# Optical properties of two complementary samples of intermediate Seyfert
galaxies

Benedetta Dalla Barba 1,2,∗, Marco Berton 3, Luigi Foschini 2,
Giovanni La Mura 4,5, Amelia Vietri 6, and Stefano Ciroi 6

###### Abstract

We present preliminary results of the analysis of optical spectra of two
complementary samples of Seyfert galaxies. The first sample was extracted from
a selection of the 4th Fermi Gamma-ray Large Area Telescope (4FGL) catalog,
and consists of 9 $\gamma$-ray emitting jetted Seyfert galaxies. The second
one was extracted from the Swift-BAT AGN Spectroscopic Survey (BASS), and is
composed of 38 hard-X ray selected Active Galactic Nuclei (AGN). These two
samples are complementary, with the former expected to have smaller viewing
angles, while the latter may include objects with larger viewing angles. We
measured emission line ratios to investigate whether the behavior of these
Seyferts can be explained in terms of obscuration, as suggested by the well-
known Unified Model (UM) of AGN, or if there are intrinsic differences due to
the presence of jets, outflows, or the evolution. We found no indications of
intrinsic differences. The UM remains the most plausible interpretation for
these classes of objects even if some results can be challenging for this
model.

1 Università degli studi dell’Insubria, 22100 Como, Italy
2 Osservatorio Astronomico di Brera, Istituto Nazionale di Astrofisica (INAF),
23807 Merate, Italy
3 European Southern Observatory (ESO), 19001 Santiago de Chile, Chile
4 Osservatorio Astronomico di Cagliari, Istituto Nazionale di Astrofisica
(INAF), 09047 Selargius, Italy
5 Laboratório de Instrumentação e Física Experimental de Partículas (LIP),
1649-003 Lisboa, Portugal
6 Dipartimento di Fisica e Astronomia, Università di Padova, 35122 Padova,
Italy
∗ Correspondence<EMAIL_ADDRESS>
Keywords: Seyfert galaxies; intermediate Seyfert; AGN; optical spectroscopy

##  1 Introduction and motivation

Seyfert galaxies were first introduced by Seyfert in 1943 [1], where he noted
the presence of unusual broad emission lines in the optical spectra of a
sample of nearby galaxies. Following this seminal work, other authors [2],
[3], [4], [5] emphasized the need to categorize Seyferts into at least two
groups: Seyfert 1 and Seyfert 2 (Sy1s and Sy2s, respectively). Sy1 correspond
to Active Galactic Nuclei (AGN) with low inclination or nearly face-on
orientations, where the Broad-Line Region (BLR) is visible in optical spectra
and produces a broad component in the permitted lines, such as the Balmer
lines. In contrast, Sy2 have high inclination or nearly edge-on orientations.
This orientation difference is not linked to the host galaxy properties, but
to the presence of a dusty torus, which completely obscures the BLR in Sy2,
allowing only the narrow lines from the Narrow-Line Region (NLR) to be visible
in the spectra. This interpretation of Seyferts is based on the well-known
Unified Model (UM) introduced by Antonucci [6], [7], [8]. According to this
model, the distinction between Sy1 and Sy2 galaxies is sharp, although there
are numerous sources with intermediate properties observed to date. These
sources have been classified as Intermediate Seyferts (IS), defined as
Seyferts with spectral properties falling between types 1 and 2 [9], [10]. In
the spectra of IS galaxies, the Balmer lines exhibit a composite profile
composed of a narrow peak from the NLR superimposed on a broad component from
the BLR. Consequently, IS galaxies can be further subdivided based on the
prominence of the BLR relative to the NLR in the permitted lines. This
division results in categories such as Sy1.2, Sy1.5, Sy1.8, and Sy1.9, with
the sequence representing a decreasing contribution of the broad component
compared to the narrow one. The determination of an IS galaxy’s type can be
achieved using two methods, the first of which is:

$R=\frac{\mathrm{[OIII]}\lambda 5007}{\mathrm{H}\beta_{\mathrm{broad}}}$ (1)

where [O III]$\lambda$5007 is the line flux of the oxygen line, the same for
H$\beta$ [11], [12]. Here we identify four classes using Table 1:

Table 1: Seyfert classification according to Whittle [11]. Seyfert type |  Condition
---|---
Sy1 |  R $\leq$ 0.3
Sy1.2 |  0.3 < R $\leq$ 1
Sy1.5 |  1 < R $\leq$ 4
Sy1.8 |  R > 4
Sy1.9 |  Broad component only in H$\alpha$

In the second we calculate directly the type through [12]:

$\mathrm{Type}\sim
1+\bigg{[}\frac{I(H\alpha_{\mathrm{n}})}{I(H\alpha_{\mathrm{t}})}\bigg{]}^{0.4}$
(2)

using the properties of the sole H$\alpha$ line. The idea, in both the cases,
is to have a comparison between the narrow and the broad/total components of
the lines.
In the UM framework, IS are positioned between Sy1 and Sy2 galaxies,
suggesting they represent AGN with intermediate viewing angles and partial
obscuration. More recently, some authors have brought attention to the
limitations and challenges of explaining Seyfert galaxy properties solely
within the UM framework [13]. This study centered on J2118–0732, which is the
first confirmed non-local IS residing within an interacting late-type galaxy,
capable of maintaining powerful relativistic jets detectable at $\gamma$-ray
energies. The probability of observing such a $\gamma$-ray source is
exceedingly low within the UM framework. More commonly, relativistic jets are
observable in Sy1 due to the alignment of the jet with the observer’s line of
sight, leading to enhanced emissions through the beaming process. However, a
small fraction of AGN, roughly 2.8% of $\gamma$-ray-detected AGN, exhibit
misaligned jets [14]. The intriguing question that arises is why these
misaligned sources constitute such a minor portion of $\gamma$-ray-emitting
AGN. Several potential explanations can be considered: these objects may
appear as IS, but their line profiles are modified by intrinsic processes
unrelated to obscuration, or their scarcity may be due to the faintness of
$\gamma$ emissions when they are not aligned with the observer’s line of
sight, possibly as a result of instrumental limitations.
From these observations, it becomes evident that the UM alone does not provide
a sufficient explanation for the differences among Seyfert classes. One
possible approach is the introduction of a hybrid model that combines the
extrinsic representation (the UM) with more intrinsic factors such as the
presence of relativistic jets, outflows, or evolutionary processes. Jets can
significantly alter the spectral characteristics of the source, especially
when dealing with Changing-Look AGN (CL-AGN). In the most common scenario, CL-
AGN initially exhibit a spectrum dominated by emission lines, which then
transforms into a featureless spectrum, and vice versa [14]. However, this
transformation is not the only possibility. Variations in obscuration or
accretion rates can also induce changes in the optical spectrum, particularly
in the BLR of the emission lines [15], [16]. Alternatively, outflows can
reshape the profiles of lines, such as [O III]$\lambda\lambda$4959, 5007.
These oxygen lines often display blue bumps or additional components, which
are associated to outflows. This can have a significant impact on the
classification process, given that [O III]$\lambda$5007 is one of the most
critical lines for both manual and automated (machine learning) classification
methods [17]. Lastly, we can consider an evolutionary scenario in which a
peculiar class of Seyferts (Narrow-Line Seyfert 1, NLS1) can be seen as young
or rejuvenated quasars that transition into their mature phase and follow an
evolution path known as the blazar sequence [19], [18].
Another aspect related to the intrinsic hypothesis is the difficulty in
classifying Seyfert types accurately. We have previously discussed the most
common classification methods, both of which rely on distinguishing between
the emission from the BLR and the NLR. However, a significant problem arises
when the spectrum has a low Signal-to-Noise ratio (S/N), as the edges of each
emission component can overlap, potentially leading to misclassification. This
highlights the limitations of the current classification methods and
underscores the need for further research to better understand the nature of
different source classes.
Additionally, there is currently a lack of dedicated research focused on IS
and an in-depth analysis of their specific properties compared to the main
Seyfert types. Occasionally, IS objects can introduce ambiguity into samples
of AGN classified as Seyferts. For example, in a study by Berton et al. (2020)
[20] a previously classified NLS1 exhibited a characteristic IS profile for
H$\beta$, featuring a combination of broad and narrow components.
This interpretational challenge is closely tied to the quality of the spectra.
When S/N is particularly low, a faint BLR emission, typical of Sy1.8/1.9, may
mix with the continuum, giving the appearance of a Sy2. Conversely, a faint
NLR emission, typical of Sy1.2, can be overwhelmed by the broad component,
resulting in a Sy1-like line shape for the Balmer lines. To address this
ambiguity, one potential solution is to prioritize the study of high S/N
spectra, ensuring more accurate classifications. However, this approach may
necessitate excluding fainter sources or observations from small telescopes.
As always, a balance must be struck between the completeness of the samples
and the precision of our results.
This study is centered on the analysis of the optical spectra of the selected
sources, while the analysis of X-ray spectra of the same sample will be
included in a forthcoming paper. Here are presented only preliminary results
on this study, the error estimation is not present, but as a reference we
calculate it for the worst spectra in the sample (lower S/N), which is CGCG
164-019. In this case we adopted H$\beta$ as representative for the other
lines. The error is calculated producing N=100 spectra summing the original
one with a random noise proportional to the one between 5050–5150 Å and
performing again the fitting obtaining a distribution for the curve parameters
(amplitude, wavelength, width). The errors on the fitting parameters are
$\leq$10%.
The paper is organized as follows: in section 2 the Sample selection, in
section 3 the Spectral analysis, in section 4 the BPT/VO diagrams, in section
5 the Line ratios, and in section 6 the Conclusions.
In this paper we assume a standard $\Lambda$CDM cosmology with H0=73.3 km s-1
Mpc-1, $\Omega_{matter}$=0.3, and $\Omega_{vacuum}$=0.7 [21].

##  2 Sample selection

To obtain a reliable sample of Seyfert galaxies, we selected two complementary
samples using their BASS or SDSS spectra, so relying on their optical emission
properties. The first sample, obtained from the Swift-BAT AGN Spectroscopic
Survey (BASS)111http://www.bass-survey.com, comprises hard X-ray-selected
sources. The hard X-ray emission primarily originates from the corona of the
accretion flow, interacting with cold material through inverse Compton
scattering of optical/UV photons on relativistic electrons [22]. At the same
time, soft X-ray are absorbed by the neutral hydrogen, hardening the spectra.
Consequently, according to the UM, these sources tend to be biased toward high
viewing angles, such as almost edge-on orientations.
The second group of Seyfert galaxies was selected from the study by Foschini
et al. [14], which is a $\gamma$-ray selected sample from the 4th Fermi Large
Area Telescope catalog (4FGL) [23]. The $\gamma$-ray emission is closely
related to the only AGN component capable of producing these energetic
photons: the relativistic jets. In the UM framework, these objects are biased
toward small viewing angles, resembling face-on orientations, with
relativistic jets almost aligned with the observer’s line of sight. Our study
involves the objects from [14] labeled as: SEY (Seyferts), MIS (misaligned),
CL (changing-look), NLS1, and AMB (ambiguous).
These sample selections complement each other due to the different selection
methods: the first sample prefers edge-on sources, often associated with Sy2
or high-intermediate types (Sy1.8/1.9), while the second favors face-on
galaxies that can be identified as Sy1 or low-intermediate types (Sy1.2/1.5).
Following the preliminary selection (see the second row in Table 2), all
spectra were meticulously inspected one by one and categorized into partial
and complete samples. The first group includes sources where only H$\beta$ and
[O III]$\lambda\lambda$4959, 5007 were visible, while the second group
encompasses sources with a spectral range from H$\beta$ to [S
II]$\lambda\lambda$6717, 6731, which is ideal for constructing the Baldwin,
Phillips & Terlevich/Veilleux & Osterbrock (BPT/VO) diagrams [24], [25]. These
diagrams can distinguish between starburst objects and AGN, as well as
classifying AGN into Seyfert and Low-Ionization Nuclear Emission-line Region
galaxies (LINERs) [26]. They rely on the ratio between the Balmer lines and
the ionized forbidden lines, revealing the presence of a more or less powerful
radiation field, which, in turn, is related to the presence or absence of a
central emitting region, the AGN.
As an initial approach, we considered only the complete spectra, enabling us
to construct the BPT/VO diagrams. At this point, the two samples consisted of
121 sources from the BASS selection and 50 objects from the 4FGL selection.
We excluded Sy1.9 sources from the sample due to the Balmer decrement
calculation, which relies on the ratio between the broad components of both
H$\alpha$ and H$\beta$, a feature absent in Sy1.9 (only present in H$\alpha$).
Consequently, we obtained 38 AGN from BASS and 7 sources from 4FGL. A summary
of these steps is provided in Table 2.
The limited number of sources in the 4FGL sample results from our reliance on
publicly available spectra from the Sloan Digital Sky Survey (SDSS). We
supplemented our dataset with two additional objects from the 4FGL sample,
namely 1H 0323+342 and PKS 2004-447, both NLS1 sources. These spectra come
from the Boller & Chivens spectrograph mounted on the Galileo telescope in
Asiago (University of Padova) and from the FOcal Reducer/low dispersion
Spectrograph 2 (FORS2) at UT1 (Very Large Telescope, European Southern
Observatory), respectively. They were added to provide a reference for the
NLS1 group in terms of obscuration, Seyfert type, and disk luminosity. Both of
them respect the same conditions for the spectra selection, even if 1H
0323+342 show very faint [S II]$\lambda\lambda$6717, 6731 lines.

Table 2: Selection steps operated to obtain the final sample using the BASS and SDSS optical spectra of the sources. 4FGL | BASS
---|---
Number | Condition |  Number | Condition
| Total number of sources |  |
2982 | from 4FGL (with | 1210 | Total number of sources
| 1H 0323+342 and |  | from BASS
| PKS 2004-447) |  |
80 | CL + MIS + SEY | 638 | With the optical spectra
| \+ NLS1 + AMB |  |
| With at least some |  | With at least some
52 | spectral features visible | 334 | spectral features visible
| (H$\alpha$/H$\beta$/[O III]) |  | (H$\alpha$/H$\beta$/[O III])
15 | Complete spectra | 121 | Complete spectra
| (from H$\beta$ to [S II]) |  | (from H$\beta$ to [S II])
| No Sy1.9, [O I]$\lambda$6300 |  | No Sy1.9, [O I]$\lambda$6300
9 | and/or [S II]$\lambda\lambda$6717, | 38 | and/or [S II]$\lambda\lambda$6717,
| 6731 significant |  | 6731 significant
Figure 1: Comparison of the H$\beta$ fitting for Mrk 79. In the left panel, a
multi-Gaussian approach is used with a curve for the BLR emission (dashed-
violet) and a second curve for the NLR one (dotted-red). The sum of the two
components is shown in blue, and in the lower panels, the residuals in green
represent the difference between the adopted model and the spectra. In the
right panel, the same analysis is conducted with a Lorentzian. It is evident
that in the second case, the residuals are more pronounced and tend to
overestimate both the center and the wings of the line.  Figure 2: Comparison
between a Lorentzian profile (in orange) and a multi-Gaussian one (in blue,
with the narrow and broad components shown in red/violet) with a FWHM of
approximately 2000 km s-1. The lower panel illustrates the residuals as the
difference between the double Gaussian model and the Lorentzian profile.
Figure 3: From left to right: fitting of H$\beta$, [O
III]$\lambda\lambda$4959, 5007, [O I]$\lambda$6300, H$\alpha$, [N
II]$\lambda\lambda$6548, 6583, [S II]$\lambda\lambda$6717, 6731 for Mrk 110.
As in Figure 2, the blue line represents the total fitting, the red dotted
curves represent the narrow component, the violet dashed curves represent the
broad component(s), and the orange dashed curve is associated with a possible
outflow component related to the [O III] emission. These blueshifted bumps are
commonly observed in the [O III] profiles [35]. In the lower panel, the
residuals are shown in light green. In the bottom-right panel of the figure,
the [O I]$\lambda$6363 line is present as a residual because it was not fitted
during the analysis.  Figure 4: Fitting of the spectrum of 1H 0323+342 in the
H$\beta$–[O III]$\lambda\lambda$4959, 5007 region is shown in the left panel,
and of the [O I]$\lambda$6300–[S II]$\lambda\lambda$6717, 6731 region for 1H
0323+432. The line styles and colors are the same as in Figure 3.
Additionally, there is a telluric absorption feature between 6450-6500 Å.

##  3 Spectral analysis

###  3.1 Preliminary steps

The spectral analysis followed the classical initial steps, including
cosmological redshift correction, de-reddening, host galaxy subtraction, and
iron line subtraction. Each of these steps was implemented using Python 3.0
scripts, and the specific tools used in each step are detailed below.
For the cosmological redshift correction, we utilized a built-in function
available on
GitHub222https://github.com/sczesla/PyAstronomy/blob/master/src/pyasl/asl/dopplerShift.py.
The de-reddening procedure was based on the Cardelli, Clayton, Mathis (CCM)
extinction law [27] and the relationship between $A(V)$ (total extinction in
magnitudes) and $N_{H}$ (column density of neutral hydrogen atoms in cm-2)
[28]. To calculate $N_{H}$, we used the dedicated NASA-HEASARC
tool333https://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3nh/w3nh.pl, which
provides a weighted value. We then converted this value into $A(V)$ using the
proportionality factor of $5.3\times 10^{22}$ mag cm-2 mentioned earlier and
applied it within the CCM extinction law. Finally, we used:

$I_{\lambda,0}=I_{\lambda}\cdot 10^{0.4\cdot A(\lambda)}$ (3)

The host galaxy subtraction was carried out exclusively for sources with a
redshift $\lesssim$ 0.3 [29], and it was applied to spectra where the
absorption features were clearly distinguishable. We employed galaxy templates
from Mannucci et al. [30], which encompass various host galaxy types,
including elliptical and spiral galaxies, with a spectral resolution of 5 Å.
Due to the template’s inherent wavelength binning, we needed to rebin the
AGN+host spectra before conducting the subtraction. In general, we used the Ca
II lines at 3934 Å and 3968 Å as a reference to renormalize the template.
These lines are exclusively produced by the galaxy’s stellar components.
For the iron subtraction, we made use of templates from Kovacevic et al.
(2010) and Shapovalova et al. (2012) [31], [32], which are accessible on the
Serbian Virtual Observatory444http://servo.aob.rs/FeII_AGN/link6.html. These
templates span the wavelength range from 4000 Å to 5500 Å and include various
Doppler widths of Fe II lines, ranging from 700 to 2800 km s-1. We performed
iron subtraction in spectra where the peaks at 4570 Å and 5270 Å exceeded the
continuum value by at least 2$\sigma$ (standard deviation).

###  3.2 Line fitting

Similar to the iron subtraction, line fitting was performed with a detection
limit set at >2$\sigma$. This constraint contributed to the limitation in the
number of sources analyzed. This procedure proved especially useful in cases
where the considered line was very faint and closely intertwined with noise
contributions to the flux.
To construct the BPT/VO diagram, we employed a multi-Gaussian model for
fitting the following lines: H$\beta$, [O III]$\lambda\lambda$4959, 5007, [O
I]$\lambda$6300, [N II]$\lambda\lambda$6548, 6583, H$\alpha$, and [S
II]$\lambda\lambda$6717, 6731. We chose the multi-Gaussian approach due to the
composite nature of the spectra, as previously mentioned, which includes both
a broad and a narrow emission component (BLR and NLR, respectively). The
narrow peaks of H$\beta$ and H$\alpha$ occur in the same region as the
forbidden lines (e.g., [O III]). Consequently, the Full Width at Half Maximum
(FWHM) of these lines remains fixed, limiting the number of free parameters.
An alternative to Gaussian fitting is the Lorentzian approach, but it does not
adequately fit the composite profile of these spectra. Specifically, it can
overestimate the contribution of the wings when the narrow component dominates
or overestimate the peak when the broad component is prominent. A comparison
between the multi-Gaussian and Lorentzian approaches is illustrated in Figure
1. However, selecting the correct model is not always straightforward, as seen
in Figure 2, where a double Gaussian and a Lorentzian yield similar profiles
when the FWHM is approximately 2000 km s-1. This is particularly relevant
because 2000 km s-1 marks the boundary between classifying an object as a NLS1
or a Sy1. The ability to distinguish between these two classes is highly
dependent on the resolution of the analyzed spectra, as emphasized by Berton
et al. [20].
To minimize the number of free parameters in the fitting process, we imposed
some constraints. For instance, starting with the [O III]$\lambda\lambda$4959,
5007 lines, we allowed the most prominent line ([O III]$\lambda$5007) to be
fitted freely while binding all parameters (center, FWHM, and amplitude) of
the other line. The center displacement and FWHM were the same for both lines,
and the amplitude of [O III]$\lambda$4959 was fixed at one-third of the
amplitude of [O III]$\lambda$5007 [33]. Similar FWHM constraints were applied
between [O III] and the narrow component of H$\beta$ and H$\alpha$. When
dealing with the H$\alpha$+[N II] complex, constraints were applied to the
amplitudes of the two nitrogen lines and the FWHM. The amplitude ratio between
[N II]$\lambda$6548 and [N II]$\lambda$6583 was set to 1/2.96 [34] (NIST
database555https://www.nist.gov/pml/atomic-spectra-database), and the FWHM was
kept consistent with the previous cases. The same procedure was employed for
the [S II]$\lambda\lambda$6717, 6731 lines, where the FWHM was identical for
both lines. An example of the line fitting using this procedure is depicted in
Figure 3. The fitting process was conducted using a dedicated Python code
based on the Levenberg-Marquardt
algorithm666https://docs.astropy.org/en/stable/modeling/ and least squares
statistics.

###  3.3 Notes on the sample

As stated at the beginning of Section 3.2, some spectra were not considered in
the sample due to the faintness of the lines, in particular [O I]$\lambda$6300
and [S II]$\lambda\lambda$6717, 6731. The 2$\sigma$ criterion, together with
the necessity of using the same sources in all the plots, further reduced the
number of objects. In the end, we obtained: 7 AGN from the 4FGL sample with
SDSS (jetted, marked with red dots in Figures 5 and 6) and 38 sources from the
BASS one (non-jetted, marked with blue dots). We also added two jetted NLS1
from 4FGL with other optical spectra: 1H 0323+342 and PKS 2004-447 (marked
with red triangles). From the BASS sample, we noticed 3 jetted AGN: 4C +29.30,
3C 227, and 3C 234.0 (again marked with red dots); and 5 objects with parsec-
scale jets (small-jet, marked with red stars): 2MASX J04234080+0408017 [36],
Mrk 590 [37], NGC 985 [38], Mrk 110 [39], and Mrk 705 [39]. It is important to
notice that there is no superposition between the two samples.
For the spectra of 1H 0323+342, it is important to stress that this source was
monitored by one of us between 2014 and 2021 with the Asiago Galileo
telescope, and so a lot of spectra were available. To increase the S/N, we
combined some of them, applied the same steps listed before (de-reddening,
cosmological redshift, and iron subtraction), and performed the line fitting.
The result is shown in Figure 4: both [O I]$\lambda$6300 and [S
II]$\lambda\lambda$6717, 6731 are very faint. As previously mentioned, these
are preliminary results, and in the near future, one of the further steps will
be including more jetted sources due to their limited number in this analysis.

Figure 5: The first three plots are the BPT/VO diagrams for the selected
sources while the last one is from Buttiglione et al. [42] and uses the LRI.
In blue the non-jetted objects (BASS) and in red the jetted ones (mainly 4FGL,
some BASS). The jetted sources are marked with different symbols according to
their properties: the red dots are the classical jetted AGN, the red squares
are from BASS, the red stars present small jets in the parsec scale, and with
the triangles we have 1H 0323+432 and PKS 2004-447. The solid lines identify
the classification limits from Kauffmann et al. [40], the dotted lines from
Kewley et al. [41], and the dashed-dotted lines (both black and grey ones)
from Buttiglione et al. [42].

##  4 BPT/VO diagrams

The BPT/VO diagrams are employed as a tool for the separation between AGN and
starburst galaxies using the properties of different ionization lines. Many
authors have proposed slightly different classifications [40], [41], but they
always identify three regions: starbursts, Seyferts, and LINERs. These
separations can be seen in Figure 5 with the solid and dotted lines. The
dashed-dotted lines represent the limits for the sources in Buttiglione et al.
[42], which are spread in the upper-right part of the curve. In their paper,
they studied the optical properties of a homogeneous sample of the 3CR radio
sources with a redshift < 0.3. Consequently, this population characterizes the
mean characteristics of jetted sources. Of course, this dashed-dotted line
serves as a rough reference used to compare a pure jetted sample with ours.
The last plot in Figure 5, the line ratio index (LRI) vs. [O III]/H$\beta$, is
also sourced from [42] and represents a composite view of the three BPT/VO
diagrams. The y-axis aligns with the diagnostic diagrams, while the x-axis is
calculated as follows:

$\mathrm{LRI}=\frac{1}{3}\bigg{(}\log\frac{\mathrm{[NII]}}{\mathrm{H}\alpha}+\log\frac{\mathrm{[SII]}}{\mathrm{H}\alpha}+\log\frac{\mathrm{[OI]}}{\mathrm{H}\alpha}\bigg{)}$
(4)

The objective here is to create a more stable index than single-line ratios,
aimed at avoiding ambiguous classifications, as emphasized by Kewley et al.
[41]. In the LRI-$\log$([O III]/H$\beta$) plot, two distinct regions are
separated by a horizontal line at approximately 1: the high Excitation Index
(EI) and low EI regions. EI represents the overall ratio of high and low
excitation emission lines in each source. While these two populations are
clearly visible in Buttiglione et al. [42], our work primarily focuses on
sources concentrated in the high ionization region. The absence of low
ionization sources may be related to the nature of the selected sources (AGN)
and the methods used to obtain the sample, based on $\gamma$ and hard-X
emissions.
In Figure 5, it is evident that the majority of the sample falls within the
Seyfert region, with some exceptions in the LINERs region. Particularly,
objects in the LINERs area can be interpreted as AGN undergoing a transition
from an efficient to an inefficient phase in the accretion process [41],
possibly nearing the end of an activity phase.
We also explored the possibility of replacing line fluxes with peak values to
address modeling issues, but this did not prove to be a valid alternative. For
further details, please refer to Appendix B.

Figure 6: Line ratios for our sample, the markers are the same as in Figure 5.
The colored areas represent the correspondence between the two Seyfert
classification methods adopted (from [O III]/H$\beta$ and H$\alpha$
components). The grey solid lines show the separations between the Seyfert
classes. The last two plots in the right part share the same y-axis and
identify the same Seyfert sub-classes indicated between the diagrams.

##  5 Line ratios

The objective is to investigate the relationship between the intrinsic
properties of AGN and obscuration. The oxygen luminosity (L[OIII]) can be
interpreted as a proxy for the disk luminosity (Ldisk) [43], [44], [45]. Some
authors criticized this statement [46], [47] demonstrating that in some cases
also [O III] is partially obscured in type 2 AGN, so it is not totally
independent from the classification type of the source. In the optical range
the [O III] luminosity is the sole quantity that we can use as a tool for the
disk one. Expanding the analysis to the X-rays we can eliminate this problem
producing more accurate conclusions. With this aim, we are analyzing also the
X-ray part of the spectra of the same objects to obtain more solid
informations about the possible connection between the disk luminosity and the
classification (B. Dalla Barba et al., in prep).
To assess obscuration we utilize the Balmer decrement and Seyfert type as
tools. The Balmer decrement is frequently employed as a means to assess the
degree of obscuration because its intrinsic value remains relatively
insensitive to variations in gas temperature and density within the region
being examined. However, the total value, which is what we observe, includes
the effects of obscuration. For the case B recombination in HII regions, the
reference value for H$\alpha$/H$\beta$ is approximately 2.9 [48]. In contrast,
for AGN, this theoretical value can vary depending on the Seyfert type, as
emphasized by various authors [49], [50], [51]. Consequently, if we wish to
disentangle the level of obscuration from the overall H$\alpha$/H$\beta$
ratio, we require a reliable estimate of the intrinsic value. This task is
particularly challenging when the sample comprises different Seyfert types.
Focusing the study on unobscured AGN, such as Sy1, some researchers [52], [53]
have discovered a general consensus between the Balmer decrement values and
the case B recombination. The values typically fall within the 3.0-3.5 range
with a dispersion of 0.2-0.6. There are, however, some outliers with
decrements of 1.6 or 7.8. Consequently, we will consider the case B
recombination as a valid reference, though it is important to bear in mind
that the conclusions regarding the Balmer decrement should account for the
Seyfert classification. Furthermore, it is worth noting that the intrinsic
value may be higher than the commonly used value of 2.9.
The link between L[OIII] and Ldisk arises from the level of activity of the
central engine and the degree of ionization produced in the NLR. Additionally,
the stellar component of the host galaxy can contribute to the oxygen flux,
particularly when the orientations of the host and AGN are not aligned [54].
This different inclination between the host and the AGN was studied by Kinney
et al. and Nagar & Wilson concluding that the difference in viewing angle to
the central engine (using the radio jet orientation) between Sy1 and Sy2 is a
direct and independent confirmation of the UM [55], [56]. Some explanations
for this misalignment can be the warping of the accretion disk by self-
irradiation instabilities or the presence of a second gravitating system as
during a merger. This does not apply to our case because the systems in
question are not experiencing significant instabilities, except for UGC 8327
which is a dual system.
As previously mentioned, the inclusion of X-ray spectral analysis will provide
an independent measure of obscuration. This method will help us disentangle
the influence of Seyfert types on the Balmer decrement and address any host
galaxy-AGN misalignment issues previously discussed.
For the moment, we have generated three plots in Figure 6, which will be
discussed in the following subsections.

###  5.1 Seyfert type vs Balmer decrement

The initial approach was to compare the Seyfert type with the BLR Balmer
decrement (H$\alpha$/H$\beta$). H$\alpha$/H$\beta$ provides an estimate of the
extinction level because the decrement’s value is closely related to gas
physical properties. Consequently, discrepancies from the theoretical value
can be attributed to obscuration effects, even if some authors argue this
choice [49], [50], [51]. In the first panel of Figure 6, the results are
displayed: the non-jetted and jetted populations do not exhibit clear
separation. This suggests that there are no significant differences in terms
of the Balmer decrement and classification. This result poses an
interpretative challenge.
The Balmer decrement is commonly used as an indicator of obscuration, with
Sy1/ Sy1.2/Sy1.5 expected to be less obscured (lower Balmer decrements) than
Sy2/ Sy1.9/Sy1.8 (higher Balmer decrements). If the separation between
jetted/non-jetted follows the UM scheme and, consequently, corresponds to
face-on/edge-on sources (e.g., Sy1/Sy2), the Balmer decrement should display
an ascending trend in the second plot of Figure 6, which is not observed. The
results for the statistical tests are as follows: Spearman (rp=-0.10,
p-value=0.46), Pearson (rp=-0.07, p-value=0.61), and Kendall (rk=-0.06,
p-value=0.53). In this case, there is no apparent correlation.
One possible explanation for the lack of a correlation between the Balmer
decrement and the Seyfert type can be the dependence of H$\alpha$/H$\beta$
with the radiation field properties. In this context we can have an excess of
H$\alpha$ emission relative to H$\beta$ produced by an intense source of high-
energy radiation (the AGN) [57].
Another point to stress is the possible bias introduced by the lack of Sy1.9
sources in the sample, for these objects it is not possible to calculate
H$\alpha_{b}$/H$\beta_{b}$ due to the missing broad component in H$\alpha$.
From the first panel of Figure 6 we can see some objects in the Sy1.9 region,
this is due to the classification method adopted. As we will explain in
section 5.2, [O III]/H$\beta$ trace better the Seyfert type than H$\alpha$.
Finally, also the influence of the Seyfert type on the intrinsic value of the
Balmer decrement can play a role as stated in the previous section.

###  5.2 Seyfert type vs [O III]/H$\beta$

A second approach involved comparing different Seyfert classification methods.
In the second panel of Figure 6, it is evident that the two methods generally
agree. For visual clarity, we identify the colored areas as regions where the
two classification methods coincide. Almost all the data points are
distributed around the diagonal, differing by at most one Seyfert class along
both axes. In particular, objects categorized as Sy1/1.2 on the y-axis exhibit
a tendency to be "over-classified" on the x-axis, indicating higher Seyfert
types. For example, 1H 0323+342, a widely recognized NLS1, is classified as a
Sy1 based on [O III]/H$\beta$ measurements, but as a Sy1.2/1.5 when
considering H$\alpha$ properties. While the difference is not extreme, it is
notable. This implies that the H$\alpha$-based classification method may not
be as reliable as the [O III]–H$\beta$ approach, especially for lower Seyfert
types. The issue with the H$\alpha$ method may stem from challenges in
accurately fitting the H$\alpha$+[N II] complex or from intrinsic
characteristics of the BLR.
Statistical tests revealed a strong trend: Spearman (rp=0.77,
p-value=3.57$\cdot 10^{-10}$), Pearson (rp=0.85, p-value=5.77 $\cdot
10^{-14}$), and Kendall (rk=0.59, p-value=6.44 $\cdot 10^{-9}$). One important
thing to stress is the absence of separation on the horizontal axis between
jetted and non-jetted AGN. Our sample is based on 4FGL and BASS objects, so
the jetted sources are selected using high-energy part of the spectrum. With
this in mind, according to the UM, this type of jetted sources should be
biased towards lower IS types (such as Sy1.2/1.5), and non-jetted ones towards
higher types (such as Sy1.8/1.9). However, this is not visible in this plot.
Globally, [O III]/H$\beta$ show a trend with the Seyfert type (with
H$\alpha$), see the second panel in Figure 6. In the UM framework, an
increasing Seyfert type is interpreted as indicating an higher level of
obscuration (from Sy1 to Sy2). Consequently, we will adopt [O III]/H$\beta$ as
a proxy for the obscuration level. The same correspondence is not observed for
the Balmer decrement.

###  5.3 L[OIII] vs [O III]/H$\beta$

The final relationship we investigated is between oxygen luminosity (L[OIII])
and [O III]$\lambda$5007/H$\beta$. As mentioned earlier, L[OIII] serves as a
tool for intrinsic properties, particularly disk luminosity, while the
obscuration level is better estimated by [O III]/H$\beta$. The pros and cons
of these quantities were discussed in section 5, here we assume that both of
them are good tracers of the intrinsic and obscuration properties respectively
of the selected AGN.
The results from the last panel of Figure 6 reveal that these two properties
are not related. There is no significant correlation observed in this
comparison and the two populations are mixed. This suggests that Seyfert type
(and subtype) does not depend on L[OIII], and consequently, the intrinsic
properties that can be derived from the optical range of the spectrum. In
other words, this implies that the properties of the central region of the AGN
do not lead to differences in terms of Seyfert type, supporting the idea of a
common structure, as predicted by the UM. The results of the statistical tests
are as follows: Spearman (rp=0.22, p-value=0.13), Pearson (rp=0.26,
p-value=0.08), and Kendall (rk=0.15, p-value=0.15).

##  6 Conclusions and future work

We conducted a search for a possible relationship between the intrinsic
properties and the obscuration of a sample of Seyfert galaxies, categorized as
jetted and non-jetted sources. Confirming their classification as AGN,
particularly as Seyfert for the majority, was achieved through the BPT/VO
diagram. Objects situated within the LINERs region can be interpreted as aging
AGN in their final activity phase. Additionally, we created the LRI-[O
III]/H$\beta$ plot, where nearly all the data points fall within the high EI
region. This trend likely stems from the nature of the sources (AGN) and the
spectral bands used for sample selection.
Next, we generated line ratio plots with the initial aim of finding a suitable
parameter related to obscuration level. Our first attempt, using the Balmer
decrement, did not yield significant results. However, in a subsequent
analysis, we employed [O III]/H$\beta$, which provided valuable insights when
compared to Seyfert type, as expected from the UM. As for intrinsic
properties, optical spectroscopy limits us to using L[OIII], which can be an
indicator of the disk luminosity and, in turn, depends on the central engine’s
activity. A comparison of these two parameters did not reveal a significant
relationship, suggesting that the central engine’s structure in Seyfert
galaxies may be consistent across all types, with obscuration remaining the
most plausible explanation. Summarizing, we do not find any clear separation
between the jetted/non-jetted sources in terms of the considered parameters,
this finding aligns with the UM.
On the contrary, there is no clear separation between jetted and non-jetted
objects in terms of IS sub-type. In panel b of Figure 6, a clear distinction
between Sy1.2/1.5 and Sy1.8/1.9, based on $\gamma$-ray jetted/non-jetted
categorization, is not evident. This implies that jetted sources do not always
correspond to lower Seyfert sub-types (such as Sy1.2 and Sy1.5) but can
sometimes correspond to Sy1.8/1.9. This unpredicted result challenges the UM
and warrants further investigation.
One potential future step involves incorporating X-ray analysis of the same
sources to calculate the parameter NH. Combining optical and X-ray analyses
allows for the use of the low-energy part as an estimator of obscuration
(using [O III]/H$\beta$ or the Balmer decrement) and the high-energy component
as a more robust indicator of accretion power.
Another approach is to focus the analysis on specific sources, utilizing
integral field spectra. Datacubes enable the study of each region of the AGN,
providing both spectra and photometry simultaneously.

Acknowledgments: BDB is grateful to the organizers of the 14th Serbian
Conference on Spectral Line Shapes in Astrophysics (14SCSLSA) for their
partial economic support for the conference. This work has been partially
funded by INAF-OA Brera Basic Research. Funding for the Sloan Digital Sky
Survey V has been provided by the Alfred P. Sloan Foundation, the Heising-
Simons Foundation, the National Science Foundation, and the Participating
Institutions. SDSS acknowledges support and resources from the Center for
High-Performance Computing at the University of Utah. The SDSS web site is
www.sdss.org. SDSS is managed by the Astrophysical Research Consortium for the
Participating Institutions of the SDSS Collaboration, including the Carnegie
Institution for Science, Chilean National Time Allocation Committee (CNTAC)
ratified researchers, the Gotham Participation Group, Harvard University,
Heidelberg University, The Johns Hopkins University, L’Ecole polytechnique
fédérale de Lausanne (EPFL), Leibniz-Institut für Astrophysik Potsdam (AIP),
Max-Planck-Institut für Astronomie (MPIA Heidelberg), Max-Planck-Institut für
Extraterrestrische Physik (MPE), Nanjing University, National Astronomical
Observatories of China (NAOC), New Mexico State University, The Ohio State
University, Pennsylvania State University, Smithsonian Astrophysical
Observatory, Space Telescope Science Institute (STScI), the Stellar
Astrophysics Participation Group, Universidad Nacional Autónoma de México,
University of Arizona, University of Colorado Boulder, University of Illinois
at Urbana-Champaign, University of Toronto, University of Utah, University of
Virginia, Yale University, and Yunnan University.

##  Appendix A Appendix

Here, we present the results for the BPT/VO diagrams calculated using the
value of the peak of the lines, instead of the flux. This idea arises from the
potential to avoid the fitting procedure when the spectrum has a low S/N by
estimating only the value of the maximum flux for each line. The results are
depicted in Figure 7. When comparing these results with those in Figure 5, a
significant difference is evident between the two approaches. To assess this,
we calculated the average shifting of the points in both the vertical and
horizontal axes to determine if there is a common trend and if applying a
correcting factor is warranted. The average shift (in both axes) and the
standard deviation are of a similar order of magnitude, as indicated in Table
3. Consequently, we can conclude that there is no preferential shift in the
positions between the classical BPT/VO and the line-peak BPT/VO plots. The
same observation applies to the LRI-$\log$([O III]/H$\beta$) plot. This
underscores the fundamental role of the fitting procedure in producing a valid
classification of the sources.

Figure 7: BPT/VO diagram of the selected sample, utilizing the line peak instead of the flux. The line styles and symbols remain consistent with those in Figure 2, but different colors are employed to emphasize the distinct method employed here. The green points correspond to the blue ones (non-jetted), and the magenta points correspond to the red ones (jetted, small-jets, and NLS1) in the previous BPT/VO diagram (Figure 5). Table 3: Results for the horizontal and vertical shift (with standard deviation) in the BPT/VO plots as difference between the peak case (Figure 7) and the flux one (Figure 5). | 4FGL | BASS
---|---|---
Axis | Shift | $\sigma$ | Shift | $\sigma$
$\quad\quad\quad\log$([N II]/H$\alpha$) |  0.24 |  0.26 |  0.24 |  0.23
$\quad\quad\quad\log$([S II]/H$\alpha$) |  0.51 |  0.13 |  0.27 |  0.19
$\quad\quad\quad\log$([O I]/H$\alpha$) |  0.53 |  0.41 |  0.46 |  0.61
$\quad\quad\quad\log$([O III]/H$\beta$) |  0.45 |  0.49 |  0.51 |  0.35

##  Appendix B Appendix

In this section the complete list of the sources included in the study. The
list of the sources (Table 4) contains informations for the BASS sample from
Koss et al. [58]. Table 5 contain the list of the fluxes of the fitted lines,
used to produce the plots.

Table 4: The horizontal lines divide the BASS and the 4FGL. In the last row the * symbol refers to the sum of six different spectra to produce the studied one: 2016-02-06, 2016-08-12, 2016-12-08, 2016-12-29, 2017-01-20, and 2017-12-18. Continue in the next page.  IAU name |  Alias |  RA |  DEC |  z |  NH |  Telescope |  Exposure |  Date
---|---|---|---|---|---|---|---|---
|  |  [∘ J2000] |  [∘ J2000] |  [$10^{-3}$] |  [$10^{20}$ cm-2] |  |  [ks] |
BASS
J0149-5015 |  2MASX J01492228-5015073 |  27.34 |  -50.25 |  2.99 |  1.75 |  Palermo |  1.20 |  2010-11-12
J0157+4715 |  2MASX J01571097+4715588 |  29.30 |  47.27 |  4.79 |  11.0 |  Perkins |  2.40 |  2011-05-01
J0206-0017 |  Mrk 1018 |  31.57 |  -0.29 |  4.27 |  2.48 |  SDSS |  2.70 |  2000-09-25
J0214-0046 |  Mrk 590 |  33.64 |  -767.00 |  2.63 |  2.77 |  SDSS |  4.20 |  2003-01-10
J0234-0847 |  NGC 985 |  38.66 |  -8.79 |  4.30 |  3.48 |  Palomar |  0.15 |  2013-08-13
J0238-4038 |  2MASX J02384897-4038377 |  39.70 |  -40.64 |  6.13 |  2.06 |  Palomar |  0.15 |  2013-08-13
J0312+5029 |  2MASX J03120291+5029147 |  48.01 |  50.49 |  6.15 |  3.45 |  Perkins |  2.40 |  2011-06-01
J0330+0538 |  2MASX J03305218+0538253 |  52.72 |  5.64 |  4.58 |  1.16 |  UH |  1.22 |  2012-11-01
J0333+3718 |  2MASX J03331873+3718107 |  53.33 |  37.30 |  5.47 |  1.48 |  Masetti |  1.80 |  2006-10-02
J0423+0408 |  2MASX J04234080+0408017 |  65.92 |  4.13 |  4.62 |  12.6 |  Palomar |  0.15 |  2012-10-18
J0503+2300 |  LEDA 097068 |  75.74 |  23.00 |  5.81 |  23.4 |  Palermo |  1.80 |  2009-01-29
J0605-2754 |  2MASX J06054896-2754398 |  91.45 |  -27.91 |  8.98 |  2.64 |  Masetti |  1.00 |  2009-11-30
J0733+4555 |  1RXS J073308.7+455511 |  113.29 |  45.92 |  14.15 |  7.19 |  SDSS |  3.30 |  2004-11-05
J0736+5846 |  Mrk 9 |  114.24 |  58.77 |  3.99 |  4.83 |  Palomar |  0.15 |  2014-02-22
J0742+4948 |  Mrk 79 |  115.64 |  49.81 |  2.21 |  5.43 |  FAST |  0.72 |  2005-12-29
J0752+1935 |  2MASX J07521780+1935423 |  118.07 |  19.60 |  11.70 |  4.06 |  SDSS |  3.30 |  2003-10-27
J0803+0841 |  2MASX J08032736+0841523 |  120.86 |  8.70 |  4.68 |  2.87 |  SDSS |  4.50 |  2007-02-08
J0804+0506 |  Mrk 1210 |  121.02 |  5.11 |  1.35 |  3.86 |  Palomar |  0.15 |  2012-10-18
J0829+4154 |  2MASX J08294266+4154366 |  127.43 |  41.91 |  12.61 |  3.57 |  SDSS |  3.50 |  2008-02-28
J0832+3707 |  FBQS J083225.3+370736 |  128.11 |  37.13 |  9.20 |  3.27 |  SDSS |  5.10 |  2002-01-23
J0840+2949 |  4C +29.30 |  130.01 |  29.82 |  6.47 |  4.56 |  SDSS |  22.89 |  2003-10-25
J0842+0759 |  2MASX J08420557+0759253 |  130.52 |  7.99 |  13.37 |  5.76 |  SDSS |  2.24 |  2003-11-21
J0843+3549 |  2MASX J08434495+3549421 |  130.94 |  35.83 |  5.39 |  2.98 |  SDSS |  2.10 |  2003-02-02
J0904+5536 |  2MASX J09043699+5536025 |  136.15 |  55.60 |  3.72 |  2.32 |  SDSS |  9.00 |  2000-12-29
J0918+1619 |  Mrk 704 |  139.61 |  16.31 |  2.95 |  2.74 |  Palomar |  0.15 |  2015-02-17
J0923+2255 |  MCG +04-22-042 |  140.93 |  22.91 |  3.30 |  3.60 |  Palomar |  0.15 |  2015-02-17
J0925+5219 |  Mrk 110 |  141.30 |  52.29 |  3.52 |  1.27 |  SDSS |  4.80 |  2006-01-05

IAU name |  Alias |  RA |  DEC |  z |  NH |  Telescope |  Exposure |  Date
---|---|---|---|---|---|---|---|---
|  |  [∘ J2000] |  [∘ J2000] |  [$10^{-3}$] |  [$10^{20}$ cm-2] |  |  [ks] |
J0926+1245 |  Mrk 705 |  141.51 |  12.73 |  2.86 |  3.43 |  SDSS |  5.10 |  2006-12-23
J0935+2617 |  2MASX J09352707+2617093 |  143.86 |  26.29 |  12.21 |  1.56 |  SDSS |  13.02 |  2005-12-28
J0942+2341 |  CGCG 122-055 |  145.52 |  23.69 |  2.13 |  2.42 |  SDSS |  3.30 |  2005-12-31
J0945+0738 |  3C 227 |  146.94 |  7.42 |  8.60 |  2.00 |  SDSS |  2.70 |  2003-03-27
J0959+1302 |  NGC 3080 |  149.98 |  13.04 |  3.54 |  2.74 |  SDSS |  2.70 |  2004-02-20
J1001+2847 |  3C 234.0 |  150.46 |  28.79 |  18.48 |  1.62 |  SDSS |  2.52 |  2009-04-17
J1023+1951 |  NGC 3227 |  155.88 |  19.87 |  0.33 |  1.86 |  FAST |  0.33 |  2006-01-05
J1043+1105 |  SDSS J104326.47+110524.2 |  160.86 |  11.09 |  4.77 |  2.38 |  SDSS |  3.90 |  2004-04-20
J1315+4424 |  UGC 08327 |  198.82 |  44.41 |  3.55 |  1.68 |  SDSS |  1.92 |  2004-03-25
J1445+2702 |  CGCG 164-019 |  221.40 |  27.03 |  2.96 |  2.37 |  Palomar |  0.15 |  2005-02-17
J1508-0011 |  Mrk 1393 |  227.22 |  -197.00 |  5.44 |  4.64 |  SDSS |  2.70 |  2001-03-22
4FGL
J0038-0207 |  3C 17 |  9.59 |  -2.13 |  22.04 |  2.48 |  SDSS |  4.50 |  2014-12-21
J0324+3410 |  1H 0323+342 |  51.17 |  34.18 |  6.29 |  11.7 |  Asiago |  1.20 ($\times$6) |  *
J0937+5008 |  GB6 J0937+5008 |  144.30 |  50.15 |  27.55 |  1.31 |  SDSS |  3.60 |  2014-12-01
J0958+3224 |  3C 232 |  149.59 |  32.40 |  53.06 |  1.57 |  SDSS |  2.70 |  2018-01-24
J1443+5201 |  3C 303 |  220.76 |  52.03 |  14.12 |  1.66 |  SDSS |  2.10 |  2003-05-27
J1516+0015 |  PKS 1514+00 |  229.17 |  0.25 |  5.26 |  4.17 |  SDSS |  4.50 |  2000-05-25
J2007-4434 |  PKS 2004-447 |  301.98 |  -44.58 |  24.00 |  2.97 |  VLT |  0.15 |  2015-10-17
J2118+0013 |  PMN J2118+0013 |  319.57 |  0.22 |  46.28 |  5.87 |  SDSS |  3.60 |  2017-10-13
J2118-0732 |  TXS 2116-077 |  319.72 |  -7.54 |  26.01 |  7.99 |  SDSS |  2.70 |  2001-08-25

Table 5: All the terms are expressed in the logarithmic form (e.g. $\log_{10}$(H$\beta_{tot}$)), and the units are erg s-1 cm-2. Continue in the next page.  Name |  H$\beta_{tot}$ |  H$\beta_{b}$ |  [OIII]$\lambda$4959 |  [OIII]$\lambda$5007 |  [OI]$\lambda$6300 |  [NII]$\lambda$6548 |  H$\alpha_{tot}$ |  H$\alpha_{b}$ |  [NII]$\lambda$6583 |  [SII]$\lambda$6717 |  [SII]$\lambda$6731
---|---|---|---|---|---|---|---|---|---|---|---
BASS
J0149–5015 |  -11.57 |  -11.59 |  -12.20 |  -11.72 |  -12.65 |  -12.71 |  -10.77 |  -10.78 |  -12.24 |  -12.41 |  -12.40
J0157+4715 |  -10.84 |  -10.87 |  -12.14 |  -11.66 |  -15.92 |  -11.77 |  -10.35 |  -10.40 |  -12.06 |  -12.28 |  -12.43
J0206–0017 |  -11.33 |  -11.36 |  -11.89 |  -11.41 |  -12.61 |  -12.27 |  -10.51 |  -10.53 |  -11.80 |  -12.46 |  -12.50
J0214–0046 |  -11.09 |  -11.13 |  -11.33 |  -10.84 |  -12.07 |  -12.01 |  -10.49 |  -10.52 |  -11.54 |  -12.23 |  -12.23
J0234–0847 |  -10.74 |  -10.76 |  -11.32 |  -10.84 |  -11.97 |  -12.00 |  -9.79 |  -9.81 |  -11.53 |  -11.71 |  -11.71
J0238–4038 |  -10.88 |  -10.91 |  -11.59 |  -11.10 |  -12.48 |  -12.64 |  -10.33 |  -10.36 |  -12.17 |  -12.83 |  -12.82
J0312+5029 |  -10.59 |  -10.66 |  -10.95 |  -10.46 |  -11.68 |  -11.8 |  -9.74 |  -9.77 |  -11.33 |  -11.51 |  -11.58
J0330+0538 |  -10.61 |  -10.79 |  -10.59 |  -10.11 |  -11.77 |  -11.79 |  -9.94 |  -10.13 |  -11.32 |  -11.78 |  -11.79
J0333+3718 |  -10.73 |  -10.75 |  -11.21 |  -10.73 |  -12.24 |  -13.02 |  -10.24 |  -10.26 |  -11.77 |  -12.12 |  -12.14
J0423+0408 |  -11.70 |  -11.99 |  -11.08 |  -10.60 |  -12.15 |  -12.27 |  -11.20 |  -11.72 |  -11.80 |  -11.93 |  -11.97
J0503+2300 |  -10.26 |  -10.31 |  -10.88 |  -10.40 |  -11.97 |  -11.97 |  -9.58 |  -9.61 |  -11.49 |  -11.63 |  -11.63
J0605–2754 |  -11.27 |  -11.32 |  -11.59 |  -11.11 |  -12.23 |  -11.77 |  -10.48 |  -10.52 |  -12.24 |  -12.16 |  -12.24
J0733+4555 |  -10.68 |  -10.69 |  -11.94 |  -11.46 |  -13.03 |  -12.31 |  -10.18 |  -10.20 |  -11.84 |  -13.04 |  -13.03
J0736+5846 |  -10.63 |  -10.64 |  -11.56 |  -11.07 |  -12.66 |  -12.44 |  -10.11 |  -10.13 |  -11.96 |  -12.49 |  -12.53
J0742+4948 |  -10.31 |  -10.35 |  -10.76 |  -10.27 |  -11.62 |  -11.49 |  -9.67 |  -9.69 |  -11.02 |  -11.45 |  -11.51
J0752+1935 |  -10.77 |  -10.80 |  -11.21 |  -10.73 |  -11.76 |  -12.09 |  -10.19 |  -10.22 |  -11.61 |  -11.68 |  -11.78
J0803+0841 |  -11.54 |  -11.61 |  -11.78 |  -11.29 |  -12.30 |  -12.50 |  -10.55 |  -10.58 |  -12.03 |  -12.29 |  -12.34
J0804+0506 |  -11.05 |  -11.22 |  -10.49 |  -10.01 |  -10.98 |  -11.44 |  -10.38 |  -10.91 |  -10.96 |  -11.40 |  -11.32
J0829+4154 |  -11.31 |  -11.36 |  -11.73 |  -11.25 |  -12.11 |  -12.57 |  -10.58 |  -10.61 |  -12.10 |  -12.30 |  -12.26
J0832+3707 |  -10.87 |  -10.89 |  -11.58 |  -11.09 |  -12.37 |  -12.26 |  -10.33 |  -10.35 |  -11.78 |  -12.27 |  -12.27
J0840+2949 |  -12.50 |  -12.83 |  -11.93 |  -11.45 |  -12.66 |  -12.26 |  -11.73 |  -12.59 |  -11.78 |  -12.26 |  -12.37
J0842+0759 |  -11.42 |  -11.49 |  -11.49 |  -11.01 |  -11.95 |  -12.33 |  -10.48 |  -10.51 |  -11.86 |  -12.09 |  -12.22
J0843+3549 |  -12.02 |  -12.36 |  -11.48 |  -11.00 |  -11.99 |  -12.00 |  -11.29 |  -11.84 |  -11.53 |  -11.78 |  -11.83
J0904+5536 |  -11.28 |  -11.32 |  -12.01 |  -11.53 |  -12.67 |  -12.35 |  -10.59 |  -10.64 |  -11.88 |  -12.28 |  -12.34
J0918+1619 |  -10.16 |  -10.19 |  -10.94 |  -10.46 |  -12.38 |  -10.86 |  -9.45 |  -9.75 |  -11.35 |  -11.97 |  -12.05
J0923+2255 |  -10.52 |  -10.60 |  -10.98 |  -10.50 |  -12.04 |  -11.32 |  -9.97 |  -10.03 |  -11.79 |  -11.92 |  -11.95
J0925+5219 |  -11.11 |  -11.26 |  -11.14 |  -10.66 |  -11.70 |  -13.03 |  -10.24 |  -10.30 |  -12.56 |  -11.83 |  -11.87
J0926+1245 |  -10.51 |  -10.54 |  -11.07 |  -10.59 |  -12.05 |  -11.83 |  -9.97 |  -10.00 |  -11.36 |  -11.84 |  -11.85

Name |  H$\beta_{tot}$ |  H$\beta_{b}$ |  [OIII]$\lambda$4959 |  [OIII]$\lambda$5007 |  [OI]$\lambda$6300 |  [NII]$\lambda$6548 |  H$\alpha_{n}$ |  H$\alpha_{b}$ |  [NII]$\lambda$6583 |  [SII]$\lambda$6717 |  [SII]$\lambda$6731
---|---|---|---|---|---|---|---|---|---|---|---
J0935+2617 |  -11.28 |  -11.30 |  -11.99 |  -11.51 |  -12.71 |  -12.99 |  -10.45 |  -10.46 |  -12.52 |  -12.77 |  -12.86
J0942+2341 |  -11.41 |  -11.45 |  -11.85 |  -11.37 |  -12.47 |  -12.23 |  -10.77 |  -10.83 |  -11.76 |  -12.35 |  -12.34
J0945+0738 |  -11.01 |  -11.02 |  -11.75 |  -11.27 |  -12.92 |  -13.18 |  -10.35 |  -10.36 |  -12.71 |  -12.68 |  -12.77
J0959+1302 |  -11.07 |  -11.08 |  -12.45 |  -11.63 |  -16.52 |  -12.73 |  -10.57 |  -10.58 |  -12.26 |  -12.95 |  -12.88
J1001+2847 |  -11.68 |  -12.2 |  -11.21 |  -10.73 |  -12.60 |  -12.36 |  -10.98 |  -11.28 |  -11.88 |  -12.48 |  -12.47
J1023+1951 |  -10.24 |  -10.30 |  -10.50 |  -10.02 |  -11.06 |  -10.18 |  -9.58 |  -9.66 |  -10.54 |  -10.84 |  -10.85
J1043+1105 |  -11.11 |  -11.18 |  -11.50 |  -11.02 |  -12.29 |  -12.77 |  -10.37 |  -10.42 |  -12.29 |  -12.41 |  -12.48
J1315+4424 |  -11.75 |  -12.30 |  -11.90 |  -11.42 |  -11.79 |  -11.63 |  -11.03 |  -11.85 |  -11.35 |  -11.45 |  -11.54
J1445+2702 |  -11.27 |  -11.54 |  -10.94 |  -10.45 |  -12.34 |  -11.92 |  -10.73 |  -10.99 |  -11.44 |  -12.04 |  -12.04
J1508–0011 |  -11.78 |  -12.47 |  -11.18 |  -10.69 |  -12.09 |  -11.75 |  -10.98 |  -11.29 |  -11.28 |  -11.86 |  -11.86
4FGL
J0038–0207 |  -12.56 |  -12.70 |  -12.62 |  -12.14 |  -12.64 |  -13.47 |  -11.87 |  -11.95 |  -13.00 |  -13.46 |  -13.33
J0324+3410 |  -10.95 |  -10.96 |  -12.25 |  -11.76 |  -13.24 |  -12.71 |  -10.59 |  -10.61 |  -12.23 |  -12.52 |  -12.59
J0937+5008 |  -12.88 |  -12.95 |  -13.39 |  -12.91 |  -13.72 |  -14.54 |  -12.30 |  -12.39 |  -14.06 |  -13.31 |  -13.41
J0958+3224 |  -11.90 |  -12.94 |  -12.00 |  -11.52 |  -13.14 |  -13.05 |  -11.19 |  -11.21 |  -12.58 |  -12.77 |  -12.92
J1443+5201 |  -12.07 |  -12.19 |  -12.66 |  -12.18 |  -12.55 |  -13.45 |  -11.33 |  -11.39 |  -12.98 |  -13.04 |  -13.00
J1516+0015 |  -12.08 |  -12.29 |  -12.19 |  -11.71 |  -12.00 |  -12.47 |  -11.39 |  -11.50 |  -12.00 |  -12.18 |  -12.26
J2007–4434 |  -12.65 |  -12.70 |  -12.96 |  -12.48 |  -13.35 |  -13.54 |  -11.89 |  -11.94 |  -13.07 |  -13.12 |  -13.16
J2118+0013 |  -12.77 |  -12.81 |  -13.07 |  -12.59 |  -13.70 |  -14.00 |  -12.16 |  -12.21 |  -13.53 |  -13.67 |  -13.78
J2118–0732 |  -12.60 |  -12.91 |  -13.48 |  -13.00 |  -13.20 |  -13.45 |  -11.82 |  -12.02 |  -12.98 |  -13.94 |  -13.99

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