
0.3 0.6
0.4 0.4
0.2
0.1 0.2 0.2
Normalised flux
Normalised flux
Normalised flux
0 0 0
4400 4500 4600 4700 4800 4900 5000 5100 5200 6200 6300 6400 6500 6600 6700 6800 6200 6300 6400 6500 6600 6700 6800
Rest-frame wavelength (Å) Rest-frame wavelength (Å) Rest-frame wavelength (Å)
0.04 0.04 0.04
Article number, page 6 of 20
0 0 0
-0.04 -0.04 -0.04
-15000 0 15000 -15000 0 -15000 0
Δvr (km s-1) Δvr (km s-1) Δvr (km s-1)
0.5 1 1.2
SDSS J080101.41+184840.7 SDSS J080101.41+184840.7 1 SDSS J080101.41+184840.7
0.4 0.8
archival Hβ archival Hα 0.8 this work Hα
0.3 0.6
0.6
0.2 0.4 0.4
0.1 0.2 0.2
Normalised flux
Normalised flux
Normalised flux
0 0 0
4400 4500 4600 4700 4800 4900 5000 5100 5200 6200 6300 6400 6500 6600 6700 6800 6200 6300 6400 6500 6600 6700 6800
Rest-frame wavelength (Å) Rest-frame wavelength (Å) Rest-frame wavelength (Å)
0.04 0.04 0.04
0 0 0
-0.04 -0.04 -0.04
-15000 0 15000 -10000 0 10000 -15000 0
Δvr (km s-1) Δvr (km s-1) Δvr (km s-1)
0.4 1 1.2
IRAS 04416+1215 IRAS 04416+1215 1 IRAS 04416+1215
0.3 0.8
archival Hβ archival Hα 0.8 this work Hα
0.6
A&A proofs: manuscript no. aa-main
0.2 0.6
0.4 0.4
0.1 0.2 0.2
Normalised flux
Normalised flux
Normalised flux
0 0 0
4400 4500 4600 4700 4800 4900 5000 5100 5200 6200 6300 6400 6500 6600 6700 6800 6200 6300 6400 6500 6600 6700 6800
Rest-frame wavelength (Å) Rest-frame wavelength (Å) Rest-frame wavelength (Å)
0.04 0.04 0.04
0 0 0
-0.04 -0.04 -0.04
-15000 0 15000 -15000 0 -15000 0
Δvr (km s-1) Δvr (km s-1) Δvr (km s-1)
Fig. 2. Natural light spectra for Mrk 1044 (top row), SDSS J080101.41+184840.7 (middle row) and IRAS 04416+1215 (bottom row). From the top: Left and middle plot: The normalized 6dF
spectrum of Mrk 1044 from (Jones et al. 2009) with fitted Hβ and Hα part. The right plot shows the normalized unpolarized spectrum obtained with FORS/VLT. Middle and lower panels show as
follows for SDSS J080101.41+184840.7 and IRAS 04416+1215: fitted Hβ, Hα part from SDSS normalized archival spectrum and the normalized unpolarized spectrum obtained with FORS/VLT.
Data are marked in grey and the model is marked in black. The fit components are given in the following: for the Hβ (left panel) Hβ broad component and [OIII] semi broad component doublet are
red, narrow components of Hβ and [OIII] doublet are black, Hβ blue component is marked in blue, power-law component is marked in blue, Fe ii pseudocontinuum is green. For the Hα (middle and
right panel) Hα narrow component and [NII] doublet are marked in black, Hα broad component is marked as red, Hα blue component is marked in blue, power-law component is marked in blue,
[SiII] doublet and Fe ii pseudocontinuum are marked in green. Lower panels of each plot corresponds to the residuals, in radial velocity units km s−1 . All spectra are normalized to Hα intensity.Marzena Śniegowska et al.: Spectropolarimetry and spectral decomposition of high-accreting Narrow Line Seyfert 1 galaxies
Table 3. Measurements in the Hβ (upper panel) and Hα (bottom panel) region for archival spectra. Columns are as follows: (1) Object’s name, (2),
(3) and (4) report the FWHM of narrow, broad and blue components in km s−1 for Balmer lines.
NAME FWHM (Hβ) NC FWHM (Hβ) BC FWHM (Hβ) BLUE
Mrk 1044 500±3 1660±61 2550±57
SDSS J080101.41+184840.7 570±43 1680±65 2540±16
IRAS 04416+1215 490±5 1400±43 2580±2
NAME FWHM (Hα) NC FWHM (Hα) BC FWHM (Hα) BLUE
Mrk 1044 500±2 1290±5 2590±2
SDSS J080101.41+184840.7 630±18 1530±14 2990±5
IRAS 04416+1215 400±7 1300±31 2790±9
Mrk 1044 SDSS J080101.41+184840.7 IRAS 04416+1215
1000 300 1000
200
500 500
I
I
I
100
0 0 0
1.0 1.0 2
P (%)
P (%)
P (%)
0.5 0.5 1
0.0 0.0 0
300 200 200
200
100 100
(°)
(°)
(°)
100
0 0 0
6300 6400 6500 6600 6700 6800 6300 6400 6500 6600 6700 6800 6300 6400 6500 6600 6700 6800
Wavelength [Å] Wavelength [Å] Wavelength [Å]
Fig. 3. From the top: the total flux, the polarization degree, and the polarization angle. From the left: for Mrk 1044, SDSS J080101.41+184840.7,
and IRAS 04416+1215. As red we mark windows in which we calculate polarization Hα. The vertical blue dotted line marks the central wavelength
for Hα
Table 4. Measurements of polarization degree (P) and the polarization angle (θ) for our sample. Columns show the name of the source, the
averaged polarization degree from the whole spectrum, the averaged θ from the whole spectrum (θmean ), the averaged P and θ from continuum
windows, and the averaged P and θ measured from emission line.
NAME Pmean θmean Pcont θcont Pline θline
[%] [◦ ] [%] [◦ ] [%] [◦ ]
Mrk 1044 0.15±0.01 107±1 0.22±0.04 129±10 0.35±0.07 65±16
SDSS J080101.41+184840.7 0.52±0.01 115±8 0.56±0.04 116±4 0.31±0.03 114±4
IRAS 04416+1215 0.16±0.02 82±1 0.26±0.06 49±11 0.05±0.05 97±9
scatterer and we investigate those two sources using STOKES Table 5. Estimation of viewing angles for our sample using the formula
radiative transfer code in this regard (see Sec. 5.2). described in Section 4.4. From left: name of the source, estimation of
FWHM based on FORS2/PMOS spectra: the FWHM of non-polarized
broad component of Hα, the FWHM of polarized broad component of
4.4. Viewing angles Hα, and the viewing angle of the source.
To estimate viewing angles, we used equation (8) from Collin
et al. 2006, NAME FWHM(Hα) BC FWHM(Hα) BC i
non-polarized polarized [◦ ]
∆Vobs ≈ VKep [ (H/R) + sin i] 2 2 1/2
. (4) Mrk 1044 1290±20 1480±20 54±2
SDSS J080101 1530±30 2500±60 31±2
We assume that FWHM of the line in polarized light measures IRAS 04416+1215 1300±20 3810±140 4±4
the true Keplerian velocity at the distance R, that is VKep , while
the observed value of ∆Vobs is measured as FWHM of the line
in non-polarized natural light, and i is the viewing angle of the 4.5. Estimation of MBH
system towards us. H/R is the ratio of the thickness of the BLR
region to its radius and represents the random motion of BLR The reverberation mapping technique allows to determine the
clouds. The physical justification for H/R higher than 0.1 comes black hole mass via virial relationship:
from Collin et al. 2006. As the authors underlined, the illumi- Rin FWHM 2
nation which is required for the BLR implies a large opening MBH = f , (5)
angle of the BLR. The most typical H/R ratio found for the BLR G
is 1/3 (e.g. Shaw et al. 2012), as we adopt this value in further where G is the gravitational constant, f is the (dimensionless)
considerations. The results are reported in Table 5. virial factor which depends on the structure, kinematic, and the
Article number, page 7 of 20A&A proofs: manuscript no. aa-main
7 2.5 6
Mrk 1044 SDSS J080101.41+184840.7 IRAS 04416+1215
Polarised flux
Polarised flux
Polarised flux
6 2 5
5 4
4 1.5
3 3
1 2
2
1 0.5 1
0 0 0
6200 6300 6400 6500 6600 6700 6800 6200 6300 6400 6500 6600 6700 6800 6200 6300 6400 6500 6600 6700 6800
Rest-frame wavelength (Å) Rest-frame wavelength (Å) Rest-frame wavelength (Å)
0.5 0.5 0.5
0 0 0
-0.5 -0.5 -0.5
-10000 0 10000 -10000 0 10000 -15000 0
Δvr (km s-1) Δvr (km s-1) Δvr (km s-1)
Fig. 4. Fitted Hα part from the polarized spectrum obtained with FORS/VLT. Data are marked in grey and the model is marked in black. For each
fit, we consider two components of Hα: broad component (red) and narrow one (black). The blue line corresponds to the power-law continuum.
Lower panels of each plot correspond to the residuals, in radial velocity units km s−1 .
viewing angle of the BLR. Rin is the emissivity-averaged radius Table 6. Comparison of mass estimations: in column (2) mass is calcu-
of BLR (also improperly called the size of the BLR), and FWHM lated using reverberation mapping technique (Du et al. 2016), in column
(3) the estimation of black hole mass is done based on this work.
is the full width at half-maximum of the emission line profile
(for details see Panda et al. 2019b). The radius of the BLR is
determined from the observed time-lag between the continuum NAME log MBH (M ) log MBH (M )
and the emission line, FWHM is determined from the spectral RM this work
modeling. The virial factor is more puzzling since it depends on Mrk 1044 6.45+0.12
−0.13 6.05±0.13
more than one property and mostly it is assumed as a fixed factor. SDSS J080101 6.78+0.34
−0.17 6.40±0.19
With the spectropolarimetry technique, it is possible to estimate IRAS 04416+1215 6.78+0.31
−0.06 6.97±0.42
the black hole mass using the FWHM of the polarized line in the
formula. The analytical formula for the virial factor in the case Table 7. Equivalent widths of FeII (2) and Hβ (3), and RFeII (4), which
of disk-like structure of thickness H (assumed Rin ) is given is the ratio of previous columns.
by:
f = [4(sin(i)2 + (H/Rin )2 )]−1 (6) NAME EW(FeII ) EW(Hβ) BC RFeII
If we use polarized line to estimate the black hole mass, we as- [Å] [Å]
sume that we are observing the line emitting medium locating Mrk 1044 40.17±0.52 42.79±1.91 0.94±0.06
the observer on the disk-plane. Thus, sin(i) = 1 (i.e., i = 90◦ ). SDSS J080101 60.89±0.42 74.41±3.22 0.82±0.07
Including the f factor with polarized line assumptions, we IRAS 04416+1215 60.88±0.24 45.79±0.92 1.33±0.08
use the following equation to compute the black hole mass for
our sources,
Hβ region, we decided to explore our sample also in the con-
Rin FWHM 2 text of RFeII . We show in Table 7 equivalent widths of Fe ii, Hβ,
MBH = , (7)
4G and RFeII . IRAS 04416+1215 is the source with the highest RFeII
value from our sample and may be called an extreme case, with
in which Rin is the radius of BLR from previous studies and RFeII > 1.3 according to Śniegowska et al. (2018) and RFeII >
FWHM is the full width at half maximum of the broad com- 1 according to Sulentic & Marziani (2015). From the latter cri-
ponent of polarized Hα line obtained in this work. terion, Mrk 1044 is a borderline high-accreting source. We ex-
We present in Table 6 the comparison between the estimation pected relatively high values of RFeII (≈ 1) for all sources from
of black hole mass for our sample from (Du et al. 2016) using our sample, since high Eddington ratio was one of our criteria to
reverberation mapping technique (column 2), and using Eq. 7 create the sample. RFeII for the other two sources, Mrk 1044 and
(column 3). However, in Eq. 7 we still used values obtained with SDSS J080101, are not extreme Fe II emitters but they are sig-
reverberation mapping, i.e. the inner radius of the BLR. From nificantly above the average value of 0.64 and the median value
our measurements, we used FWHM of polarized Hα. After us- of 0.38 in the Shen et al. (2011a) catalog, if only the measure-
ing our method, IRAS 04416+1215 the mass estimation is higher ments with errors below 20% are considered (Śniegowska et al.
in comparison to the reverberation mapping method. In the case 2018).
of IRAS 04416+1215, for which the value of the sin(i) is the
lowest, the obtained black hole mass is higher than for the stan-
dard assumption of the virial factor equal to 1. For the other two
sources, with higher viewing angle, computed black hole mass 5. STOKES modeling
is lower than from (Du et al. 2016).
We have used the radiative transfer code STOKES (Goosmann
4.6. Determination of RFeII & Gaskell 2007; Marin 2018; Savić et al. 2018) to understand
the polarization in the Hα emitting broad-line region in these
RFeII is defined as Fe ii optical flux from 4434–4684Å range three objects. STOKES is a 3D radiative transfer code that in-
to Hβ ratio of equivalent widths. RFeII seems to be an indirect corporates the Monte Carlo approach to trace every single pho-
tracer of the Eddington ratio (e.g. Shen & Ho 2014; Marziani ton emitted from the source, of which a considerable number
et al. 2018a; Panda et al. 2018, 2019a,c; Martínez-Aldama et al. are scattered by the intervening media, including physical pro-
2021a). Since we have already modeled optical spectra around cesses such as the electron or dust scattering until they become
Article number, page 8 of 20Marzena Śniegowska et al.: Spectropolarimetry and spectral decomposition of high-accreting Narrow Line Seyfert 1 galaxies
absorbed or they escape eventually to reach the distant observer. observed data. Table 8 enlists the various input parameters used
We use the latest version (v1.2) of the code5 . in our modeling for the three objects.
5.1. Parameters of the model
In our modeling, we assume the continuum source to be point-
like located at the center and emitting isotropic unpolarized ra-
diation. The flux thus generated takes the form of a power-law
spectrum Fcont ∝ ν−α with α = 2. This assumption of the value
for the spectral index is justified from prior works (Savić et al.
2018; Jiang et al. 2021) where the focus is to reproduce the po-
larization properties around a specific emission line (in our case
Hα) and the continuum around this line. The continuum source
is surrounded by the line-emitting BLR of cylindrical geometry
which is specified by the three spatial parameters, namely (a)
the distance between the continuum source to the center of the
BLR (Rmid ), (b) the height of extension of the BLR above (or
Fig. 5. The schematic view of models with equatorial scatterer. The
below) the mid-plane joining the source and the BLR (a), and continuum source is marked in grey, the broad line region marked in
(c) the half-width of the BLR (b). These three spatial terms can blue, equatorial scattering region is marked in red. As a dashed curve
be estimated using the information of the BLR’s inner (Rin ) and between gray and black dashed lines, we marked the half-opening angle
outer (Rout ) radii which we extract from previous studies. Thus of the scattering structure.
for the inner radius of the BLR, we use the reverberation mapped
estimates compiled by (Du et al. 2016) for the three objects con-
sidered in this work, while we estimated the outer radius of the We show in Figure 5 the schematic illustrations of the model
BLR using the analytical relation by Netzer & Laor (1993) for setup for our equatorial scatterer inclusive of the continuum
the dust sublimation radius which has the following form: source, the BLR, and the equatorial scattering region.We utilize
the equatorial scatterer option, as in Jiang et al. (2021), but also
out = 0.2Lbol,46
RBLR 0.5
(8) consider the option of the polar scatterer (see e.g., Smith et al.
2005). To model the polar scatterers, we assume the scatterers
where Lbol,46 is the bolometric luminosity in units of 1046 erg to be located at the same distance from the continuum source
s−1 . The Lbol for our objects are estimated by scaling the ob- and having the same size (Rout - Rin ) analogous to the equato-
served monochromatic luminosity at 5100Å compiled in Du rial scattering region. The structure of the polar scatterers is also
et al. (2016) by a luminosity-dependent bolometric correction kept identical to their equatorial counterparts. We modulate the
factor (Netzer 2019). The remaining parameter to complete the net density of the polar scatterer until we recover a good fit for
information for the BLR is the velocity distribution for the Hα the polarized spectrum and the polarization fraction. This setup
line. This requires the value of the average velocity (vavg ) along is illustrated in Figure 6.
the x-y direction (or φ direction, here we assume the z-axis to
be aligned with the BH spin axis). We estimated it by assuming
a Keplerian velocity distribution at a given radius (here, the ra-
dius considered is the mean value between the BLR’s inner and
outer radii)
√ from the source with a pre-determined BH mass, i.e.
vavg = GM/R, where G is the gravitational constant.
For the scattering region that represents the torus, we assume
a geometry of a flared disk which is characterized by inner and
outer radii as well as a half-opening angle. The light gets scat-
tered predominantly due to free electrons (i.e. Thomson scatter-
ing) where the content of this medium is regulated by the number
density parameter. In addition to these, the velocity distribution
is again set similar to the scenario for the BLR but assuming
the appropriate inner and outer radii consistent for the torus. The
inner radius for the scattering region is set based on the R-L re-
lation from the infrared reverberation mapping (Kishimoto et al.
2007; Koshida et al. 2014). We adopt the prescription of Savić
et al. (2018) for the outer radius for this region, i.e. the location
of the outer radius of the torus is set at the location where the
outer radius of the BLR subtends a half-angle of 25o . The com-
bined information from the inner and outer radius of the scat-
tering region plus the electron number density allows us to es- Fig. 6. The schematic view of models with two scatterers: polar and
timate the optical depth of this region. We obtained a grid of equatorial. The various components are colored identical to Figure 5.
electron number densities solutions to arrive at a preferred so- The polar scattering region is marked in purple. As dashed curves be-
lution that fits best the spectral information gathered from the tween gray and black dashed lines, we marked a half-opening angle of
the scattering structures.
5
publicly available at http://astro.u-
strasbg.fr/marin/STOKES_web/index.html
Article number, page 9 of 20A&A proofs: manuscript no. aa-main
5.2. Comparison with the observed polarization properties The viewing angle for Mkn 1044 of 54◦ measured in this pa-
per comparing the natural light width and the width in the polar-
To reconstruct the observed spectrum and the corresponding po- ized light is comparable to i = 47.2+1.0
larized emission pertaining to the Hα emission line and the con- −2.5 obtained by Mallick et al.
(2018) from the joint fitting of Swift, XMM-Newton and NuS-
tinuum around the line, we generate a simulated spectrum us- TAR X-ray spectra for Mrk 1044. The viewing angle of SDSS
ing STOKES for each of the three objects. Specifically, we con- J080101 using our method is 31◦ , and we do not find any in-
sider the spectrum in natural light (Fλ ), the polarized spectrum dependent measurements of the viewing angle for this source
(p × Fλ ), the polarization fraction (p) and the polarization angle in the literature. As for IRAS 04416+1215, our determination of
(χ). We compare our simulated results obtained using the param- the viewing angle gave a very small value of 4◦ . Our formal (sys-
eterization as enlisted in Table 8 for these aforementioned four tematic) error on this value is small but we do not include here
entities to the observed results. a possible systematic error due to uncertainty in the H/R factor
In addition to the above setup, we also include a polar scat- in Equation 4. On the other hand, Tortosa et al. (2022) analyzed
tering region that helps to recover the peak distribution in the XMM-Newton and NuSTAR observations and were only able to
polarized spectra and fractional polarization for Mrk 1044 and get an upper limit for the viewing angle for this source, in the
SDSS J080101.41+184840.7. These two objects are the ones range from 25 to 45 degrees, highly model-dependent. If indeed
where we recover a value of viewing angle that is slightly the viewing angle is higher than 4◦ , that would mean that the ran-
on the higher side, i.e. ∼52o for Mrk 1044 and ∼31o for the dom motions in this source are less important, i.e. H/R in Equa-
SDSS J080101.41+184840.7. On the other hand, for IRAS tion 4 is smaller than 1/3 adopted in the current paper. We discuss
04416+1215, which is the source for which we recover the more the issue of the viewing angle of IRAS 04416+1215 in the
smallest viewing angle, i.e. ∼4o , we can reconstruct the polar- Sec. 6.3.
ized emission using the model with only the equatorial scattering Our results confirmed that the sources are high accretors. The
region. values of the black hole masses for all three sources obtained by
The polarization fraction (p%) for the three sources (Mrk us from the polarized spectra are not systematically higher than
1044, SDSS J080101.41+184840.7, and IRAS 04416+1215) in- the masses obtained earlier from reverberation mapping (see Ta-
clusive of the polar scattering region are 0.40, 0.36, and 0.3, re- ble 6.). The high accretor character of all three sources is also
spectively. For IRAS 04416+1215, the p% for the case with only supported now with their X-ray spectra analysis which revealed
the equatorial scatterer is found to be similar, i.e., 0.28. We note features characteristic for high Eddington ratio sources: steep X-
that the parameterization for the polar scattering region is unre- ray spectra/large soft X-ray excess (Tortosa et al. 2022; Mallick
strained as we lack observational pieces of evidence to break the et al. 2018; Liu et al. 2021).
degeneracy between the optical depths, the density, the location,
and the extension of these media. In addition, the half-opening
angle for the polar scatterers was assumed to be identical to their 6.1. Polarization level
counterparts in the equatorial region. With the assumption that
The overall level of polarization in the three sources is low, at
the density of the medium remains rather unaffected there is a
0.5% or less, typical for Type 1 AGNs Smith et al. (see 2004);
coupling between the opening angle and the optical depth of the
Afanasiev et al. (see 2019) Polarization in Mrk 1044 has been
scattering medium, i.e., an increase/decrease in the opening an-
previously measured by Grupe et al. (1998) giving a polariza-
gle will lead to a decrease/increase in the optical depth along
tion level of 0.52%+/−0.05 and the polarization angle 144◦ .
a given line of sight. The parameters used to model this polar
This is roughly consistent with our results for the continuum
scattering region for the two objects are tabulated in Table 9.
(although our value is even lower, 0.22%+/−0.04). Their ob-
We illustrate the performance of the recovery of the spec-
servations were performed with a 2.1-m telescope at McDon-
trum in natural and polarized light for our three objects - Mrk
ald Observatory which did not allow to study specifically the
1044, SDSS J080101.41+184840.7 and IRAS 04416+1215, in
wavelength-dependent polarization.
Figures 7, 8 and 9, respectively. For each panel in these Figures,
High-quality polarization measurements of NGC 3783 and
we provide the residuals comparing the observed and simulated
Mrk 509 show higher polarization levels, of the order of 0.8%
entities and evaluate the quality of the fit using a reduced χ2 es-
in continuum and 0.5% in Balmer lines (Lira et al. 2021). Older
timate (quoted on each panel for these Figures). For comparison
papers by Smith et al. (2004) and Smith et al. (2005) contained
purposes we show the results of our STOKES modeling for Mrk
many objects with polarization larger than 0.5%. The recent
1044 and SDSS J080101.41+184840.7 excluding a polar scatter-
sample by Capetti et al. (2021) shows the mean level of polariza-
ing region in Figures A.1 and A.2. We can notice that these mod-
tion at 0.75% (median is at 0.59%), but their sources do not pref-
els without the polar scatterers are unable to recover the “peaky”
erentially belong to the high Eddington ratio class. What seems
profiles obtained in the observed spectrum, both in natural and
to be a difference between the sources of previous spectropo-
polarized light. In addition, the polarization fraction is substan-
larimetry sources studied in previous survey such as the ones of
tially under-predicted in these models, especially at the position
Afanasiev & Popović (2015); Afanasiev et al. (2019) is the re-
marking the line-center for Hα.
quirement of an additional polar scatterer. We can hypothesise
that the occurrence of a polar scatterer could be a genuine dif-
6. Discussion and conclusions ference between sources radiating at high Eddington ratio and
sources radiating at lower Eddington ratios (Population B fol-
We performed the VLT-FORS2 spectropolarimetric observations lowing Sulentic et al. 2000b).
of three AGN which belong to the NLSy1 class and were claimed The low polarization in our sample still requires both the
to accrete at the super-Eddington rates. Our goal was to check presence of the equatorial as well as polar scatterers in two ob-
whether the Eddington rates are not overestimated by the spe- jects (Mrk 1044 and SDSS J080101.41+184840.7), in the third
cific orientation of the sources. Polarimetric measurements, even object the polar scatter is not required, but also not excluded.
for top view sources, can reveal the true range of the Keplerian High Eddington ration sources in general are expected to have
velocities of the BLR. more massive outflows which might lead to rather higher than
Article number, page 10 of 20Marzena Śniegowska et al.: Spectropolarimetry and spectral decomposition of high-accreting Narrow Line Seyfert 1 galaxies
Table 8. STOKES modeling parameters (without polar scattering region)
Broad-line region Scattering region (Equatorial)
viewing angle
Object Geometry Rin Rout vavg Geometry Rin Rout half-opening angle vavg type density optical depth
[deg.] [light days] [light days] [km s−1 ] [light days] [light days] [deg.] [km s−1 ] [cm−3 ]
Mrk 1044 52 Cylindrical 10.5 41.503 745.394 Flared-disk 57.018 89.0 35 373.266 electron 2.53×105 0.229
SDSS J080101.41+184840.7 31 Cylindrical 8.3 121.920 688.783 Flared-disk 219.287 549.947 35 283.395 electron 1.06×105 0.34
IRAS 04416+1215 4 Cylindrical 13.3 146.580 621.37 Flared-disk 276.068 661.182 35 256.727 electron 1.06×105 0.396
Table 9. STOKES modeling parameters for the polar scattering region for Mrk 1044 and SDSS J080101+184840.7
Scattering region (Polar)
Object Geometry Rin Rout half-opening angle vavg type density optical depth
[light days] [light days] [deg.] [km s−1 ] [cm−3 ]
Mrk 1044 Double-cone 57.018 89.0 35 373.266 electron 1.11×106 .1.0
5
SDSS J080101.41+184840.7 Double-cone 219.287 549.947 35 283.395 electron 6.24×10 .2.0
IRAS 04416+1215 Double-cone 276.068 661.182 35 256.727 electron 1.06×105 .0.396
lower polarization but this is not what we observe. The upper For color-coding: the viewing angle and FWHM(Hα) BC from
limits for the optical depth of the polar scatter are quite high, Table 5, RFeII from Table 7.
but apparently, the geometric location of the scatterer along the We plot the sample from Capetti et al. 2021 using values
symmetry axis does not imprint high polarization. Pline /Pcont from Table 3 in their paper (ratio of column 5 and
column 2). For our sample, we used values from Table 4 (for
Mrk 1044 Pline /Pcont = 1.59, for SDSS J080101.41+184840.7
6.2. Comparison of sources’ properties with the sample from
Pline /Pcont = 0.55, and for IRAS 04416+1215 Pline /Pcont = 0.19).
Capetti et al. 2021
For Fairall 9 we use values from (Jiang et al. 2021) analysis
We decided to compare the properties of our sample to a big- Pline /Pcont = 0.89 (Pline = 0.94, Pcont = 1.06). For calculating the
ger set of spectropolarimetric measurements of 25 AGN from Eddington ratio of those sources with the same method as for our
the same instrument carried out recently by Capetti et al. 2021 sample, we use M BH and L5100Å from Table 1 from Capetti et al.
(24 objects) and by Jiang et al. 2021 (1 object). In the paper 2021. For color-coding we used:
of Capetti et al. 2021, the whole sample is a set of 25 ob- – the viewing angle we calculated from Eq. 4, wherein for the
jects. However, the BLR polarization measurements of source VKep we put values of inter-quartile width (between 25%
J145108 were affected by a cosmic ray, as the authors men- and 75%) of the BLR in total intensity from Capetti et al.
tioned, and we did not include this source. For J140700, two 2021 (Table 3 column 8 in their paper) and for ∆Vobs , we put
observations were performed, and we chose observation with FWHM of non-polarized Hα BC (Capetti et al. 2021 (Table
higher S/N = 413. Pline from this work is the same quantity as 1 column 7 in their paper))
PBLR from Capetti et al. 2021. To illustrate this, we calculate the – FWHM of non-polarized Hα BC
window from the definition postulated in Capetti et al. 2021 for – RFeII we calculated from EW(FeII ) and EW(Hβ) BC from
Mrk 1044. FWHM(Hα) from our spectral fitting is 1290 km/s, (Shen et al. 2011b) catalog for each source from Capetti et al.
which gives 28Å. After doubling this value as in Capetti et al. 2021 sample
2021, we obtain 56Å, which is comparable with ≈50Å, taken in For Fairall 9, Jiang et al. (2021) determined the viewing an-
this work. We decided to use Pline term, since as we show for gle of the source as 50◦ , from the polarization level and polar-
our objects, they may contain a narrow line component in this ization angle fit to the data. For our sources, we determined this
part of polarized spectrum. In Figure 10 we show 3 plots of the viewing angle from a simple comparison of the FWHM in natu-
Pline /Pcont - λEdd plane. Each plot contains sample from Capetti ral and in polarized light, following Equation 4. To achieve con-
et al. 2021, Fairall 9 measurements from (Jiang et al. 2021) and sistency with our sample, and to gain insight into the accuracy
measurements for our sources. For Fairall 9 and our sources we of the viewing angle determination, we used the data from (Jiang
labelled the points. We color-coded the Pline /Pcont - λEdd with et al. 2021) and adopted our method, based on Equation 4. In this
3 quatities: the viewing angle (left panel), FWHM Hα (middle case, we obtained for the viewing angle the value ∼27◦ , which
panel) and RFeII . For calculating Pline /Pcont from our sample we is smaller compared to the original determination by Jiang et al.
use measurements from Table 4. To calculate the Eddington ratio (2021). We used both values in the plot, marking them as Fairall
of sources from our sample we use: 9 (our value) and Fairall (II) - original value. However, as Jiang
– for bolometric luminosity we use bolometric correction of et al. (2021) comment, if the object has a more complex struc-
Netzer (2019) (eq. 3 therein), which in our case (for L5100Å ) ture (i.e a warped structure), a single value of the viewing an-
is: gle may be just illusory. We compute the Eddington ratio with
the method used for the rest of the sample. We used values for
kBOL = 40 × [L5100Å /1042 ergs−1 ]−0.2 (9) L5100Å and M BH from Du et al. 2015. This value of Eddington
ratio is plotted in the middle and right panel of the Figure 10.
– for Eddington luminosity we use masses which we calculate The outlying source on the Pline /Pcont - λEdd diagram with
and present in Table 6. the highest Pline /Pcont ratio is J110538. This high ratio is caused
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