smaller

Jet-induced feedback in the [O III] lines of early evolution stage active
galactic nuclei
Marco Bertona,b,111email<EMAIL_ADDRESS>& Emilia
<EMAIL_ADDRESS>
a Finnish Centre for Astronomy with ESO - University of Turku, Vesilinnantie
5, FIN-20500, Turku (Finland)
b Aalto University Metsähovi Radio Observatory, Metsähovintie 114, FIN-02450,
Kylmälä (Finland)
c European Space Agency (ESA), European Space Astronomy Centre (ESAC), Camino
Bajo del Castillo s/n, 28692 Villanueva de la Cañada, Madrid, Spain

###### Abstract

It is well known that active galactic nuclei (AGN) show various forms of
interaction with their host galaxy, in a number of phenomena generally called
AGN feedback. In particular, the relativistic plasma jets launched by a
fraction of AGN can strongly affect their environment. We present here a study
of the [O III] $\lambda\lambda$4959,5007 lines in a diverse sample of early
evolution stage AGN – specifically narrow-line Seyfert 1 galaxies. Radio
imaging observations of all of the sources enable a division to jetted and
non-jetted sources, and exploiting this we show that the ionized gas
properties are significantly influenced by the presence of the jets, as we
often find the [O III] lines (blue-)shifted with respect to their restframe
wavelength. We also show how the radio morphology and the radio spectral index
do not seem to play a role in the origin of the [O III] shifts, thus
suggesting that the source inclination is not relevant to the lines
displacement. We do not find a strong relation between the [O III] line
properties and the bolometric luminosity, suggesting that within our sample
radiatively driven outflows do not seem to have a significant contribution to
the [O III] line kinematics. We finally suggest that [O III] shifts may be a
good proxy to identify the presence of relativistic jets. Additional studies,
especially with integral-field spectroscopy, will provide a deeper insight
into the relation between jets and their environment in early evolution stage
AGN.

Keywords: Active galactic nuclei; relativistic jets; AGN feedback

##  1 Introduction

Active galactic nuclei (AGN) play an essential role in the evolution of
galaxies across cosmic time (Taylor & Kobayashi, 2015), as the central black
hole (BH) and its host galaxy undergo co-evolution (Kormendy & Ho, 2013). The
nucleus can experience multiple activity phases throughout its life cycle, and
during every phase there is interplay between the AGN and its host galaxy. The
latter regulates the gas supply to the nucleus, enabling or hampering the
triggering of the AGN, and inducing changes to its activity state. On the
other hand, the AGN also affects its environment. The radiation produced by
the process of accretion onto the BH can heat or ionize the gas up to several
kiloparsecs (kpc) away (e.g., Schmitt et al., 2001; Congiu et al., 2017b;
Venturi et al., 2018), and by means of outflows or relativistic jets it can
suppress or in some cases enhance (Fragile et al., 2017) the star formation
activity across the host galaxy (Morganti, 2017). All the examples of mutual
interaction between an AGN and its host galaxy are collectively known as AGN
feedback.
Feedback is still a relatively poorly understood aspect of galaxy evolution,
and it is particularly interesting to study if and how this phenomenon occurs
in AGN in early phases of their life cycle, that is when they are ”young”.
Carrying out such research at high redshift, when all sources were young, is
not an easy task. An alternative approach is to focus on sources that are the
low redshift analogs to the first quasars. Narrow-line Seyfert 1 (NLS1)
galaxies are in this sense ideal laboratories to study the early phases of AGN
feedback since, like high-redshift quasars, they are characterized by fast
growing black holes with high Eddington ratios, and also by high metallicity,
especially prominent Fe II multiplets (Mathur, 2000).
NLS1s are mostly non-jetted sources, but some of them do actually harbor
powerful relativistic jets (Komossa et al., 2006; Yuan et al., 2008; Abdo et
al., 2009; Foschini, 2011; Foschini et al., 2015). Typically hosted by spiral
galaxies (Krongold et al., 2001; Crenshaw et al., 2003; Deo et al., 2006; Ohta
et al., 2007; Antón et al., 2008; Orban de Xivry et al., 2011; Mathur et al.,
2012; Kotilainen et al., 2016; Olguín-Iglesias et al., 2017; Järvelä et al.,
2018; Berton et al., 2019; Olguín-Iglesias et al., 2019), due to their low-
mass BHs and high Eddington ratios NLS1s are believed to be the progenitors of
other AGN classes, such as broad-line Seyfert 1 galaxies and flat-spectrum
radio quasars (Grupe, 2000; Mathur, 2000; Sulentic et al., 2000; Mathur et
al., 2001; Peterson, 2011; Berton et al., 2016a; Paliya, 2019; Berton et al.,
2020a; Foschini, 2020). Since their classification is exclusively based on the
width of the H$\beta$ line (Osterbrock & Pogge, 1985), NLS1s can currently be
found only up to z$\sim$1 (Rakshit et al., 2017), although unclassified
objects with similar properties are present also at higher redshift (Sulentic
et al., 2002; Marziani et al., 2018b). A feature which is often observed in
the optical spectra of this class of AGN is a shift, usually toward lower
wavelength, in the emission lines of [O III]$\lambda\lambda$4959, 5007,
corresponding to a velocity $|v|>150$ km s-1 (Komossa et al., 2008). Such
sources, in the literature, are called blue outliers (Zamanov et al., 2002;
Marziani et al., 2003; Komossa et al., 2008). They are mainly found among
NLS1s but, more in general, they are observed almost exclusively in population
A sources on the quasar main sequence (MS). The MS is the locus on the plane
defined by the flux ratio between the Fe II multiplets and H$\beta$, known as
R4570, and the FWHM(H$\beta$), where all type 1 AGN lie. This sequence was
identified by means of principal component analysis, and it seems to be mainly
driven by the Eddington ratio (Boroson & Green, 1992). However, some other
factors seem to play a role in it, such as metallicity, and the inclination
compared to the line of sight (Shen & Ho, 2014; Panda et al., 2019). Sources
on the MS can be divided into two main populations according to the full width
at half maximum (FWHM) of their H$\beta$ line. Population A sources are
characterized by FWHM $<$4000 km s-1, and they often exhibit prominent Fe II
emission (Sulentic et al., 2000; Marziani et al., 2001; Sulentic et al., 2002;
Marziani et al., 2003, 2018a; Czerny et al., 2018). The origin of blue
outliers is still unclear (Marziani et al., 2016a; Komossa et al., 2018; Ganci
et al., 2019), and it could be connected to the high Eddington ratio observed
in population A sources and NLS1s (Boroson, 2002; Grupe et al., 2004, 2010; Xu
et al., 2012). However, their presence is often associated with strong radio
emission indicative of relativistic jets (Komossa et al., 2018; Berton et al.,
2016b, hereafter B16). If this is true, blue outliers may be a remarkable
example of feedback between kinematically young relativistic jets, as those
present in NLS1s, and the ionized gas of their host galaxy.
In this work, we analyzed the properties of the [O III] lines, with particular
attention to blue outliers, in a sample of NLS1s whose radio morphology has
been observed with the Karl G. Jansky Very Large Array (JVLA), to investigate
their relation with radio emission. In Section 2 we describe our sample
selection and the data analysis process, in Section 3 we present our results,
in Section 4 we discuss our findings, and in Section 5 we present a brief
summary of this work. Throughout the paper, we adopt the standard $\Lambda$CDM
cosmology, with a Hubble constant H${}_{0}=70$ km s-1 Mpc-1, and
$\Omega_{\Lambda}=0.73$ (Komatsu et al., 2011).

##  2 Sample selection and data analysis

Table 1: Observational details of the spectra observed with the Asiago 1.22m telescope. (1) Short name; (2) exposure time (seconds); (3) resolution; (4) observing date. The asterisk indicates that the spectra were observed on five separate dates, 2013-12-06, 2014-09-21, 2015-08-12, 2015-11-15, 2018-07-07, and later combined (Tripathi et al., 2020). This operation is possible because the interval between these spectra is shorter than the typical scale of the NLR. Short name | Exp. Time | $R$ | Date
---|---|---|---
J0006+2012 | 25500 | 700 | 2013-2018*
J0347+0105 | 4800 | 700 | 2015-11-16
J0629$-$0545 | 6000 | 700 | 2015-12-10
J2242+2943 | 18000 | 700 | 2015-08-06

Our sample includes the 74 NLS1s that were observed with the JVLA at 5 GHz in
its most extended A configuration (longest baseline 36 km, resolution 0.5”) in
the survey described in Berton et al. (2018, hereafter B18). We also decided
to include into our sample seven additional NLS1s studied in the same JVLA
configuration. These peculiar objects show mJy-level or lower flux density at
5 GHz, but they were observed flaring at Jy-level at 37 GHz, indicating that
they harbor relativistic jets (Lähteenmäki et al., 2018). The reason for the
weak radio emission at low frequencies may be a form of absorption, either
free-free or synchrotron self-absorption (Berton et al., 2020b).

We were able to retrieve an optical spectrum for 80 out of 81 sources (all but
J0100-0200). In another object, J1348+2622, the [O III] lines are not visible
in the optical spectrum due to its high redshift ($z$ = 0.917), and it was
therefore excluded from this analysis. For 62 of these NLS1s, a measurement of
the [O III] line properties was already available in B16. For the remaining 17
objects, 13 spectra were obtained from the Sloan Digital Sky Survey (SDSS)
Data Release 16 (DR16) (Blanton et al., 2017), while the remaining four were
observed with the 1.22m telescope of the Asiago Astrophysical Observatory. The
details of these new observations are reported in Table 1. The measurement
process we used is the same as the one described in detail in B16. In the
following we provide a brief summary.
The first challenge was obtaining a precise estimate of the redshift of each
source. Ideally, stellar absorption lines would be the best proxy for the
recessional velocity of a galaxy, but in type 1 AGN they are typically not
visible. Emission lines are easily detectable, but not all of them are
reliable redshift indicators. Broad lines and high-ionization lines often
originate from outflows and appear shifted with respect to the galaxy frame,
so the ideal solution is to use low-ionization lines, such as [S
II]$\lambda\lambda$6717, 6731, [O I] $\lambda$6300, or [O II]$\lambda$3727
(Komossa et al., 2008; Sulentic et al., 2007). It is worth noting that the [O
II] line is an unresolved doublet in low-resolution spectra, therefore the
other lines, if available, grant a better estimate of the redshift. However,
due to the redshift, [O II] is often the only available reference. We measured
the wavelength of the reference line by using a fitting procedure in Python.
In this procedure there are two sources of systematic error. The first one is
due to wavelength calibration errors. For SDSS its value is 2 km s-1
(Abazajian et al., 2009), while for the other spectra it is of the order of 20
km s-1 (B16). The other error is produced by the fit of the lines, and we
computed it using a Monte Carlo method. We added a Gaussian noise, whose
amplitude was proportional to the noise in the AGN continuum measured around
$\lambda$5100, to the line. We then repeated the fitting procedure 100 times
with different noise, getting a median value of redshift and a standard
deviation. On average, the latter is rather small, $\sim$2 km s-1 as it is
obtained by a simple fit with one Gaussian of a relatively strong emission
line.

Figure 1: [O III] line fitting of the source J1114+3241. The black solid line
represents the restframe spectrum after continuum and Fe II subtraction, the
green dashed line represents the [O III] core components, the orange dashed
line represents the wing components, and the best fit is shown as the red
dashed line. The blue solid line shows the residuals, and the black horizontal
dashed line corresponds to a zero-level flux.

After moving the spectra to restframe, we subtracted the AGN continuum using
IRAF and, when necessary, we removed the Fe II multiplets using an online
software333http://servo.aob.rs/FeII_AGN/(Kovačević et al., 2010; Shapovalova
et al., 2012). At this point, we decomposed each of the [O III] lines using
two Gaussian components, one to reproduce the line core, and one for the blue-
shifted wing of the line. The line core component originates from the ionized
gas of the narrow-line region (NLR). Its kinematics is dominated by the
gravitational potential of the bulge, and therefore it usually has the same
recessional velocity as the host galaxy. The wing component, instead, probably
comes from a turbulent outflow in the inner part of the NLR. Since in type 1
AGN we preferentially have a pole-on view of the central engine, the outflow
appears as a blue-shifted (approaching) component in the line with a FWHM
larger than that of the core due to its turbulence. Our aim was obtaining the
position of the [O III] core component to derive the velocity of the ionized
gas with respect to the galaxy frame, the position of the line wing to measure
the velocity of the outflow, and the FWHM of both components to derive the gas
velocity dispersion. To reduce the number of free parameters in the fitting
procedure, we fixed the ratio between each component of the $\lambda$5007 and
$\lambda$4959 lines to its theoretical ratio of 3.0 (Dimitrijević et al.,
2007). Furthermore, we forced the FWHM of the line core, the FWHM of the wing,
and the wing velocity, expressed in km s-1, to be the same in both lines. The
errors were computed adopting the previously described Monte Carlo procedure.
In conclusion, for each line we measured four parameters: the core shift with
respect to the laboratory wavelength of [O III] (5006.843Å), the line core
FWHM, the wing velocity with respect to the core component, and the wing FWHM.
An example of a fit is shown in Fig. 1. We successfully reproduced the [O III]
lines in 76 sources, while the remaining three did not show any oxygen in
their spectra (see Discussion). The results of our calculation are shown in
Table LABEL:tab:oiii. In the following, we also considered some additional
physical properties of the sources such as their bolometric luminosity, radio
luminosity, and radio loudness parameter. The bolometric luminosity was
estimated using its scaling relation with the H$\beta$ luminosity, as in
(Berton et al., 2015). The radio luminosities are derived from B18 and (Berton
et al., 2020b). The radio loudness parameter is defined as the ratio between
the radio flux density at 5 GHz and the flux density in the optical B band.
When the ratio is larger than 10, the source is radio loud. Conversely, it is
radio quiet (Kellermann et al., 1989). To estimate it, we either used the
values published from (Foschini et al., 2015; Berton et al., 2015), or we
calculated it using the B-band magnitude values reported in the NASA
Extragalactic Database (NED) and our own measurements at 5 GHz.

###  2.1 Correlations with the whole sample

Table 2: Results of the correlations between the [O III] core shift and the physical properties of the sources. Columns: (1) First tested sample; (2) second tested sample; (3) Pearson correlation coefficient; (4) p-value of the Pearson coefficient; (5) Spearman’s rank correlation coefficient; (6) p-value of the Spearman coefficient. Q1 | Q2 | rP | p-value(P) | rS | p-value(S)
---|---|---|---|---|---
L${}_{\rm[O~{}III]}$ | vc | -0.05 | 0.68 | 0.10 | 0.38
L${}_{\rm[O~{}III]}$ | —vc— | 0.07 | 0.54 | 0.24 | 0.03
Lbol | vc | -0.10 | 0.37 | 0.01 | 0.97
Lbol | —vc— | 0.15 | 0.18 | 0.28 | 0.01
L${}_{\rm 5~{}GHz}$ | vc | 0.01 | 0.97 | 0.01 | 0.92
L${}_{\rm 5~{}GHz}$ | —vc— | 0.03 | 0.78 | 0.39 | 0.0004
RL | vc | 0.18 | 0.12 | 0.03 | 0.81
RL | —vc— | 0.12 | 0.29 | 0.36 | 0.001

The property we are most interested in is the shift of the whole [O III] line,
that is the blue outlier phenomenon. Our aim is to have a better understanding
of what may be causing this shift. We tested the core velocity against four
physical properties that may be associated with it, that is the actual [O III]
line luminosity, the bolometric luminosity, the radio luminosity, and the
radio loudness parameter. The first two, indeed, have been found to possibly
affect the [O III] core shift especially at high redshift (Marziani et al.,
2016b). The radio luminosity and the radio loudness parameter, instead, may be
related with the presence of relativistic jets (Padovani, 2017; Järvelä et
al., 2017, with some caveats, see e.g.,). The results of these tests are
reported in Table 2. None of the tested quantities is strongly correlated to
the velocity of the [O III] core. This is true also when the absolute value of
the shift is considered, neglecting whether the line is blueshifted or
redshifted. An example is shown in Fig. 2. Since it is known that [O III]
shift can also be radiatively driven in high luminosity sources (Marziani et
al., 2016b), and a weak, but significant correlation is identified by the
Spearman’s rank correlation coefficient between Lbol and —vc—, we wanted to
investigate whether an Lbol threshold exists above which a source has a higher
possibility of showing an [O III] shift or a higher FWHM([O III]). As seen in
the left panel of Fig. 2, however, no clear threshold seems to be present to
produce an [O III] shift. In the FWHM distribution, the highest values of FWHM
are found mostly, but not exclusively, among the highest luminosity sources,
although the statistic is too poor to draw any meaningful conclusion.
It is also worth mentioning that the Spearman’s rank correlation coefficient
has a rather low p-value when the radio loudness parameter and, even more
noteworthy, the radio luminosity are tested against the absolute value of the
core shift velocity, suggesting that these two weak correlations may be
significant. The radio luminosity, however, contains some dependence on the
redshift. Indeed, B18 already identified a correlation between the radio
luminosity and the redshift, present also in our sample (rP = 0.51,
p-value=3$\times 10^{-6}$), which is due to the sample selection. For this
reason, we carried out a partial correlation analysis, using the redshift as a
control parameter. In this case, the Spearman’s rank correlation coefficient
becomes 0.26 (p-value = 0.03), so the correlation seems to disappear. We
carried out the same analysis on the other correlations involving the radio
luminosity shown in Table 2, but we found no statistically significant result.
The radio loudness parameter instead is redshift independent, and it also
shows a weak but statistically significant correlation with the [O III] core
shift. As we already mentioned, radio loudness is often used as a proxy to
detect the presence of relativistic jets. This, therefore, suggests that
sources harboring relativistic jets may tend to present [O III] core shift
more often than those without a jet. To verify this hypothesis, we decided to
split the sample of sources according to the presence or not of a relativistic
jet and of other radio properties derived in B18, and to carry out some
statistical tests to verify if and how these features affect the optical
lines’ profile.

Figure 2: Left: [O III] core velocity shift in km s-1 as a function of the
bolometric luminosity in erg s-1. The horizontal dashed lines indicate the
limits for red or blue outliers at 150 km s-1. Right: FWHM([O III]) core
component in km s-1 as a function of the bolometric luminosity in erg s-1.

###  2.2 Radio classification

The first radio-based classification, jetted (J) or non-jetted (NJ), is based
on the data from B18, where each source was analyzed individually. In general,
the classification criteria were the radio luminosity, which is typically
higher than 1040 erg s-1in jetted NLS1s, the spectral index, since a flat
radio spectrum with $\alpha_{\nu}<0.5$ indicates the presence of relativistic
beamed jets, and the multiwavelength properties of the sources. In previous
studies as B16, each source could be classified only according to the radio-
loudness parameter (Kellermann et al., 1989), defined as the ratio between the
flux density at 5 GHz and in the optical B band. Sources with radio-loudness
larger than 10 are classified as radio-loud, conversely as radio-quiet.
Objects with no previous detection in radio surveys are often labeled as
radio-silent. Usually, radio-loud sources harbor relativistic jets, while the
origin of radio emission in radio-quiet objects is strongly debated (e.g.,
Panessa et al., 2019). However, several sources do not obey this rule,
therefore a classification criterion based on physical properties is needed
(Padovani, 2017). Based on the radio observations, B18 established whether
relativistic jets are present or not in each source. An example is J1038+4227,
which was classified as radio-loud by Foschini et al. (2015) (and so,
estimated to be jetted), but was later found to not have relativistic jets,
since the radio emission originates from prominent star forming activity
(Caccianiga et al., 2015).
The second subdivision is based on the radio morphology of our NLS1s. B18
classified them with a compactness criterion. They used the ratio
$\mathcal{R}=S_{p}/S_{i}$, that is the ratio between the peak flux density and
the integrated flux density of the whole emitting region. Sources with
$\mathcal{R}\geq 0.95$ were classified as compact (C), those with
$0.75\leq\mathcal{R}<0.95$ were labeled as intermediate (I), and finally those
with $\mathcal{R}<0.75$ were classified as extended (E). While jetted sources
were mostly (but no exclusively) compact, non-jetted sources typically had an
extended morphology. The morphological parameters were calculated using the
flux densities reported in B18 and (Berton et al., 2020b). For the sources
without a detection in (Berton et al., 2020b), we assumed a compact
morphology. It is worth noting that since the radio observations were carried
out in the most extended configuration of JVLA, some large-scale extended
emission may had been resolved out. However, a sample of NLS1s were
investigated at the same frequency by Chen et al. (2020) using the JVLA in a
more compact configuration. Their results seem to indicate that the morphology
of jetted sources is typically compact, with no signs of large-scale radio
lobes as those seen in other classes of jetted AGN. Furthermore, the spatial
scale of these structures even in A configuration would be several kpc and
possibly outside of the host galaxy.
The last classification was used only for jetted sources, and is based on the
division to sources with flat (F) or steep (S) radio spectral index
$\alpha_{\nu}$ ($F_{\nu}\propto\nu^{-\alpha_{\nu}}$). As in other AGN, in
jetted NLS1s the radio spectral index is a rather good indicator of the source
inclination, that is the angle between the line of sight and the relativistic
jet axis. In most AGN the spectrum of the radio core appears flat due to the
$\tau$=1 surface, where $\tau$ is the optical depth, moving closer to the BH
as a function of frequency (Falcke & Biermann, 1995). When our line of sight
is close to the axis of the system, the emission of the radio core experiences
relativistic beaming, that can result in the radio core emission to dominate
the whole spectrum. If we are seeing the radio core at larger angles, the
relativistic effects diminish and the contribution of extended, often
optically thin, emission becomes more important. A properly formed, powerful
jet is often characterized by several components, launched at different times,
that each have a convex spectrum, that is synchrotron self-absorbed below the
turnover frequency, $\nu_{\textrm{t}}$, and shows an inverted spectral index
($\alpha_{\nu}\sim$ -2.5), and optically thin with a spectral index of
$\alpha_{\nu}\sim$ 0.7 above $\nu_{\textrm{t}}$. As a component ages it loses
energy and expands, and the peak moves to lower frequencies. A superposition
of several components with different ages results in a seemingly flat, but
variable, radio spectrum (Blandford & Königl, 1979). This is especially
pronounced in sources whose jet axis is close to our line of sight, since the
emission is enhanced by relativistic beaming, and in these sources the beamed
jet and radio core emission dominates the whole radio spectrum. In sources
that we observe at larger angles, the beaming effects become less significant,
and whereas the radio core still appears flat, the overall radio spectrum
becomes dominated by extended, optically thin, emission, with a steeper
spectral index. The spectral classification used in this work is derived from
B18 and Järvelä et al. (in prep.). All the classes are reported in Table
LABEL:tab:sample.

###  2.3 Statistical tests

We studied the properties of the [O III] lines by grouping them according to
the radio properties as described in the previous section. For each
distribution, we estimated its median value, its standard deviation, and its
interquartile range (IQR). We remark that the standard deviation is not used
here as an estimate of the error associated to the median, but instead as a
proxy of how much each distribution is concentrated around its median value.
Furthermore, to examine whether two distributions are statistically different
or not, we used both the Kolmogorov-Smirnov (K-S) and Anderson-Darling (A-D)
test. The A-D test is more powerful than the K-S in most cases (Hou et al.,
2009), but we decided to adopt both tests to provide results comparable to the
literature and B16. The null hypothesis of these tests is that the two
distributions originate from the same population of sources. In the following,
we will reject it only if p-value$\leq$0.05. Finally, to account for the
measurement errors, in addition to the observed samples we created simulated
samples of sources. In each distribution, to each measurement we added a
Gaussian noise whose amplitude was equal to the measurement error, thus
creating new samples with the same size as the observed ones. We repeated this
operation 100 times and performed every time both statistical tests, to obtain
a median p-value for both K-S and A-D test. The goal of these simulations is
to verify the stability of the p-values we measured on the observed samples.
Therefore, we also estimated the standard deviation of the distributions of
p-values obtained from the simulations.

###  2.4 Jets and [O III]

####  2.4.1 Line core component

Figure 3: Left: distributions of the [O III] core velocity shift in km s-1, in non-jetted (top panel) and jetted sources (bottom panel). The vertical dashed lines indicate the limits for red or blue outliers at 150 km s-1. Binning of 25 km s-1. Right: distributions of the FWHM([O III]) core component in km s-1, in non-jetted (top panel) and jetted sources (bottom panel). Binning of 25 km s-1. Table 3: Results of the statistical tests for the spectral index samples. Columns: (1) Tested sample; (2) median velocity; (3) standard deviation of the velocity; (4) interquartile range; (5) p-value of the K-S tests against the other sample, observed in top row, simulated in middle row, and standard deviation of the simulated p-value in the bottom row; (6) p-value of the A-D tests against the other sample, observed in top row, simulated in middle row, and standard deviation of the simulated p-value in the bottom row. Sample |  Median (km s-1) |  Std (km s-1) |  IQR (km s-1) |  K-Sobs K-Ssim K-Sstd |  A-Dobs A-Dsim A-Dstd
---|---|---|---|---|---
vc(NJ) | -26.5 | 48.5 | 71.0 |  0.06 0.06 0.02 |  0.01 0.02 0.01
vc(J) | -11.9 | 226.4 | 249.1 |  0.06 0.06 0.02 |  0.01 0.02 0.01
FWHMc(NJ) | 329.4 | 85.9 | 69.1 |  0.05 0.03 0.02 |  0.03 0.02 0.01
FWHMc(J) | 385.7 | 289.8 | 328.3 |  0.05 0.03 0.02 |  0.03 0.02 0.01
vw(NJ) | -268.9 | 127.7 | 98.8 |  0.06 0.10 0.09 |  0.05 0.07 0.05
vw(J) | -335.7 | 244.1 | 298.6 |  0.06 0.10 0.09 |  0.05 0.07 0.05
FWHMw(NJ) | 868.6 | 346.2 | 311.4 |  0.03 0.06 0.05 |  0.07 0.07 0.04
FWHMw(J) | 1180.3 | 427.8 | 607.7 |  0.03 0.06 0.05 |  0.07 0.07 0.04

First of all, we examined the properties of the [O III] lines separately in
jetted and non-jetted sources, as indicated in the second to last column of
Table LABEL:tab:sample. We first compared the velocity of the core component
of [O III] with respect to the galaxy rest frame in the two samples. We remark
that, by definition, a source is considered a red or blue outlier when the
velocity of the line core component is $|v_{c}|>150$ km s-1, following the
definition by Komossa et al. (2008). The distributions are shown in the left
panel of Fig. 3. Among our 28 non-jetted sources, we found only one blue
outlier ($\sim$4%) and no red outliers. In our sample of 48 jetted sources,
instead, we found 15 blue outliers ($\sim$31%) and 3 red outliers ($\sim$6%).
The statistical properties of the two distributions are shown in Table 3. It
is evident that while the median values of the two samples are similar, both
the standard deviation and the IQR are very different. Both quantities are
much larger in jetted sources, suggesting that the [O III] core velocity
distribution in this subsample is significantly more spread out. Despite this
difference, the K-S test indicates that the null hypothesis cannot be
rejected, but the A-D test provides the opposite result (p-value = 0.01). We
repeated the tests on the simulated samples, and the results did not change
significantly. However, since our goal is to understand the nature of blue
outliers, and they are concentrated in the distributions’ tail, the A-D test
is more reliable than the K-S test. We can therefore conclude that the null
hypothesis can be rejected, and that the core shift distributions do not
originate from the same population of sources.
We then tested the distributions of FWHM of the [O III] core component in the
two samples, shown in the right panel of Fig. 3. As we did before, we tested
the differences between the two samples, and also in this case the
distribution of jetted sources has larger standard deviation and IQR. In this
case, both statistical tests suggest that the null hypothesis can be rejected
in the observed samples, and also the simulated ones. Therefore, we can
conclude that the two distributions of [O III] core FWHM do not originate from
the same population of sources.
We also studied whether a correlation is present among the core shift and the
FWHM of the [O III] core component. The result is shown in Fig. 4, while the
values we obtained are reported in Table 6. The Pearson correlation
coefficient for all sources is $r=-0.70$, and its p-value, that is the
probability that this result is due to a random distribution of sources, is
8$\times 10^{-13}$. However, as we have seen before, a core shift is present
almost exclusively in jetted sources. Therefore, we decided to test the two
samples separately. In non-jetted sources, no correlation is present
($r=0.02$, p-value = 0.93), indicating that the FWHM of the [O III] core and
its position are independent of each other. Conversely, in jetted sources the
Pearson correlation coefficient is higher ($r=-0.73$), and the p-value is
$4\times 10^{-9}$, indicating that the correlation is present only in this
sample of objects. It is worth noting that this correlation is utterly
dominated by the presence of the blue outliers. Indeed, when the Pearson
coefficient is evaluated in this sample rejecting the blue outliers, its value
becomes $r=-0.28$ (p-value = 0.13). Vice versa, when only the blue outliers
are tested, the correlation remains, although with a lower p-value (r=$-0.67$,
p-value = $3\times 10^{-3}$). In conclusion, [O III] core velocity and FWHM do
not seem to correlate in the bulk of NLS1s population, but they are related
among blue outliers.

Figure 4: Correlation between the FWHM of the [O III] core component and its
shift. Non-jetted sources are represented with the red squares, while jetted
sources are indicated with the blue circles.

####  2.4.2 Wing component

As previously mentioned, the wing component is likely associated to an outflow
coming from the inner part of the NLR. A wing component could be identified
during the decomposition process in all but ten sources, three non-jetted and
seven jetted. We tested the distribution of wing velocity in jetted and non-
jetted sources, shown in the left panel of Fig. 5, and whose parameters are
summarized in Table 3. As for the distributions of the [O III] core
properties, the wing velocity distribution is somewhat peaked in non-jetted
sources, with a large majority of sources having velocities between -400 and
-200 km s-1, while it is more spread out in the jetted sample.

However, the statistical tests seem to indicate that the null hypothesis
cannot be rejected with the required confidence level. The same is true for
the distributions of wing FWHM, displayed in the right panel of Fig. 5. The
presence of relativistic jets does not ultimately seem to affect the FWHM of
the [O III] wing, which remains as turbulent in the jetted NLS1s as in the
non-jetted NLS1s. However, it is worth noting that the standard deviation of
the p-value derived in the simulations is rather large because of the
relatively large measurement errors on the wing properties. Therefore, our
data may not be enough to reach a robust conclusion.
Finally, to better understand if and how the physical quantities we measured
are related to each other, we tested various correlations by evaluating the
Pearson coefficients. The results are shown in Table 6. We first tested the
presence of a correlation between the wing velocity and both [O III] core
velocity and FWHM velocity, but we do not find any strong correlation, except
for a trend between the wing velocity and the [O III] core FWHM in jetted
sources (r = -0.50, p-value = 8$\times 10^{-4}$). In a similar fashion, the
FWHM of the wing component does not appear to correlate with any property of
the [O III] core component, except again for the [O III] core FWHM in jetted
sources (r = 0.54, p-value = 2$\times 10^{-4}$). We finally tested the
presence of a correlation between the wing velocity and its FWHM, and we find
another possible trend among these two quantities in non-jetted sources (r =
-0.51, p-value = 0.01).

Figure 5: Left: distributions of the [O III] wing velocity shift in km s-1, in
non-jetted (top panel) and jetted sources (bottom panel). Binning of 25 km
s-1. Right: distributions of the FWHM([O III]) wing component in km s-1, in
non-jetted (top panel) and jetted sources (bottom panel). Binning of 50 km
s-1.

###  2.5 Radio morphology and [O III]

We tried to investigate whether the radio morphology is somehow connected to
the [O III] properties. As before, we first analyzed the [O III] core velocity
distributions, which are shown in the left side of Fig. 6. Outliers do not
seem to appear preferentially in any specific sample. Among the 26 C sources,
we found one red outlier and six blue, among 27 I sources there is one red
outlier and six blue, and finally in E sources we found one red outlier and
four blue. To check whether the distributions were different from each other,
we used the K-S and A-D tests both on the observed samples and on the
simulated samples. All tests, even when accounting for the p-values
distributions, confirm that the null hypothesis can never be rejected. The
results of all tests, as well as the median velocity, its standard deviation,
and the IQR are all reported in Table 4.

Figure 6: Left: distributions of the [O III] core velocity shift in km s-1, in compact (top panel), intermediate (middle panel) and extended sources (bottom panel). The vertical dashed line indicate the limit for red or blue outliers at 150 km s-1. Binning of 25 km s-1. Right: distributions of the FWHM([O III]) core component in km s-1 in the same samples. Binning of 25 km s-1. Table 4: Statistical tests results for the morphological samples. Columns: (1) Tested sample; (2) median velocity; (3) standard deviation of the velocity; (4) interquartile range; (5) p-value of the K-S tests against the corresponding C sample, observed in top row, simulated in middle row, and standard deviation of the simulated p-value in bottom row; (6) p-value of the A-D tests against the corresponding C sample, observed in top row, simulated in middle row, and standard deviation of the simulated p-value in bottom row; (7) K-S tests against the corresponding I sample, rows as before; (8) A-D tests against the corresponding I sample, rows as before; (9) K-S tests against the corresponding E sample, rows as before; (10) A-D tests against the corresponding E sample, rows as before. Sample |  Median (km s-1) |  Std (km s-1) |  IQR (km s-1) |  K-S${}^{obs}_{C}$ K-S${}^{sim}_{C}$ K-S${}^{std}_{C}$ |  A-D${}^{obs}_{C}$ A-D${}^{sim}_{C}$ A-D${}^{std}_{C}$ |  K-S${}^{obs}_{I}$ K-S${}^{sim}_{I}$ K-S${}^{std}_{I}$ |  A-D${}^{obs}_{I}$ A-D${}^{sim}_{I}$ A-D${}^{std}_{I}$ |  K-S${}^{obs}_{E}$ K-S${}^{sim}_{E}$ K-S${}^{std}_{E}$ |  A-D${}^{obs}_{E}$ A-D${}^{sim}_{E}$ A-D${}^{std}_{E}$
---|---|---|---|---|---|---|---|---|---
vc(C) | -11.6 | 210.4 | 125.7 | $-$ | $-$ |  0.87 0.87 0.12 |  0.89 0.74 0.11 |  0.36 0.33 0.15 |  0.29 0.30 0.08
vc(I) | -31.3 | 207.2 | 116.1 |  0.87 0.87 0.12 |  0.89 0.74 0.11 | $-$ | $-$ |  0.42 0.52 0.18 |  0.32 0.36 0.10
vc(E) | -11.0 | 101.0 | 97.0 |  0.36 0.33 0.15 |  0.29 0.30 0.08 |  0.42 0.52 0.18 |  0.32 0.36 0.10 | $-$ | $-$
FWHMc(C) | 323.8 | 232.2 | 147.3 | $-$ | $-$ |  0.36 0.20 0.18 |  0.15 0.15 0.09 |  0.63 0.60 0.15 |  0.65 0.61 0.12
FWHMc(I) | 385.0 | 285.0 | 193.4 |  0.36 0.20 0.18 |  0.15 0.15 0.09 | $-$ | $-$ |  0.64 0.43 0.20 |  0.56 0.57 0.12
FWHMc(E) | 341.2 | 194.6 | 176.3 |  0.63 0.60 0.15 |  0.65 0.61 0.12 |  0.64 0.43 0.20 |  0.56 0.57 0.12 | $-$ | $-$
vw(C) | -309.8 | 191.9 | 272.2 | $-$ | $-$ |  0.95 0.69 0.23 |  0.92 0.70 0.21 |  0.68 0.44 0.28 |  0.54 0.57 0.25
vw(I) | -302.4 | 229.5 | 276.6 |  0.95 0.69 0.23 |  0.92 0.70 0.21 | $-$ | $-$ |  0.59 0.59 0.24 |  0.60 0.55 0.24
vw(E) | -270.3 | 210.8 | 199.1 |  0.68 0.44 0.28 |  0.54 0.57 0.25 |  0.59 0.59 0.24 |  0.60 0.55 0.24 | $-$ | $-$
FWHMw(C) | 965.2 | 425.0 | 472.4 | $-$ | $-$ |  0.65 0.89 0.14 |  0.81 0.93 0.17 |  1.00 0.99 0.05 |  1.00 1.00 0.10
FWHMw(I) | 1082.1 | 404.8 | 561.9 |  0.65 0.89 0.14 |  0.81 0.93 0.17 | $-$ | $-$ |  0.75 0.93 0.16 |  0.89 0.84 0.15
FWHMw(E) | 999.8 | 383.0 | 567.4 |  1.00 0.99 0.05 |  1.00 1.00 0.10 |  0.75 0.93 0.16 |  0.89 0.84 0.15 | $-$ | $-$

We also tested the distributions of FWHM of the [O III] core component across
the different morphological samples. The results are shown in the right side
of Fig. 6. As for the [O III] core velocity, however, this parameter does not
seem to be related to the radio morphology. All the tests we carried out, and
that are summarized in Table 4, seem to point out that the distributions of
FWHM are essentially the same across the three samples. Indeed, median values,
standard deviation and IQR are all very similar.
There is only one minor detail possibly worth noting, that is the blue
outliers in the E sample are those with velocity closest to the threshold.
Indeed, the [O III] core velocity distribution of E sources has the lowest
standard deviation and IQR. Since most sources with extended morphology do not
have relativistic jets, this result could be a (weak) reflection of the
difference between the distributions of jetted and non-jetted sources we found
in the previous section. However, the samples we are examining do not allow us
to draw any statistically robust conclusions.
The behavior of [O III] wings, as for the core component, does not seem to be
related in any way to the radio morphology. We tested the wing velocity and
FWHM among the three samples, and found that there is no evidence against the
null hypothesis. Detailed results are shown in Table 4.

###  2.6 Effect of the redshift

Figure 7: Left: Left: Redshift distribution of the jetted (top panel) and the
non-jetted (bottom panel) sources. Binning of 0.025. Right: distributions of
the logarithm of the [O III] luminosity (in erg s-1) for the jetted (top
panel) and the non-jetted (bottom panel) sources. Binning of 0.25 dex. Figure
8: The bolometric luminosity of our sources as a function of their redshift.
Non-jetted objects are shown in red, while jetted sources are plotted in blue.

Before moving any further, we tried to quantify the impact that the redshift
could have on some possibly redshift-dependent properties. The redshift
distribution of jetted and non-jetted sources is indeed rather different, as
shown in the left panel of Fig. 7, with jetted sources detected up to higher
redshift. As shown in B18, the same effect is present in the morphological
samples, since sources with extended morphologies are usually not detected at
high redshift due to their low surface brightness densities. We do not expect
significant cosmological evolution of the sources within the redshift range of
our samples, both jetted and non-jetted. Indeed, in the radio luminosity
function of jetted NLS1s, studied by (Berton et al., 2016a), no evolution is
observed up to $z$ = 0.6. Only five out of 48 of our jetted sources lie above
this threshold, therefore we do not expect them to affect significantly our
results. Furthermore, even if some galaxy evolution were present, it is
unlikely to affect the interaction between the relativistic jets and the NLR,
since the physical mechanisms at play remain the same.

We tried also to verify whether the distribution of [O III] luminosity is
different among the jetted and non-jetted sources. Indeed, since jetted
sources are located at higher redshift, there could be a selection bias toward
more luminous sources. The two distributions are shown in the right panel of
Fig. 7, and they are actually different, as both K-S and A-D test show
p-values$\sim 10^{-4}$. However, we tried to verify whether the [O III]
luminosity within the jetted sample is dependent on the redshift, and we found
that only a weak correlation is present (r = 0.39, p-value = 0.01). It is also
worth noting that, at z$<$0.6, even this weak correlation disappears. Sources
with bright [O III], therefore, are not necessarily located at higher
redshift. This suggests that the different luminosity observed in jetted and
non-jetted sources is not a selection effect, but a real physical difference
between the two classes. This should not come as a surprise, since the
presence of a correlation between the [O III] luminosity and the radio
luminosity, which in turn is related to the presence of a relativistic jet,
has been known since the seventies (de Bruyn & Wilson, 1978).
We investigate whether it is possible that the higher [O III] luminosity
instead of the relativistic jets affects the presence of blue outliers. A
higher line luminosity could originate because of a higher disk luminosity
which, in turn, may induce outflows via radiation pressure, and disturbed and
asymmetric profiles are often observed in the [O III] and C IV lines of high-
luminosity sources (Marziani et al., 2016a, b; Sulentic et al., 2017).
However, this may not be the case for our sources. As found already by B16,
and shown by the correlation coefficients of Tables 2 and 6, there is no
evidence that the [O III] core shift correlates with the [O III] line
luminosity, neither when the whole sample nor when the jetted and non jetted
subsamples are considered. We also investigated the redshift dependency of the
bolometric luminosity. Naturally, a strong correlation is present (see Fig. 8)
with higher luminosity sources seen at higher redshift (Spearman’s rs = 0.67,
p-value 3 $\times$ 10-11). However, as seen in Sec. 2.1, the bolometric
luminosity also does not appear to correlate with the [O III] line properties.
Therefore, the presence of the relativistic jets remains the most likely
explanation for blue outliers and the other observed properties, at least for
the sources included in our sample.

###  2.7 [O III] vs spectral index

Testing the [O III] properties as a function of the spectral index will allow
us to understand, at least at a first approximation, whether they are
dependent on the source inclination or not.

Figure 9: Left: distributions of the [O III] core velocity shift in km s-1, in
steep-spectrum (top panel) and flat-spectrum jetted sources (bottom panel).
The vertical dashed line indicate the limit for red or blue outliers at 150 km
s-1. Binning of 25 km s-1. Right: distributions of the FWHM([O III]) wing
component in km s-1 in the same samples. Binning of 25 km s-1.

The resulting distributions of jetted NLS1s are shown in the left panel of
Fig. 9. The number of blue and red outliers in the two samples is essentially
the same. Out of 21 steep-spectrum sources, we found two red outliers and six
blue, while out of 27 flat-spectrum sources we found one red outlier and nine
blue. Neither the K-S nor the A-D test allows us to reject the null hypothesis
that the two distributions of core shift velocities originate from the same
sample of sources. The p-value is always above 0.9, both in our samples and in
the simulated ones, indicating that the two distributions are very similar.
Also median velocity, standard deviation and IQR are comparable among the two
samples. We also analyzed the distributions of FHWM of the [O III] core
component, shown on the right side of Fig. 9. As for the core shift
distributions, the null hypothesis can never be rejected. In a similar way,
the properties of the [O III] wings among the two samples, both their velocity
and FWHM, show that they originate from the same population of sources. All
the values we measured are shown in Table 5.
It is worth noting that we decided to include the seven sources with absorbed
radio jets in the flat-spectrum sample even if, formally, at low radio
frequency they either show a steep spectrum, or they are not detected. The
reason for this is that their rapid variability observed at high radio
frequency indicates that they harbor relativistic jets aligned with the line
of sight, as seen in blazars. In fact, their spectral indexes are steep
because the radio emission at low frequency is dominated by star formation,
but their intrinsic spectrum, if corrected for absorption, has to be flat.
However, our conclusions are independent of this assumption. Even when all of
them are included in the steep-spectrum sample, the statistical tests still do
not provide any evidence against the null hypothesis. This indicates that our
conclusion is rather robust, and that the radio spectral index is not directly
connected to the physical properties of the [O III] line.

Table 5: Results of the statistical tests for the spectral index samples. Columns: (1) Tested sample; (2) median velocity; (3) standard deviation of the velocity; (4) interquartile range; (5) p-value of the K-S tests against the other sample, observed in top row, simulated in middle row, and standard deviation of the simulated p-value in bottom row; (6) p-value of the A-D tests against the other sample, observed in top row, simulated in middle row, and standard deviation of the simulated p-value in bottom row. Sample |  Median (km s-1) |  Std (km s-1) |  IQR (km s-1) |  K-Sobs K-Ssim K-Sstd |  A-Dobs A-Dsim A-Dstd
---|---|---|---|---|---
vc(S) | -4.97 | 226.6 | 252.1 |  0.71 0.92 0.14 |  0.97 0.95 0.11
vc(F) | -12.8 | 225.8 | 226.2 |  0.71 0.92 0.14 |  0.97 0.95 0.11
FWHMc(S) | 408.9 | 311.6 | 266.5 |  0.15 0.14 0.11 |  0.25 0.28 0.09
FWHMc(F) | 321.1 | 268.0 | 253.5 |  0.15 0.14 0.11 |  0.25 0.28 0.09
vw(S) | -362.1 | 246.5 | 355.6 |  0.50 0.50 0.23 |  0.26 0.32 0.23
vw(F) | -326.1 | 229.6 | 339.9 |  0.50 0.50 0.23 |  0.26 0.32 0.23
FWHMw(S) | 1062.6 | 399.4 | 678.3 |  0.68 0.68 0.05 |  0.67 0.69 0.10
FWHMw(F) | 1180.3 | 445.3 | 467.8 |  0.68 0.68 0.05 |  0.67 0.69 0.10

##  3 Discussion

###  3.1 Origin of the bulk outflows

The [O III] lines of our sample could be modeled with two components in 67 out
of 77 sources in which the line is detected (87%), a fraction which is in
agreement with previous findings in the literature (e.g., Mullaney et al.,
2013; Cracco et al., 2016). In type 2 and intermediate type AGN, as well as in
star forming galaxies, instead, the fraction of two-components [O III] lines
is significantly lower (Vaona et al., 2012; Concas et al., 2017). This
indicates that in type 1 AGN two separate kinematic components are directly
visible. The first one is the [O III] core, which is usually not moving with
respect to the galaxy restframe, and the second one is the wing, associated to
an outflowing and more turbulent component closer to the central engine (Bian
et al., 2005).
In some cases both the [O III] core and wing components are shifted toward
bluer wavelengths, although outflows in the opposite direction can also be
observed, likely depending on the geometry of the source. The presence of such
wavelength shift is interpreted as a sign of a bulk outflowing motion in their
NLR. The existence of these outflows in the NLR of AGN has been known for a
long time, and it is particularly common in NLS1s and other population A
sources (Boroson & Green, 1992; Zamanov et al., 2002; Negrete et al., 2018).
In particular, Ganci et al. (2019), analyzing high-quality optical spectra,
found that line displacements are often found in sources with prominent Fe II
multiplets, a feature which is typically observed in NLS1s. This spectral
property has often been associated to the high Eddington ratio that is
characteristic of population A sources, but past studies showed that these two
quantities are not directly correlated (Aoki et al., 2005; Cracco et al.,
2016, B16). An alternative explanation is that the origin of this phenomenon
resides in the relativistic jets crossing the ISM and interacting with it
(Komossa et al., 2008). Such phenomenon, indeed, has been observed not only in
NLS1s, but in several galaxies harboring compact jets (Venturi et al., 2020).
Our results seem to support the latter scenario. As seen in the left panel of
Fig. 3, the number of blue (and red) outliers observed in sources with
relativistic jets is significantly higher than what is found in the non-jetted
sample. This finding is also in agreement with the main result of B16. In that
work, however, the adopted radio classification was based on the radio-
loudness criterion, which has been repeatedly proven to be rather ineffective
(Caccianiga et al., 2015; Padovani, 2017; Järvelä et al., 2017; Lähteenmäki et
al., 2018; Berton et al., 2020b; Chen et al., 2020; Foschini, 2020). In this
work, instead, we make use of a more rigorous classification based directly on
the observed radio properties and morphology of the sources. The relativistic
jets of NLS1s are usually rather compact (Gu et al., 2015; Berton, 2018; Liao
& Gu, 2020; Chen et al., 2020), although some exceptions are known (e.g.,
Gliozzi et al., 2010; Richards & Lister, 2015; Congiu et al., 2017a; Rakshit
et al., 2018; Gabányi et al., 2018; Congiu et al., 2020). This has led some
authors to suggest that NLS1s may be connected with classes of young radio
galaxies such as compact steep-spectrum (CSS) sources and gigahertz-peaked
sources (GPS) (Oshlack et al., 2001; Gallo, 2006; Komossa et al., 2008;
Caccianiga et al., 2014; Berton et al., 2017; Caccianiga et al., 2017; Liao &
Gu, 2020; O’Dea & Saikia, 2020; Yao & Komossa, 2020). Indeed, several radio
galaxies show signs of line shifts and broadening induced by the presence of
jets (Tadhunter et al., 2001; Marziani et al., 2003; Holt et al., 2006;
Venturi et al., 2020), and this seems particularly common in young radio
galaxies, where the jet is directly interacting with the ISM (Morganti et al.,
2015; Orienti, 2016). In this framework, the detection of a large fraction of
outliers among jetted NLS1s is not surprising.
As previously mentioned, Ganci et al. (2019) found that line displacements are
commonly, although not exclusively, present in sources with prominent Fe II
emission. Our data do not allow to test this conclusion, because accurate Fe
II measurements would require high signal-to-noise (S/N) ratio spectra which
are currently unavailable. On the other hand, their sources are not prominent
radio emitters according to the FIRST and NVSS radio surveys at 1.4 GHz
(Becker et al., 1995; Condon et al., 1998). Potentially, this could indicate
that they do not harbor relativistic jets, and that two different physical
mechanisms causing outflows can (co-)exist, a first one connected to the
relativistic jet, and a second one connected to the presence of iron
multiplets. In particular, the iron abundance may be connected to a high
Eddington ratio, as postulated by the quasar main sequence (Marziani et al.,
2018b). This property can produce strong outflows that can ultimately produce
the NLR bulk motion we observed. Indeed, Marziani et al. (2016b) observed that
at intermediate redshift (0.9 $<$ z $<$ 2.2), blue outliers seem to be
significantly more common than what is seen at low redshifts, with a whopping
detection rate of $\sim$40%. Those quasars, however, have bolometric
luminosity in the range 47 $<\log{L_{\rm Bol}}<$ 48.5 erg s-1, that is one
order of magnitude more than the brightest source included in our sample. We
therefore speculate that, at such extreme luminosity, radiation pressure
driven outflows are the most likely mechanism to produce blue outliers. This,
of course, preferably happens at relatively high redshift, since in the local
Universe such extremely bright objects are usually not observed. When,
instead, the bolometric luminosity is smaller, it is possible that the
formation of blue outliers is driven mostly by the presence of relativistic
jets, although the radiation pressure channel cannot be fully ruled out, since
we do observe outliers among non-jetted AGN.
It is worth noting that a weak radio emission at low frequencies does not
necessarily rule out the presence of relativistic jets. As found by Berton et
al. (2020b), the radio emission of some small-scale jets can be entirely
absorbed via free-free or synchrotron self-absorption, thus making these
features invisible in the low-frequency domain ($<$10 GHz). Their interaction
with the ISM is instead well visible, since two out of seven of these peculiar
objects have shifted [O III] lines (J1232+4957, J1509+6137). Therefore, we
suggest that blue outliers, especially those showing the most prominent line
displacements, could be potential targets in which an absorbed jet could be
present. Only high-frequency radio observations will allow us to clarify this
aspect.
As in Komossa et al. (2018), we found a rather strong correlation between the
FWHM and the shift of the [O III] core component (Fig. 4). This seem to
indicate that a bulk motion in the NLR (the core shift) is characterized by a
larger intrinsic turbulence of the gas (the [O III] core FHWM). This result
was observed not only in [O III] lines, but also in other NLR lines (such as
[Fe VII], [Ne III], [Ne V], [Fe X]), and in the BLR line C IV $\lambda$1549
(Komossa et al., 2008; Spoon & Holt, 2009; Gaskell & Goosmann, 2013; Cracco et
al., 2016). In our analysis, this correlation is dominated by the outliers
found in the jetted sources, again suggesting that this behavior can be
produced by the relativistic jets crossing the NLR. Another effect of this
interaction may be seen in the wing component, since its velocity distribution
is different in jetted and non-jetted sources, with the former having a larger
scatter than the latter.
Another interesting result of our analysis is the non-detection of the [O III]
line in three of our sources, J0744+5149, J0945+1915, and J1421+2824. Except
for J0744+5149, where the S/N ratio in the continuum is low ($\sim 3$), the
other sources have relatively good spectra (S/N$>$10). Therefore, the non-
detection in these cases is not due to a poor data quality, but to an
intrinsic weakness of the [O III] lines. All of these sources have
relativistic jets with a compact morphology and a flat radio spectrum. This
may not be a coincidence. The presence of relativistic jets can decrease the
covering factor of the clouds where high ionization lines such as [O III]
originate, and as a consequence reduce the equivalent width and luminosity of
the lines (Cecil et al., 2002; Ludwig et al., 2012). Therefore, the non
detection of these lines may be another sign of jet/NLR interaction.
The reason why all jetted sources are not blue outliers is still not fully
understood. A possible explanation for this behavior, however, was provided in
the simulations by Wagner et al. (2012). The interaction between the
relativistic jets and the NLR can indeed occur only if the ratio between the
total jet power and the Eddington luminosity of the source exceeds a threshold
value (log(Pj/LEdd) $>$ -4). When this condition is met, the velocity of the
outflow produced by the relativistic jet of power Pj exceeds the gas velocity
dispersion derived from the relation between the black hole mass and the
stellar velocity dispersion (Ferrarese & Merritt, 2000) of a galaxy for which
the Eddington luminosity is LEdd. Therefore, the gas can escape the
gravitational potential of the black hole. Unfortunately, our data do not
allow us to test this hypothesis for all of our sources. However, for a few
sources the jet power was calculated by Foschini et al. (2015), and all of
them respect this condition, as already shown by B16. In conclusion, we
speculate that only the high-power jets launched by relatively low-mass - and
hence low Eddington luminosity - sources, such as in NLS1s, can produce blue
outliers.

###  3.2 [O III] and radio properties

A direct comparison between the radio morphology and the [O III] properties in
high Eddington ratio AGN such as NLS1s has never been carried out, since the
number of radio surveys aimed at studying the kpc-scale properties of these
objects is very limited. The role of the radio morphology was studied by
dividing the sources in compact, intermediate, and extended, based on the
concentration of the flux density in the radio core. As found by B18, jetted
sources often have compact or intermediate morphology, while non-jetted ones
are mostly extended. However, our results seem to indicate that this is not
connected to any of the [O III] properties we measured. The number of outliers
found in the different morphological samples is basically the same, and also
the other properties do not show preferences for any sample. This is somewhat
surprising taking into account that jetted sources preferably have a compact
morphology, and we therefore would have expected to see this trend reflected
on the morphological distributions.
Interestingly, we also found that the [O III] properties in jetted sources are
the same when the radio spectral index is flat or steep. As previously
described, the spectral index can be used as a rough proxy to estimate the
source inclination. Flat-spectrum objects are preferentially observed at low
inclination with respect to the line of sight, while in steep-spectrum sources
the relativistic jets are observed at larger angles. This suggests that the
perturbations in the ionized gas kinematics can be seen regardless of the jet
inclination. Since the structure of AGN has a cylindrical symmetry, also this
result is slightly counterintuitive, as one would expect to observe outflows
preferentially at low inclination (Komossa et al., 2006; Yuan et al., 2008),
but it is not utterly surprising. If orientation were the sole responsible for
blue outliers, we should see them in sources with the highest relativistic
beaming and, as a consequence, the highest radio luminosity. However, this has
never been the case, as clearly shown in the past (Berton et al., 2016b;
Komossa et al., 2018). Our work confirms this result. The Pearson correlation
coefficient between the core shift velocity and the total radio luminosity at
5 GHz is 0.005, with a p-value of 0.97, providing no evidence that these two
quantities are connected. Also no correlation was found in the jetted and non-
jetted subsamples (see Table 6).
Determining the exact inclination of an AGN, especially non-jetted, is not an
easy task. The most reliable technique relies on the geometry of the
ionization cones, but it can be applied only for a limited number of sources
(Marin, 2016). Blue outliers, until now, were commonly considered low-
inclination sources. Indeed, Boroson (2011) suggested that they could be used,
along with the red-shift of Fe II multiplets ((Hu et al., 2008), but see (Bon
et al., 2018)), as inclination indicators. However, our findings seriously
challenge this conclusion, since blue outliers seem to be uniformly
distributed among sources with different inclinations.
The cloud acceleration process, therefore, seems to occur in every direction,
and not only along the jet propagation axis. This behavior is similar to what
was predicted by hydrodynamical simulations by Wagner & Bicknell (2011) and
Wagner et al. (2012)444A very explanatory video can be found at
https://www.youtube.com/watch?v=6-CO4hp6PYg. In their simulations, the jet can
accelerate the clouds to velocities up to $10^{3}$ km s-1, which are
comparable to the highest values we observed in our sample. Similarly, Cielo
et al. (2018) found that young jets can accelerate the gas clouds up to these
velocities much more efficiently than what nuclear outflows induced by disk
winds can do. Therefore, we can conclude that faster blue outliers are the
most likely to originate from jet-ISM interaction.

##  4 Summary

In this paper, we analyzed the properties of the [O III] lines in a sample of
low-mass high Eddington ratio AGN belonging to the NLS1 class, and compared
them with their radio properties reported by a recent JVLA survey at 5 GHz. We
found strong evidence that the presence of relativistic jets in these AGN can
produce a bulk outflowing motion in the NLR, causing the [O III] lines to
appear blue-shifted with respect to their restframe wavelength. Among the non-
jetted AGN we found one outlier out of 28 sources, while amidst the jetted AGN
we identified 18 outliers out of 48 sources. No strong correlation between the
bolometric luminosity of the sources and the [O III] line properties was
found, suggesting only a minor contribution of radiation pressure driven
outflows in producing [O III] shifts. We therefore suggest that the most
common origin for these so-called blue outliers is the interaction between the
NLR gas clouds producing [O III] lines and the relativistic jets. The presence
of blue outliers seems not to be connected either with the radio morphology of
the source, or with the inclination of radio jets. This possibly indicates
that the acceleration of the clouds produced by the propagation of the jets in
the NLR medium does not have a preferred direction. The association between
the blue outliers and the relativistic jets might indicate that the properties
of the [O III] lines can be used as a powerful tool to identify the presence
of jets whose low-frequency radio emission is absorbed, and that have gone
undetected in the past radio surveys. Further studies, especially by means of
spatially resolved integral-field spectroscopy, are needed to fully understand
the interplay between the relativistic jets and the NLR clouds.

## References

  * Abazajian et al. (2009) Abazajian, K. N., Adelman-McCarthy, J. K., Agüeros, M. A., et al. 2009, ApJS, 182, 543
  * Abdo et al. (2009) Abdo, A. A., Ackermann, M., Ajello, M., et al. 2009, ApJ, 699, 976
  * Antón et al. (2008) Antón, S., Browne, I. W. A., & Marchã, M. J. 2008, A&A, 490, 583
  * Aoki et al. (2005) Aoki, K., Kawaguchi, T., & Ohta, K. 2005, ApJ, 618, 601
  * Becker et al. (1995) Becker, R. H., White, R. L., & Helfand, D. J. 1995, ApJ, 450, 559
  * Berton (2018) Berton, M. 2018, in Proceedings of Science, vol. Revisiting narrow-line Seyfert 1 galaxies and their place in the Universe. 9-13 April 2018, 26
  * Berton et al. (2020a) Berton, M., Björklund, I., Lähteenmäki, A., et al. 2020a, Contributions of the Astronomical Observatory Skalnate Pleso, 50, 270
  * Berton et al. (2016a) Berton, M., Caccianiga, A., Foschini, L., et al. 2016a, A&A, 591, A98
  * Berton et al. (2019) Berton, M., Congiu, E., Ciroi, S., et al. 2019, AJ, 157, 48
  * Berton et al. (2018) Berton, M., Congiu, E., Järvelä, E., et al. 2018, A&A, 614, A87
  * Berton et al. (2017) Berton, M., Foschini, L., Caccianiga, A., et al. 2017, Frontiers in Astronomy and Space Sciences, 4, 8
  * Berton et al. (2016b) Berton, M., Foschini, L., Ciroi, S., et al. 2016b, A&A, 591, A88
  * Berton et al. (2015) Berton, M., Foschini, L., Ciroi, S., et al. 2015, A&A, 578, A28
  * Berton et al. (2020b) Berton, M., Järvelä, E., Crepaldi, L., et al. 2020b, A&A, 636, A64
  * Bian et al. (2005) Bian, W., Yuan, Q., & Zhao, Y. 2005, MNRAS, 364, 187
  * Blandford & Königl (1979) Blandford, R. D. & Königl, A. 1979, ApJ, 232, 34
  * Blanton et al. (2017) Blanton, M. R., Bershady, M. A., Abolfathi, B., et al. 2017, AJ, 154, 28
  * Bon et al. (2018) Bon, E., Marziani, P., Berton, M., et al. 2018, in Revisiting Narrow-Line Seyfert 1 Galaxies and their Place in the Universe, 7
  * Boroson (2002) Boroson, T. A. 2002, ApJ, 565, 78
  * Boroson (2011) Boroson, T. A. 2011, ApJL, 735, L14
  * Boroson & Green (1992) Boroson, T. A. & Green, R. F. 1992, ApJS, 80, 109
  * Caccianiga et al. (2014) Caccianiga, A., Antón, S., Ballo, L., et al. 2014, MNRAS, 441, 172
  * Caccianiga et al. (2015) Caccianiga, A., Antón, S., Ballo, L., et al. 2015, MNRAS, 451, 1795
  * Caccianiga et al. (2017) Caccianiga, A., Dallacasa, D., Antón, S., et al. 2017, MNRAS, 464, 1474
  * Cecil et al. (2002) Cecil, G., Dopita, M. A., Groves, B., et al. 2002, ApJ, 568, 627
  * Chen et al. (2020) Chen, S., Järvelä, E., Crepaldi, L., et al. 2020, MNRAS, 498, 1278
  * Cielo et al. (2018) Cielo, S., Bieri, R., Volonteri, M., Wagner, A. Y., & Dubois, Y. 2018, MNRAS, 477, 1336
  * Concas et al. (2017) Concas, A., Popesso, P., Brusa, M., et al. 2017, A&A, 606, A36
  * Condon et al. (1998) Condon, J. J., Cotton, W. D., Greisen, E. W., et al. 1998, AJ, 115, 1693
  * Congiu et al. (2017a) Congiu, E., Berton, M., Giroletti, M., et al. 2017a, A&A, 603, A32
  * Congiu et al. (2017b) Congiu, E., Contini, M., Ciroi, S., et al. 2017b, Frontiers in Astronomy and Space Sciences, 4, 27
  * Congiu et al. (2020) Congiu, E., Kharb, P., Tarchi, A., et al. 2020, MNRAS, 499, 3149
  * Cracco et al. (2016) Cracco, V., Ciroi, S., Berton, M., et al. 2016, MNRAS, 462, 1256
  * Crenshaw et al. (2003) Crenshaw, D. M., Kraemer, S. B., & Gabel, J. R. 2003, AJ, 126, 1690
  * Czerny et al. (2018) Czerny, B., Panda, S., Śniegowska, M., et al. 2018, in Revisiting Narrow-Line Seyfert 1 Galaxies and their Place in the Universe, 30
  * de Bruyn & Wilson (1978) de Bruyn, A. G. & Wilson, A. S. 1978, A&A, 64, 433
  * Deo et al. (2006) Deo, R. P., Crenshaw, D. M., & Kraemer, S. B. 2006, AJ, 132, 321
  * Dimitrijević et al. (2007) Dimitrijević, M. S., Popović, L. Č., Kovačević, J., Dačić, M., & Ilić, D. 2007, MNRAS, 374, 1181
  * Falcke & Biermann (1995) Falcke, H. & Biermann, P. L. 1995, A&A, 293, 665
  * Ferrarese & Merritt (2000) Ferrarese, L. & Merritt, D. 2000, ApJL, 539, L9
  * Foschini (2011) Foschini, L. 2011, in Narrow-Line Seyfert 1 Galaxies and their Place in the Universe, Proc. of Science, Vol. NLS1, id. 24, 24
  * Foschini (2020) Foschini, L. 2020, Universe, 6, 136
  * Foschini et al. (2015) Foschini, L., Berton, M., Caccianiga, A., et al. 2015, A&A, 575, A13
  * Fragile et al. (2017) Fragile, P. C., Anninos, P., Croft, S., Lacy, M., & Witry, J. W. L. 2017, ApJ, 850, 171
  * Gabányi et al. (2018) Gabányi, K. É., Frey, S., Paragi, Z., et al. 2018, MNRAS, 473, 1554
  * Gallo (2006) Gallo, L. C. 2006, MNRAS, 368, 479
  * Ganci et al. (2019) Ganci, V., Marziani, P., D’Onofrio, M., et al. 2019, A&A, 630, A110
  * Gaskell & Goosmann (2013) Gaskell, C. M. & Goosmann, R. W. 2013, ApJ, 769, 30
  * Gliozzi et al. (2010) Gliozzi, M., Papadakis, I. E., Grupe, D., et al. 2010, ApJ, 717, 1243
  * Grupe (2000) Grupe, D. 2000, New Astron. Rev., 44, 455
  * Grupe et al. (2010) Grupe, D., Komossa, S., Leighly, K. M., & Page, K. L. 2010, ApJS, 187, 64
  * Grupe et al. (2004) Grupe, D., Mathur, S., & Komossa, S. 2004, AJ, 127, 3161
  * Gu et al. (2015) Gu, M., Chen, Y., Komossa, S., et al. 2015, ApJS, 221, 3
  * Holt et al. (2006) Holt, J., Tadhunter, C., Morganti, R., et al. 2006, MNRAS, 370, 1633
  * Hou et al. (2009) Hou, A., Parker, L. C., Harris, W. E., & Wilman, D. J. 2009, ApJ, 702, 1199
  * Hu et al. (2008) Hu, C., Wang, J.-M., Ho, L. C., et al. 2008, ApJ, 687, 78
  * Järvelä et al. (2018) Järvelä, E., Lähteenmäki, A., & Berton, M. 2018, A&A, 619, A69
  * Järvelä et al. (2017) Järvelä, E., Lähteenmäki, A., Lietzen, H., et al. 2017, A&A, 606, A9
  * Kellermann et al. (1989) Kellermann, K. I., Sramek, R., Schmidt, M., Shaffer, D. B., & Green, R. 1989, AJ, 98, 1195
  * Komatsu et al. (2011) Komatsu, E., Smith, K. M., Dunkley, J., et al. 2011, ApJS, 192, 18
  * Komossa et al. (2006) Komossa, S., Voges, W., Xu, D., et al. 2006, AJ, 132, 531
  * Komossa et al. (2008) Komossa, S., Xu, D., Zhou, H., Storchi-Bergmann, T., & Binette, L. 2008, ApJ, 680, 926
  * Komossa et al. (2018) Komossa, S., Xu, D. W., & Wagner, A. Y. 2018, MNRAS, 477, 5115
  * Kormendy & Ho (2013) Kormendy, J. & Ho, L. C. 2013, ARA&A, 51, 511
  * Kotilainen et al. (2016) Kotilainen, J. K., León-Tavares, J., Olguín-Iglesias, A., et al. 2016, ApJ, 832, 157
  * Kovačević et al. (2010) Kovačević, J., Popović, L. Č., & Dimitrijević, M. S. 2010, ApJS, 189, 15
  * Krongold et al. (2001) Krongold, Y., Dultzin-Hacyan, D., & Marziani, P. 2001, AJ, 121, 702
  * Lähteenmäki et al. (2018) Lähteenmäki, A., Järvelä, E., Ramakrishnan, V., et al. 2018, A&A, 614, L1
  * Liao & Gu (2020) Liao, M. & Gu, M. 2020, MNRAS, 491, 92
  * Ludwig et al. (2012) Ludwig, R. R., Greene, J. E., Barth, A. J., & Ho, L. C. 2012, ApJ, 756, 51
  * Marin (2016) Marin, F. 2016, MNRAS, 460, 3679
  * Marziani et al. (2018a) Marziani, P., del Olmo, A., D’Onofrio, M., et al. 2018a, in Proceedings of Science, vol. Revisiting narrow-line Seyfert 1 galaxies and their place in the Universe, 2
  * Marziani et al. (2018b) Marziani, P., Dultzin, D., Sulentic, J. W., et al. 2018b, Frontiers in Astronomy and Space Sciences, 5, 6
  * Marziani et al. (2016a) Marziani, P., Martínez Carballo, M. A., Sulentic, J. W., et al. 2016a, Ap&SS, 361, 29
  * Marziani et al. (2016b) Marziani, P., Sulentic, J. W., Stirpe, G. M., et al. 2016b, Ap&SS, 361, 3
  * Marziani et al. (2001) Marziani, P., Sulentic, J. W., Zwitter, T., Dultzin-Hacyan, D., & Calvani, M. 2001, ApJ, 558, 553
  * Marziani et al. (2003) Marziani, P., Zamanov, R. K., Sulentic, J. W., & Calvani, M. 2003, MNRAS, 345, 1133
  * Mathur (2000) Mathur, S. 2000, MNRAS, 314, L17
  * Mathur et al. (2012) Mathur, S., Fields, D., Peterson, B. M., & Grupe, D. 2012, ApJ, 754, 146
  * Mathur et al. (2001) Mathur, S., Kuraszkiewicz, J., & Czerny, B. 2001, NewA, 6, 321
  * Morganti (2017) Morganti, R. 2017, Frontiers in Astronomy and Space Sciences, 4, 42
  * Morganti et al. (2015) Morganti, R., Oosterloo, T., Oonk, J. B. R., Frieswijk, W., & Tadhunter, C. 2015, A&A, 580, A1
  * Mullaney et al. (2013) Mullaney, J. R., Alexander, D. M., Fine, S., et al. 2013, MNRAS, 433, 622
  * Negrete et al. (2018) Negrete, C. A., Dultzin, D., Marziani, P., et al. 2018, A&A, 620, A118
  * O’Dea & Saikia (2020) O’Dea, C. P. & Saikia, D. J. 2020, arXiv e-prints, arXiv:2009.02750
  * Ohta et al. (2007) Ohta, K., Aoki, K., Kawaguchi, T., & Kiuchi, G. 2007, ApJS, 169, 1
  * Olguín-Iglesias et al. (2019) Olguín-Iglesias, A., Kotilainen, J., & Chavushyan, V. 2019, arXiv e-prints, arXiv:1912.02746
  * Olguín-Iglesias et al. (2017) Olguín-Iglesias, A., Kotilainen, J. K., León Tavares, J., Chavushyan, V., & Añorve, C. 2017, MNRAS, 467, 3712
  * Orban de Xivry et al. (2011) Orban de Xivry, G., Davies, R., Schartmann, M., et al. 2011, MNRAS, 417, 2721
  * Orienti (2016) Orienti, M. 2016, Astronomische Nachrichten, 337, 9
  * Oshlack et al. (2001) Oshlack, A. Y. K. N., Webster, R. L., & Whiting, M. T. 2001, ApJ, 558, 578
  * Osterbrock & Pogge (1985) Osterbrock, D. E. & Pogge, R. W. 1985, ApJ, 297, 166
  * Padovani (2017) Padovani, P. 2017, Nature Astronomy, 1, 0194
  * Paliya (2019) Paliya, V. S. 2019, Journal of Astrophysics and Astronomy, 40, 39
  * Panda et al. (2019) Panda, S., Marziani, P., & Czerny, B. 2019, ApJ, 882, 79
  * Panessa et al. (2019) Panessa, F., Baldi, R. D., Laor, A., et al. 2019, Nature Astronomy, 3, 387
  * Peterson (2011) Peterson, B. M. 2011, in Narrow-Line Seyfert 1 Galaxies and their Place in the Universe, 32
  * Rakshit et al. (2017) Rakshit, S., Stalin, C. S., Chand, H., & Zhang, X.-G. 2017, ApJS, 229, 39
  * Rakshit et al. (2018) Rakshit, S., Stalin, C. S., Hota, A., & Konar, C. 2018, ApJ, 869, 173
  * Richards & Lister (2015) Richards, J. L. & Lister, M. L. 2015, ApJL, 800, L8
  * Schmitt et al. (2001) Schmitt, H. R., Ulvestad, J. S., Antonucci, R. R. J., & Kinney, A. L. 2001, ApJS, 132, 199
  * Shapovalova et al. (2012) Shapovalova, A. I., Popović, L. Č., Burenkov, A. N., et al. 2012, ApJS, 202, 10
  * Shen & Ho (2014) Shen, Y. & Ho, L. C. 2014, Nature, 513, 210
  * Spoon & Holt (2009) Spoon, H. W. W. & Holt, J. 2009, ApJL, 702, L42
  * Sulentic et al. (2007) Sulentic, J. W., Bachev, R., Marziani, P., Negrete, C. A., & Dultzin, D. 2007, ApJ, 666, 757
  * Sulentic et al. (2017) Sulentic, J. W., del Olmo, A., Marziani, P., et al. 2017, A&A, 608, A122
  * Sulentic et al. (2002) Sulentic, J. W., Marziani, P., Zamanov, R., et al. 2002, ApJL, 566, L71
  * Sulentic et al. (2000) Sulentic, J. W., Zwitter, T., Marziani, P., & Dultzin-Hacyan, D. 2000, ApJL, 536, L5
  * Tadhunter et al. (2001) Tadhunter, C., Wills, K., Morganti, R., Oosterloo, T., & Dickson, R. 2001, MNRAS, 327, 227
  * Taylor & Kobayashi (2015) Taylor, P. & Kobayashi, C. 2015, MNRAS, 448, 1835
  * Tripathi et al. (2020) Tripathi, S., McGrath, K. M., Gallo, L. C., et al. 2020, MNRAS, 499, 1266
  * Vaona et al. (2012) Vaona, L., Ciroi, S., Di Mille, F., et al. 2012, MNRAS, 427, 1266
  * Venturi et al. (2020) Venturi, G., Cresci, G., Marconi, A., et al. 2020, arXiv e-prints, arXiv:2011.04677
  * Venturi et al. (2018) Venturi, G., Nardini, E., Marconi, A., et al. 2018, A&A, 619, A74
  * Wagner & Bicknell (2011) Wagner, A. Y. & Bicknell, G. V. 2011, ApJ, 728, 29
  * Wagner et al. (2012) Wagner, A. Y., Bicknell, G. V., & Umemura, M. 2012, ApJ, 757, 136
  * Xu et al. (2012) Xu, D., Komossa, S., Zhou, H., et al. 2012, AJ, 143, 83
  * Yao & Komossa (2020) Yao, S. & Komossa, S. 2020, arXiv e-prints, arXiv:2012.01785
  * Yuan et al. (2008) Yuan, W., Zhou, H. Y., Komossa, S., et al. 2008, ApJ, 685, 801
  * Zamanov et al. (2002) Zamanov, R., Marziani, P., Sulentic, J. W., et al. 2002, ApJL, 576, L9

Table 6: Results of the correlations tested throughout this work. Columns: (1) First tested sample; (2) second tested sample; (3) Pearson correlation coefficient; (4) p-value. Q1 | Q2 | r | p-value
---|---|---|---
v${}_{c}^{tot}$ | FWHM${}_{c}^{tot}$ | -0.70 | 8$\times 10^{-13}$
v${}_{c}^{NJ}$ | FWHM${}_{c}^{NJ}$ | 0.02 | 0.93
v${}_{c}^{J}$ | FWHM${}_{c}^{J}$ | -0.73 | 4$\times 10^{-9}$
v${}_{w}^{tot}$ | v${}_{c}^{tot}$ | 0.01 | 0.92
v${}_{w}^{NJ}$ | v${}_{c}^{NJ}$ | 0.20 | 0.35
v${}_{w}^{J}$ | v${}_{c}^{J}$ | 0.24 | 0.13
v${}_{w}^{tot}$ | FWHM${}_{c}^{tot}$ | -0.09 | 0.43
v${}_{w}^{NJ}$ | FWHM${}_{c}^{NJ}$ | -0.06 | 0.79
v${}_{w}^{J}$ | FWHM${}_{c}^{J}$ | -0.50 | 8$\times 10^{-4}$
FWHM${}_{w}^{tot}$ | v${}_{c}^{tot}$ | 0.02 | 0.85
FWHM${}_{w}^{NJ}$ | v${}_{c}^{NJ}$ | -0.21 | 0.32
FWHM${}_{w}^{J}$ | v${}_{c}^{J}$ | -0.36 | 0.02
FWHM${}_{w}^{tot}$ | FWHM${}_{c}^{tot}$ | -0.01 | 0.95
FWHM${}_{w}^{NJ}$ | FWHM${}_{c}^{NJ}$ | 0.34 | 0.09
FWHM${}_{w}^{J}$ | FWHM${}_{c}^{J}$ | 0.55 | 2$\times 10^{-4}$
v${}_{w}^{tot}$ | FWHM${}_{w}^{tot}$ | -0.47 | 1$\times 10^{-5}$
v${}_{w}^{NJ}$ | FWHM${}_{w}^{NJ}$ | -0.51 | 0.01
v${}_{w}^{J}$ | FWHM${}_{w}^{J}$ | -0.08 | 0.61
v${}_{c}^{NJ}$ | L${}^{5~{}GHz}_{int}$ | -0.30 | 0.12
v${}_{c}^{J}$ | L${}^{5~{}GHz}_{int}$ | 0.03 | 0.84
v${}_{c}^{NJ}$ | L${}_{[O~{}III]}$ | 0.22 | 0.26
v${}_{c}^{J}$ | L${}_{[O~{}III]}$ | -0.01 | 0.94

Table 7: Sample studied in this work. Columns: (1) short name; (2) NED alias; (3) right ascension; (4) declination; (5) redshift; (6) jet (Y stands for yes, N for no); (7) radio morphology (C stands for compact, I for intermediate, E for extended); (8) spectral index for jetted sources (F stands for flat spectrum, S stands for steep spectrum). Short Name | NED Alias | R.A. | Dec. | z | Jet | Morph | Sp.Ind.
---|---|---|---|---|---|---|---
J0006$+$2012 | Mrk 335 | 00 06 19.50 | +20 12 10.0 | 0.026 | N | C | $-$
J0146$-$0040 |  2MASX J01464481$-$0040426 | 01 46 44.80 |  $-$00 40 43.0 | 0.083 | N | C | $-$
J0347$+$0105 | IRAS 03450+0055 | 03 47 40.20 | +01 05 14.0 | 0.031 | N | E | $-$
J0629$-$0545 |  IRAS 06269$-$0543 | 06 29 24.70 |  $-$05 45 30.0 | 0.117 | Y | I | S
J0632$+$6340 | UGC 3478 | 06 32 47.20 | +63 40 25.0 | 0.013 | N | E | $-$
J0706$+$3901 | FBQS J0706+3901 | 07 06 25.15 | +39 01 51.6 | 0.086 | Y | I | S
J0713$+$3820 | FBQS J0713+3820 | 07 13 40.29 | +38 20 40.1 | 0.123 | Y | E | S
J0744$+$5149 | NVSS J074402+514917 | 07 44 02.24 | +51 49 17.5 | 0.460 | Y | C | F
J0752$+$2617 | FBQS J0752+2617 | 07 52 45.60 | +26 17 36.0 | 0.082 | N | C | $-$
J0758$+$3920 | FBQS J075800.0+392029 | 07 58 00.05 | +39 20 29.1 | 0.096 | Y | E | S
J0804$+$3853 | FBQS J0804+3853 | 08 04 09.24 | +38 53 48.7 | 0.212 | Y | E | S
J0806$+$7248 | RGB J0806+728 | 08 06 38.96 | +72 48 20.4 | 0.098 | Y | I | S
J0814$+$5609 | SDSS J081432.11+560956.6 | 08 14 32.13 | +56 09 56.6 | 0.510 | Y | I | F
J0849$+$5108 | SBS 0846+513 | 08 49 57.99 | +51 08 28.8 | 0.585 | Y | C | F
J0850$+$4626 | SDSS J085001.17+462600.5 | 08 50 01.17 | +46 26 00.5 | 0.524 | Y | C | F
J0902$+$0443 | SDSS J090227.16+044309.5 | 09 02 27.15 | +04 43 09.4 | 0.533 | Y | C | F
J0913$+$3658 | RX J0913.2+3658 | 09 13 13.70 | +36 58 17.0 | 0.107 | N | I | $-$
J0925$+$5217 | Mrk 110 | 09 25 12.90 | +52 17 11.0 | 0.035 | N | E | $-$
J0926$+$1244 | Mrk 705 | 09 26 03.30 | +12 44 04.0 | 0.029 | N | I | $-$
J0937$+$3615 | SDSS J093703.03+361537.2 | 09 37 03.01 | +36 15 37.3 | 0.180 | N | E | $-$
J0945$+$1915 | SDSS J094529.23+191548.8 | 09 45 29.21 | +19 15 48.9 | 0.284 | Y | C | F
J0948$+$5029 | Mrk 124 | 09 48 42.60 | +50 29 31.0 | 0.056 | N | C | $-$
J0952$-$0136 | Mrk 1239 | 09 52 19.10 |  $-$01 36 43.0 | 0.020 | Y | I | S
J0957$+$2433 | RX J0957.1+2433 | 09 57 07.20 | +24 33 16.0 | 0.082 | N | C | $-$
J1029$+$5556 | SDSS J102906.69+555625.2 | 10 29 06.69 | +55 56 25.2 | 0.451 | Y | C | F
J1031$+$4234 | SDSS J103123.73+423439.3 | 10 31 23.73 | +42 34 39.4 | 0.377 | Y | C | S
J1034$+$3938 | KUG 1031+398 | 10 34 38.60 | +39 38 28.0 | 0.042 | Y | I | S
J1037$+$0036 | SDSS J103727.45+003635.6 | 10 37 27.45 | +00 36 35.8 | 0.595 | Y | C | F
J1038$+$4227 | SDSS J103859.58+422742.3 | 10 38 59.59 | +42 27 42.0 | 0.220 | N | E | $-$
J1047$+$4725 | SDSS J104732.68+472532.0 | 10 47 32.65 | +47 25 32.2 | 0.799 | Y | E | F
J1048$+$2222 | SDSS J104816.57+222238.9 | 10 48 16.56 | +22 22 40.1 | 0.330 | Y | I | S
J1102$+$2239 | SDSS J110223.38+223920.7 | 11 02 23.36 | +22 39 20.7 | 0.453 | Y | I | S
J1110$+$3653 | SDSS J111005.03+365336.3 | 11 10 05.03 | +36 53 36.1 | 0.629 | Y | I | F
J1114$+$3241 | B2 1111+32 | 11 14 38.89 | +32 41 33.4 | 0.189 | Y | C | F
J1121$+$5351 | SBS 1118+541 | 11 21 08.60 | +53 51 21.0 | 0.103 | N | I | $-$
J1138$+$3653 | SDSS J113824.54+365327.1 | 11 38 24.54 | +36 53 27.1 | 0.356 | Y | I | S
J1146$+$3236 | SDSS J114654.28+323652.3 | 11 46 54.30 | +32 36 52.2 | 0.465 | Y | C | F
J1159$+$2838 | SDSS J115917.32+283814.5 | 11 59 17.31 | +28 38 14.8 | 0.210 | Y | I | S
J1203$+$4431 | NGC 4051 | 12 03 09.60 | +44 31 53.0 | 0.002 | N | E | $-$
J1209$+$3217 | RX J1209.7+3217 | 12 09 45.20 | +32 17 01.0 | 0.144 | N | I | $-$
J1215$+$5442 | SBS 1213+549A | 12 15 49.40 | +54 42 24.0 | 0.150 | N | I | $-$
J1218$+$2948 | Mrk 766 | 12 18 26.50 | +29 48 45.8 | 0.013 | N | E | $-$
J1227$+$3214 | SDSS J122749.14+321458.9 | 12 27 49.15 | +32 14 59.0 | 0.136 | Y | I | S
J1228$+$5017 | SDSS J122844.81+501751.2 | 12 28 44.81 | +50 17 51.2 | 0.262 | Y | I | F
J1232$+$4957 | SDSS J123220.11+495721.8 | 12 32 20.11 | +49 57 21.8 | 0.262 | Y | C | F
J1238$+$3942 | SDSS J123852.12+394227.8 | 12 27 49.15 | +39 42 27.6 | 0.622 | Y | C | F
J1242$+$3317 | WAS 61 | 12 42 10.60 | +33 17 03.0 | 0.044 | N | E | $-$
J1246$+$0222 | PG 1244+026 | 12 46 35.20 | +02 22 09.0 | 0.048 | N | E | $-$
J1246$+$0238 | SDSS J124634.65+023809.0 | 12 46 34.68 | +02 38 09.0 | 0.367 | Y | C | F
J1302$+$1624 | Mrk 783 | 13 02 58.80 | +16 24 27.0 | 0.067 | Y | E | S
J1305$+$5116 | SDSS J130522.74+511640.2 | 13 05 22.75 | +51 16 40.3 | 0.788 | Y | E | F
J1317$+$6010 | SBS 1315+604 | 13 17 50.30 | +60 10 41.0 | 0.137 | N | E | $-$
J1333$+$4141 | SDSS J133345.47+414127.7 | 13 33 45.47 | +41 41 28.2 | 0.225 | Y | C | S
J1337$+$2423 | IRAS 13349+2438 | 13 37 18.70 | +24 23 03.0 | 0.108 | N | I | $-$
J1346$+$3121 | SDSS J134634.97+312133.7 | 13 46 35.07 | +31 21 33.9 | 0.246 | Y | C | F
J1355$+$5612 | SBS 1353+564 | 13 55 16.50 | +56 12 45.0 | 0.122 | N | C | $-$
J1358$+$2658 | SDSS J135845.38+265808.4 | 13 58 45.40 | +26 58 08.3 | 0.331 | Y | I | S
J1402$+$2159 | RX J1402.5+2159 | 14 02 34.40 | +21 59 52.0 | 0.066 | N | I | $-$
J1421$+$2824 | SDSS J142114.05+282452.8 | 14 21 14.07 | +28 24 52.2 | 0.539 | Y | C | F
J1443$+$4725 | SDSS J144318.56+472556.7 | 14 43 18.56 | +47 25 56.7 | 0.705 | Y | C | F
J1505$+$0326 | SDSS J150506.47+032630.8 | 15 05 06.47 | +03 26 30.8 | 0.408 | Y | C | F
J1509$+$6137 | SDSS J150916.18+613716.7 | 15 09 16.18 | +61 37 16.7 | 0.201 | Y | C | F
J1510$+$5547 | SDSS J151020.06+554722.0 | 15 10 20.06 | +55 47 22.0 | 0.150 | Y | E | F
J1522$+$3934 | SDSS J152205.41+393441.3 | 15 22 05.41 | +39 34 41.3 | 0.077 | Y | I | F
J1536$+$5433 | Mrk 486 | 15 36 38.30 | +54 33 33.0 | 0.039 | N | I | $-$
J1537$+$4942 | SDSS J153732.61+494247.5 | 15 37 32.60 | +49 42 48.0 | 0.280 | N | C | $-$
J1548$+$3511 | SDSS J154817.92+351128.0 | 15 48 17.92 | +35 11 28.4 | 0.479 | Y | C | F
J1555$+$1911 | Mrk 291 | 15 55 07.90 | +19 11 33.0 | 0.035 | N | E | $-$
J1559$+$3501 | Mrk 493 | 15 59 09.60 | +35 01 47.0 | 0.031 | N | E | $-$
J1612$+$4219 | SDSS J161259.83+421940.3 | 16 12 59.83 | +42 19 40.0 | 0.233 | Y | E | F
J1629$+$4007 | SDSS J162901.30+400759.9 | 16 29 01.31 | +40 07 59.6 | 0.272 | Y | C | F
J1633$+$4718 | SDSS J163323.58+471858.9 | 16 33 23.58 | +47 18 59.0 | 0.116 | Y | I | F
J1634$+$4809 | SDSS J163401.94+480940.2 | 16 34 01.94 | +48 09 40.1 | 0.495 | Y | C | F
J1641$+$3454 | SDSS J164100.10+345452.7 | 16 41 00.10 | +34 54 52.7 | 0.164 | Y | E | F
J1703$+$4540 | SDSS J170330.38+454047.1 | 17 03 30.38 | +45 40 47.2 | 0.060 | Y | I | S
J1709$+$2348 | SDSS J170907.80+234837.7 | 17 09 07.82 | +23 48 38.2 | 0.254 | Y | C | S
J1713$+$3523 | FBQS J1713+3523 | 17 13 04.48 | +35 23 33.4 | 0.083 | Y | I | S
J2242$+$2943 | Ark 564 | 22 42 39.30 | +29 43 31.0 | 0.025 | N | E | $-$
J2314$+$2243 | RX J2314.9+2243 | 23 14 55.70 | +22 43 25.0 | 0.169 | Y | I | S

Table 8: [O III] and radio properties of the sample. Columns: (1) short name; (2) logarithm of the [O III] luminosity (erg s-1); (3) wavelength of the [O III] core component (Å); (4) velocity of the [O III] core component (km s-1); (5) FWHM of the [O III] core component (km s-1); (6) wavelength of the [O III] wing component (Å); (7) velocity of the [O III] wing component (km s-1); (8) FWHM of the [O III] wing component (km s-1); (9) logarithm of the integrated radio luminosity at 5 GHz (erg s-1); (10) logarithm of the peak luminosity at 5 GHz (erg s-1); (11) logarithm of the diffuse luminosity at 5 GHz (erg s-1). Short Name | $\log L_{[O~{}III]}$ | $\lambda_{c}$ |  vc |  FWHMc | $\lambda_{w}$ |  vw |  FWHMw | $\log L_{int}$ | $\log L_{peak}$ | $\log L_{diff}$
---|---|---|---|---|---|---|---|---|---|---
J0006+2012 | 41.47 |  5005.65$\pm$0.02 |  -71.63$\pm$1.03 |  458.28$\pm$4.77 |  5001.29$\pm$0.11 |  -332.20$\pm$6.84 |  1387.51$\pm$17.02 | 38.42 | 38.41 | 36.88
J0146-0040 | 41.16 |  5006.67$\pm$0.05 |  -10.36$\pm$2.99 |  143.83$\pm$4.19 |  5005.78$\pm$0.22 |  -53.34$\pm$14.37 |  575.32$\pm$19.16 | 37.95 | 37.94 | 36.40
J0347+0105 | 41.32 |  5004.23$\pm$0.12 |  -156.55$\pm$7.33 |  328.92$\pm$26.55 |  5000.64$\pm$0.15 |  -371.31$\pm$9.11 |  1074.27$\pm$27.42 | 39.17 | 39.01 | 38.66
J0629-0543 | 42.63 |  5007.91$\pm$0.25 |  63.91$\pm$15.15 |  621.47$\pm$60.71 |  4999.09$\pm$0.92 |  -463.97$\pm$55.21 |  1627.96$\pm$52.20 | 40.43 | 40.36 | 39.59
J0632+6340 | 40.28 |  5007.03$\pm$0.22 |  11.20$\pm$13.17 |  547.12$\pm$4.19 |  5002.18$\pm$3.92 |  -290.47$\pm$242.50 |  999.76$\pm$115.56 | 37.76 | 37.54 | 37.39
J0706+3901 | 41.33 |  5003.61$\pm$0.07 |  -193.58$\pm$4.19 |  736.07$\pm$2.99 | 0.00 | 0.00 | 0.00 | 39.33 | 39.23 | 38.63
J0713+3820 | 42.39 |  5011.86$\pm$0.14 |  300.4$\pm$8.38 |  504.81$\pm$12.57 |  4999.46$\pm$0.58 |  -741.81$\pm$41.31 |  1494.70$\pm$17.96 | 39.82 | 39.68 | 39.27
J0744+5149 | N.D. |  |  |  |  |  |  | 40.98 | 40.97 | 39.53
J0752+2617 | 41.00 |  5006.13$\pm$0.06 |  -42.69$\pm$3.59 |  227.02$\pm$2.40 |  5001.64$\pm$0.31 |  -268.66$\pm$20.36 |  720.56$\pm$8.98 | 39.19 | 38.17 | 36.64
J0758+3920 | 42.67 |  5005.40$\pm$0.07 |  -86.40$\pm$4.19 |  644.41$\pm$2.99 |  4993.81$\pm$0.40 |  -694.39$\pm$26.35 |  468.15$\pm$17.96 | 39.68 | 39.39 | 39.36
J0804+3853 | 42.19 |  5007.01$\pm$0.04 |  10.00$\pm$2.40 |  157.93$\pm$1.20 |  5000.52$\pm$0.15 |  -388.22$\pm$9.58 |  760.04$\pm$3.59 | 39.69 | 39.33 | 39.43
J0806+7248 | 41.81 |  5001.90$\pm$0.20 |  -295.97$\pm$11.98 |  375.08$\pm$8.98 |  4994.79$\pm$1.49 |  -426.03$\pm$99.40 |  733.25$\pm$61.67 | 40.17 | 40.15 | 38.93
J0814+5609 | 41.73 |  4997.51$\pm$0.51 |  -558.83$\pm$30.54 |  740.30$\pm$14.97 | 0.00 | 0.00 | 0.00 | 42.19 | 42.13 | 41.28
J0849+5108 | 41.74 |  5011.26$\pm$0.18 |  264.47$\pm$10.78 |  231.26$\pm$8.98 |  5002.93$\pm$1.09 |  -498.77$\pm$74.25 |  703.64$\pm$50.30 | 43.21 | 43.20 | 41.22
J0850+4626 | 42.32 |  5008.52$\pm$0.16 |  100.41$\pm$9.58 |  248.18$\pm$14.97 |  5005.76$\pm$0.52 |  -164.70$\pm$38.92 |  1558.15$\pm$28.74 | 41.77 | 41.76 | 39.98
J0902+0443 | 42.20 |  5007.75$\pm$0.27 |  54.31$\pm$16.17 |  286.25$\pm$26.35 |  4999.03$\pm$0.66 |  -521.95$\pm$54.49 |  1231.01$\pm$20.36 | 42.77 | 42.77 | 0.00
J0913+3658 | 41.43 |  5005.90$\pm$0.08 |  -56.46$\pm$4.79 |  329.96$\pm$5.39 |  5000.78$\pm$0.32 |  -306.62$\pm$22.15 |  799.52$\pm$7.78 | 38.66 | 38.61 | 37.74
J0925+5217 | 41.99 |  5007.78$\pm$0.53 |  56.1$\pm$31.73 |  215.74$\pm$1.80 | 0.00 | 0.00 | 0.00 | 38.72 | 38.42 | 38.42
J0926+1244 | 41.14 |  5006.32$\pm$0.61 |  -31.32$\pm$36.52 |  337.01$\pm$31.73 |  5005.00$\pm$0.96 |  -79.02$\pm$92.21 |  846.06$\pm$32.93 | 38.52 | 38.45 | 37.70
J0937+3615 | 41.76 |  5006.98$\pm$0.14 |  8.20$\pm$8.38 |  404.70$\pm$11.98 |  5002.80$\pm$0.69 |  -249.94$\pm$48.50 |  999.76$\pm$25.15 | 39.71 | 39.58 | 39.11
J0945+1915 | N.D. |  |  |  |  |  |  | 40.90 | 40.88 | 39.52
J0948+5029 | 41.48 |  5006.74$\pm$0.05 |  -6.17$\pm$2.99 |  337.01$\pm$1.80 |  5001.76$\pm$0.17 |  -298.60$\pm$11.38 |  868.62$\pm$5.39 | 38.91 | 38.89 | 37.52
J0952-0136 | 41.67 |  5005.96$\pm$0.08 |  -52.87$\pm$4.79 |  558.40$\pm$4.19 |  4994.51$\pm$0.52 |  -685.82$\pm$33.53 |  1583.53$\pm$13.77 | 39.04 | 38.99 | 38.02
J0957+2433 | 41.36 |  5005.64$\pm$0.06 |  -72.03$\pm$3.59 |  327.14$\pm$2.40 |  5001.11$\pm$0.27 |  -271.66$\pm$17.96 |  958.86$\pm$9.58 | 38.32 | 38.31 | 36.73
J1031+4234 | 41.83 |  5009.07$\pm$0.30 |  133.35$\pm$17.96 |  231.26$\pm$10.78 |  5005.45$\pm$0.92 |  -216.94$\pm$65.86 |  744.53$\pm$10.18 | 41.14 | 41.13 | 39.55
J1034+3938 | 41.03 |  5006.76$\pm$0.04 |  -4.97$\pm$2.40 |  228.44$\pm$1.80 |  5002.54$\pm$0.17 |  -253.01$\pm$10.78 |  901.05$\pm$4.79 | 39.24 | 39.18 | 38.34
J1037+0036 | 41.60 |  5008.56$\pm$0.42 |  102.81$\pm$25.15 |  280.61$\pm$7.78 |  4999.72$\pm$0.58 |  -529.11$\pm$58.08 |  356.75$\pm$22.15 | 42.19 | 42.19 | 0.00
J1038+4227 | 41.62 |  5005.11$\pm$0.12 |  -103.77$\pm$7.19 |  314.45$\pm$7.78 |  4999.87$\pm$0.52 |  -313.96$\pm$36.52 |  1146.41$\pm$28.14 | 40.67 | 40.16 | 40.51
J1047+4725 | 42.89 |  5006.66$\pm$0.07 |  -10.96$\pm$4.19 |  280.61$\pm$4.79 |  5004.82$\pm$0.39 |  -110.43$\pm$25.75 |  844.65$\pm$24.55 | 43.77 | 43.51 | 43.42
J1048+2222 | 41.79 |  5003.35$\pm$0.31 |  -209.15$\pm$18.56 |  389.19$\pm$19.76 |  4994.49$\pm$0.75 |  -530.88$\pm$61.07 |  1238.06$\pm$16.77 | 39.75 | 39.72 | 38.59
J1102+2239 | 42.38 |  4997.40$\pm$0.36 |  -565.41$\pm$21.56 |  879.90$\pm$11.38 |  4984.90$\pm$2.41 |  -749.79$\pm$164.06 |  1319.85$\pm$83.23 | 40.47 | 40.44 | 39.35
J1110+3653 | 41.73 |  5007.53$\pm$0.22 |  41.14$\pm$13.17 |  400.47$\pm$16.77 | 0.00 | 0.00 | 0.00 | 41.89 | 41.83 | 40.98
J1114+3241 | 42.13 |  4993.71$\pm$0.08 |  -786.31$\pm$5.06 |  971.78$\pm$22.39 |  4979.31$\pm$0.67 |  -1648.58$\pm$40.06 |  2101.97$\pm$36.32 | 41.04 | 41.03 | 39.35
J1121+5351 | 42.01 |  5007.09$\pm$0.06 |  14.79$\pm$3.59 |  305.99$\pm$2.99 |  5003.61$\pm$0.15 |  -208.07$\pm$10.78 |  734.66$\pm$4.79 | 39.25 | 39.17 | 38.45
J1138+3653 | 41.68 |  5005.39$\pm$0.17 |  -87.00$\pm$10.18 |  234.08$\pm$16.17 |  5004.59$\pm$0.43 |  -48.40$\pm$34.73 |  923.61$\pm$28.74 | 41.02 | 41.00 | 39.77
J1146+3236 | 42.00 |  5007.10$\pm$0.10 |  15.39$\pm$5.99 |  307.40$\pm$4.19 |  5001.47$\pm$0.40 |  -336.86$\pm$28.74 |  620.44$\pm$7.78 | 41.90 | 41.89 | 39.94
J1159+2838 | 42.36 |  5007.75$\pm$0.04 |  54.32$\pm$2.17 |  397.77$\pm$7.33 |  5002.81$\pm$0.05 |  -241.67$\pm$2.80 |  1201.68$\pm$5.55 | 39.76 | 39.74 | 38.51
J1203+4431 | 39.94 |  5006.94$\pm$0.45 |  5.81$\pm$26.94 |  318.68$\pm$5.99 |  5002.02$\pm$0.91 |  -294.61$\pm$59.88 |  606.34$\pm$81.43 | 36.46 | 35.93 | 36.31
J1209+3217 | 41.74 |  5006.17$\pm$0.19 |  -40.30$\pm$11.38 |  410.34$\pm$6.59 |  5002.46$\pm$0.24 |  -221.92$\pm$23.95 |  1405.86$\pm$12.57 | 39.31 | 39.26 | 38.35
J1215+5442 | 41.78 |  5005.69$\pm$0.16 |  -69.04$\pm$9.58 |  384.96$\pm$7.19 |  4997.34$\pm$0.58 |  -500.47$\pm$42.51 |  661.33$\pm$15.57 | 39.30 | 39.25 | 38.34
J1218+2948 | 41.26 |  5007.70$\pm$0.38 |  51.31$\pm$22.75 |  341.24$\pm$2.40 |  5005.81$\pm$0.47 |  -113.22$\pm$50.90 |  774.14$\pm$6.59 | 38.47 | 38.34 | 37.87
J1227+3214 | 42.11 |  5007.84$\pm$0.05 |  59.70$\pm$2.99 |  242.54$\pm$1.80 |  5004.35$\pm$0.12 |  -209.05$\pm$8.38 |  592.24$\pm$3.59 | 39.94 | 39.90 | 38.92
J1238+3942 | 41.98 |  5005.75$\pm$0.15 |  -65.45$\pm$8.98 |  380.73$\pm$15.57 | 0.00 | 0.00 | 0.00 | 42.16 | 42.16 | 0.00
J1242+3317 | 41.65 |  5006.86$\pm$0.04 |  1.02$\pm$2.40 |  307.40$\pm$1.20 |  5003.30$\pm$0.10 |  -213.29$\pm$6.59 |  888.36$\pm$2.99 | 38.74 | 38.51 | 38.35
J1246+0222 | 41.12 |  5006.03$\pm$0.07 |  -48.68$\pm$4.19 |  310.22$\pm$3.59 |  5002.24$\pm$0.49 |  -227.11$\pm$31.73 |  738.89$\pm$13.77 | 38.31 | 38.12 | 37.87
J1246+0238 | 41.64 |  5007.51$\pm$0.35 |  39.94$\pm$20.96 |  382.14$\pm$20.36 |  4999.10$\pm$1.96 |  -503.49$\pm$136.52 |  1245.11$\pm$70.06 | 41.69 | 41.67 | 40.20
J1302+1624 | 42.31 |  5003.65$\pm$0.05 |  -191.19$\pm$2.99 |  408.93$\pm$1.20 | 0.00 | 0.00 | 0.00 | 39.91 | 39.28 | 39.80
J1305+5116 | 43.14 |  5003.62$\pm$1.36 |  -192.98$\pm$81.43 |  1012.45$\pm$77.84 |  4993.35$\pm$2.45 |  -615.00$\pm$226.33 |  1394.58$\pm$43.11 | 42.91 | 42.71 | 42.47
J1317+6010 | 41.74 |  5007.03$\pm$0.03 |  11.31$\pm$1.55 |  341.35$\pm$6.69 |  5000.55$\pm$0.34 |  -376.82$\pm$20.53 |  1435.93$\pm$33.88 | 39.42 | 39.28 | 38.86
J1333+4141 | 42.34 |  5005.35$\pm$0.08 |  -89.4$\pm$4.79 |  521.73$\pm$3.59 |  5001.35$\pm$0.11 |  -239.37$\pm$10.18 |  1745.69$\pm$5.39 | 39.85 | 39.84 | 38.21
J1337+2423 | 42.24 |  5005.74$\pm$0.25 |  -66.04$\pm$14.97 |  352.52$\pm$16.17 |  4995.76$\pm$1.29 |  -597.81$\pm$90.41 |  2240.64$\pm$25.75 | 40.26 | 40.23 | 39.18
J1346+3121 | 41.39 |  5005.55$\pm$0.51 |  -77.42$\pm$30.54 |  248.18$\pm$55.09 |  5001.72$\pm$2.60 |  -229.05$\pm$184.42 |  920.79$\pm$102.39 | 40.63 | 40.63 | 0.00
J1355+5612 | 42.22 |  5007.64$\pm$0.04 |  47.72$\pm$2.40 |  407.52$\pm$1.20 |  5003.83$\pm$0.10 |  -228.43$\pm$7.19 |  1013.86$\pm$3.59 | 39.67 | 39.65 | 38.15
J1358+2658 | 42.36 |  5007.86$\pm$0.05 |  60.89$\pm$2.99 |  448.41$\pm$3.59 |  5003.64$\pm$0.24 |  -252.57$\pm$16.17 |  1366.38$\pm$5.39 | 40.05 | 39.99 | 39.16
J1402+2159 | 41.49 |  5006.18$\pm$0.05 |  \- 39.70$\pm$2.99 |  291.89$\pm$1.80 |  5003.45$\pm$0.13 |  -163.17$\pm$8.98 |  762.86$\pm$5.39 | 38.34 | 38.24 | 37.67
J1421+2824 | N.D. |  |  |  |  |  |  | 42.20 | 42.18 | 40.66
J1443+4725 | 42.32 |  4999.40$\pm$0.54 |  -445.66$\pm$32.33 |  1199.99$\pm$33.53 | 0.00 | 0.00 | 0.00 | 42.98 | 42.97 | 41.43
J1505+0326 | 41.77 |  5000.96$\pm$0.99 |  -352.25$\pm$59.28 |  597.88$\pm$60.48 |  4995.52$\pm$4.34 |  -326.51$\pm$317.35 |  1180.25$\pm$131.13 | 43.09 | 43.09 | 0.00
J1536+5433 | 41.06 |  5007.23$\pm$0.75 |  23.17$\pm$44.91 |  455.46$\pm$16.17 | 0.00 | 0.00 | 0.00 | 37.98 | 37.92 | 37.13
J1537+4942 | 41.67 |  5006.48$\pm$0.08 |  -21.74$\pm$4.79 |  153.70$\pm$6.59 |  5000.77$\pm$0.37 |  -341.82$\pm$25.15 |  837.60$\pm$16.17 | 40.09 | 40.09 | 0.00
J1548+3511 | 42.40 |  5006.63$\pm$0.07 |  -12.75$\pm$4.19 |  265.10$\pm$2.40 |  4999.95$\pm$0.59 |  -400.03$\pm$37.72 |  971.55$\pm$27.54 | 37.85 | 37.23 | 37.73
J1555+1911 | 40.82 |  5005.91$\pm$0.43 |  -55.86$\pm$25.75 |  349.70$\pm$4.79 | 0.00 | 0.00 | 0.00 | 38.19 | 37.75 | 38.00
J1559+3501 | 40.58 |  5005.48$\pm$0.07 |  -81.61$\pm$4.19 |  293.30$\pm$3.59 |  4999.65$\pm$0.38 |  -349.15$\pm$25.15 |  806.57$\pm$14.97 | 42.46 | 42.46 | 0.00
J1612+4219 | 42.31 |  5004.23$\pm$0.14 |  -156.46$\pm$8.38 |  741.71$\pm$12.57 |  5005.96$\pm$0.36 |  103.66$\pm$28.14 |  1913.50$\pm$20.96 | 39.86 | 39.45 | 39.65
J1629+4007 | 42.06 |  5009.08$\pm$0.06 |  133.94$\pm$3.59 |  211.51$\pm$2.40 |  5004.52$\pm$0.34 |  -272.60$\pm$22.15 |  830.54$\pm$19.16 | 41.71 | 41.70 | 39.06
J1633+4718 | 41.65 |  5006.57$\pm$0.05 |  -16.35$\pm$2.99 |  255.23$\pm$1.80 |  5001.91$\pm$0.23 |  -278.77$\pm$14.97 |  1185.89$\pm$13.17 | 40.67 | 40.64 | 39.46
J1634+4809 | 41.96 |  5002.47$\pm$0.44 |  -261.84$\pm$26.35 |  389.19$\pm$24.55 |  4993.44$\pm$1.69 |  -541.39$\pm$125.14 |  1109.74$\pm$56.28 | 41.42 | 41.40 | 40.03
J1703+4540 | 41.55 |  5010.34$\pm$1.23 |  209.39$\pm$73.65 |  335.60$\pm$49.10 |  5004.73$\pm$3.96 |  -335.68$\pm$309.56 |  645.82$\pm$110.77 | 40.20 | 40.16 | 39.04
J1709+2348 | 42.14 |  5008.43$\pm$0.09 |  95.02$\pm$5.39 |  401.88$\pm$3.59 |  5003.06$\pm$0.39 |  -321.53$\pm$26.94 |  786.83$\pm$9.58 | 39.55 | 39.55 | 0.00
J1713+3523 | 40.98 |  4995.59$\pm$0.88 |  -673.79$\pm$52.69 |  1641.35$\pm$47.90 | 0.00 | 0.00 | 0.00 | 39.54 | 39.51 | 38.33
J2242+2943 | 41.34 |  5006.77$\pm$0.01 |  -4.46$\pm$0.72 |  226.27$\pm$3.75 |  4997.42$\pm$1.64 |  -564.31$\pm$98.63 |  1089.89$\pm$279.06 | 38.88 | 38.62 | 38.54
J2314+2243 | 42.31 |  5006.86$\pm$0.81 |  1.02$\pm$48.50 |  602.11$\pm$25.75 |  4989.64$\pm$2.20 |  -1031.42$\pm$178.43 |  1579.30$\pm$55.69 | 40.48 | 40.42 | 39.58
J1029+5556 | 41.82 |  5006.96$\pm$0.22 |  6.97$\pm$13.39 |  335.37$\pm$37.37 |  5004.16$\pm$1.65 |  -160.73$\pm$98.89 |  2004.90$\pm$387.54 | 38.46 | 38.46 | 0.00
J1228+5017 | 41.83 |  5008.56$\pm$0.03 |  102.76$\pm$1.71 |  218.21$\pm$5.78 |  5003.58$\pm$0.27 |  -195.21$\pm$16.48 |  978.24$\pm$47.52 | 39.00 | 38.91 | 38.27
J1232+4957 | 41.66 |  5003.32$\pm$1.28 |  -211.09$\pm$76.44 |  578.97$\pm$156.31 |  4994.82$\pm$4.45 |  -719.77$\pm$266.80 |  1142.04$\pm$297.77 | 38.23 | 38.23 | 0.00
J1509+6137 | 40.94 |  5003.87$\pm$0.12 |  -178.29$\pm$7.34 |  320.40$\pm$24.24 |  5005.87$\pm$1.91 |  -58.51$\pm$114.53 |  1690.48$\pm$189.36 | 37.68 | 37.68 | 0.00
J1510+5547 | 41.29 |  5006.85$\pm$0.03 |  0.67$\pm$1.56 |  123.77$\pm$8.75 |  5004.09$\pm$0.29 |  -164.99$\pm$17.14 |  547.56$\pm$34.91 | 37.69 | 37.48 | 37.27
J1522+3934 | 40.30 |  5007.43$\pm$0.36 |  35.10$\pm$21.3 |  321.12$\pm$79.06 |  5001.67$\pm$3.29 |  -310.03$\pm$197.07 |  1306.38$\pm$324.23 | 37.95 | 37.88 | 37.12
J1641+3454 | 41.45 |  5006.61$\pm$0.07 |  -14.01$\pm$4.46 |  448.44$\pm$19.38 |  5006.19$\pm$0.62 |  -39.20$\pm$37.30 |  1274.45$\pm$136.69 | 38.92 | 38.79 | 38.33