JVLA 1.4 GHz flux density, suggesting strong star formation activity. Also
$S_{\textrm{5.2GHz, CDFS}}$ is $\sim$0.8 mJy higher than $S_{\textrm{int}}$,
supporting star formation as the main producer of radio emission in this
galaxy. However, the contribution of AGN-heated dust to the mid-infrared
emission cannot be ruled out, especially since an [O III] wing is present,
with a velocity of -500 km s-1 and a FWHM of 661 km s-1.

####  5.2.26 J1218+2948

J1218+2948 (Mrk 766) is a well-studied nearby ($z$ = 0.013) radio-loud NLS1
galaxy. Its peak flux density is 11.17 mJy beam-1, the integrated flux density
15.46 mJy, and the flux densities of the tapered maps almost the same. Its
luminosity is quite low with log $\nu L_{\nu,\textrm{int}}$ = 38.47 erg s-1.
The radio morphology (Fig. 31(a)) is mostly featureless, but exhibits a slight
elongation in south-east - north-west direction. The tapered maps (Figs. 32(c)
and 32(d)) reveal additional extended emission toward east. The radio emission
is confined within the central parts of a barred spiral host (Fig. 32(e) and
32(f)). The $\alpha$ map shows a steep core, with a spectral index of -0.84.
The total spectral index is -0.83, so the source is thoroughly steep.

In archival radio observations obtained with the VLA the flux densities were
found to be 35.9 mJy at 1.46 GHz, and 15.4 mJy at 4.89 GHz, in OVRO
observations at 20 GHz the flux density was 5.3 mJy (Edelson 1987), and
observations with Arecibo yielded 29 mJy at 2.38 GHz (Dressel & Condon 1978).
Bicay et al. (1995) obtained 17 $\pm$ 4 mJy at 4.755 GHz with the GBT, and
Parra et al. (2010) 14.5 mJy at 4.8 GHz with the VLA. All the archival values
are consistent with what we obtain, and indicate that the emission is not
variable, at least at these time-scales. Using the archival and our data the
spectral index between 1.46 and 5.2 GHz is -0.66, and between 5.2 and 20 GHz
-0.79, consistent with our $\alpha$ map. J1218+2948 has been unresolved in
most radio imaging observations, but at 8.4 GHz with the VLA in A
configuration a slight extension toward south-east was seen (Kukula et al.
1995), and 1.5 GHz VLA observations also detected extended emission on the
north-west side of the nucleus, in addition to the south-east emission, was
detected (Nagar et al. 1999). These structures correspond to the slightly
elongated morphology we see in J1218+2948. The source was not detected at 22
GHz with the JVN (Doi et al. 2016), but was detected at 95 GHz with the
Combined Array for Research in Millimeter-wave Astronomy (CARMA) with a flux
density of 1.98 mJy (Behar et al. 2015). Behar et al. (2015) note that the
high frequency emission follows the $L_{\textrm{R}}$ \- $L_{\textrm{X}}$
relation found in stellar coronae, and argue that the high frequency radio
emission is similarly produced in radio-quiet and moderately radio-loud AGN in
the accretion disk corona closer to the black hole and within a smaller region
than the large-scale radio emission. Esparza-Arredondo et al. (2020) classify
this source as a fading nucleus.

Structures resembling the radio morphology have also been observed in optical
and IR emission lines. The [O III] emission is found to be rather circular
(Mulchaey et al. 1996; Schmitt et al. 2003) and limited to the innermost
$\sim$0.5 kpc of the galaxy, but an excitation map shows a biconical south-
east – north-west structure, with the south-east emission being brighter.
Gonzalez Delgado & Perez (1996) find that the H$\alpha$ emission is extended
1.5 kpc toward north-east, and conclude that the spectrum of this region can
be fitted with H II region models. They also find both a blue-shifted and a
red-shifted component of [O III] on opposite sides of the nucleus, and argue
that the kinematics of this gas is a result of a nuclear outflow. Schönell et
al. (2014) mapped the innermost 450 pc of the galaxy with the Gemini Near
Infrared Integral Field Spectrograph in $J$ and $Ks$ bands, and found that
also the [Fe II] emission is extended toward south-east. Based on line ratios
they deduce that they have been ionised by a mixture of nuclear and starburst
emission. Toward south-east they detect signs of shocked gas and also an
increase in the [Fe II] velocity distribution, as well as red- and blue-
shifted components they think are caused by a jet/ourflow. Also a water maser
has been detected in J1218+2948, but whether it is associated with the nucleus
is unclear (Tarchi et al. 2011).

Star formation diagnostics for this source can be unreliable due to the JVLA
observation not mapping the whole galaxy, so the flux density can be somewhat
underestimated. Furthermore, Rodríguez-Ardila et al. (2005) found in the near-
infrared spectrum an excess they attribute to hot dust with a blackbody
temperature of 1200 K they think resides very close to the nucleus, at the
dust sublimation radius. If also the toroidal/poloidal cooler dust is bright
it can contribute to the observed mid-infrared flux densities. Keeping these
in mind, the star formation diagnostics point to this source being star-
forming. The W3-W4 is 2.87, $q22$ is 1.40, the W3 flux density is 269.8 mJy
higher than $S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{5.2GHz, CDFS}}$ is
4.6 mJy higher than $S_{\textrm{int}}$. The star formation of the host is not
extraordinarily high: based on far-infrared observations it is 5.42
$M_{\odot}$ yr-1 (Terrazas et al. 2017), and based on the infrared luminosity
it is 2.24 $M_{\odot}$ yr-1 (Lamperti et al. 2020). However, Rodríguez-Ardila
& Viegas (2003) detect a 3.3 $\mu$m PAH feature within 150 pc from the nucleus
at a level seen in circumnuclear starbursts in other sources. Also ultraviolet
(UV) observations reveal several bright star-forming regions and star
clusters, and the most pronounced star forming region seems to be associated
with the east part of the galaxy bar (Muñoz Marín et al. 2007). This could
explain the east-ward extended emission seen in the tapered maps.

Interestingly two QPOs with a frequency ratio of 3:2 have been reported in
J1218+2948 in X-rays (Boller et al. 2001; Zhang et al. 2017a). The periods of
the QPOs are 4200 and 6450 s, and seem to be transient. Also ultrafast
outflows (UFO) have been detected in this source, with velocities around
$\sim$0.1 c (Igo et al. 2020, and references therein).

All in all, it seems like the radio emission of J1218+2948 is a mixture of
star formation and AGN related emission. In the radio maps there are signs of
possible jets or outflows in the south-east - north-west direction, but due to
the lack of variability in the radio flux densities or radio morphology the
jets/outflows appear lethargic. Furthermore, according to several studies the
host galaxy is forming stars, but the level of this activity, and thus its
contribution to the radio emission, is uncertain. Further observations would
be needed to clarify the role of each component in this galaxy.

####  5.2.27 J1227+3214

J1227+3214 (SDSS J122749.14+321458.9) is a radio-loud NLS1 ($z$ = 0.136) with
a peak flux density of 2.96 mJy beam-1, the integrated flux density of 3.34
mJy, and the flux densities of the tapered maps only slightly higher. This
results in a moderate luminosity of log $\nu L_{\nu,\textrm{int}}$ = 39.94 erg
s-1. The source was detected in FIRST with a flux density of 6.42 mJy,
yielding a spectral index of -0.50 between 1.4 and 5.2 GHz. Our $\alpha$ map
in Fig. 33(a) shows a steeper spectrum: the total spectral index is -0.67, and
the core spectral index is -0.74. The difference might be due to variability
since the mean epoch of the FIRST observations is in 1994, however, due to the
lack of further radio observations we cannot investigate the possible
variability. The radio morphology seems somewhat resolved, but without any
clear structures. In the normal map the morphology seems slightly extended
toward north-east and west, whereas the tapered maps reveal faint extended
emission toward north and south-east. Whereas the bulk of the emission is
confined within the host galaxy, this extended emission seems to be reaching
outside it (Fig. 33(e)).

The star formation diagnostics indicate some activity in the host: the W3-W4
colour is 2.62, $q22$ is 1.06, and the W3 flux density is $\sim$25.4 mJy
higher than the $S_{\textrm{1.4GHz, JVLA}}$. However, $S_{\textrm{int}}$ is
1.9 mJy higher than $S_{\textrm{5.2GHz, CDFS}}$. Caccianiga et al. (2015) find
the star formation of this source to be high, 54 $M_{\odot}$ yr-1, so star
formation contribution to the radio emission can be expected. Interestingly,
this source was labelled as a red quasar in LaMassa et al. (2016). They argue
that it is in a dusty, possibly post-merger state, when powerful nuclear winds
start unveiling a heavily obscured nucleus. However, no shift was detected in
[O III] lines so a large-scale outflow does not seem to be present, and only a
moderate wing component with a velocity of -209 km s-1 and a FWHM of 592 km
s-1 was seen in Berton & Järvelä (2021). Also Zhang et al. (2017b) studied its
dust properties, and conclude that only the innermost parts of the AGN – the
accretion disk, the broad-line region, and the inner edge of the torus – seem
to be obscured, and not the narrow-line region. Based on Cloudy analysis they
claim that the obscuring medium can actually be the torus itself. The host
galaxy of J1227+3214 has been observed a few times, but despite its moderate
redshift, the attemps to model the morphology of the host galaxy have been
inconclusive, thus nothing can be said about the state of the host (Järvelä et
al. 2018; Olguín-Iglesias et al. 2020).

J1227+3214 would be very luminous for a galaxy dominated purely by star
formation, so it is probable that also the AGN contributes to the overall
radio emission. Based on its luminosity and the steep spectral index, it could
be a low-luminosity CSS-like source. With these data it is impossible to say
whether the extended emission that reaches outside the host is of nuclear or
star formation origin, and further observations will be needed to reveal the
nature of this source.

####  5.2.28 J1242+3317

J1242+3317 (WAS 61) is a radio-quiet source at $z$ = 0.044 with a peak flux
density of 1.32 mJy beam-1, integrated flux density of 2.23 mJy, and the flux
densities of the tapered maps similar. Its luminosity is quite low with log
$\nu L_{\nu,\textrm{int}}$ = 38.74 erg s-1. Its FIRST flux density is 6.32
mJy, giving a spectral index of -0.79 between 1.4 and 5.2 GHz. This is
consistent with what we see in the $\alpha$ map in Fig. 34(a): the total
spectral index is -0.74, and the core spectral index is -0.88. Brinkmann et
al. (1995) report a value of 45 mJy at 4.85 GHz observed with the GBT in 1989,
but this record is not found anywhere else, for example, in the NASA/IPAC
Extragalactic Database (NED), the High Energy Astrophysics Science Archive
Research Center’s (HEASARC) Xamin, or SIMBAD’s VizieR, so the value is highly
dubious. The morphology of the source is clearly elongated toward east, and it
seems like the spectral index flattens toward east as well, but this might be
an edge-effect, especially since also the errors increase toward this
direction. The tapered maps (Figs. 34(c) and 34(d)) show the same extended
emission toward east, and are additionally also slightly extended toward
south-west. The maximum extent of the region of continuous emission is 1.3 kpc
in the normal map and 2.9 kpc in the 90k$\lambda$ tapered map. The extended
emission can be caused by either a low-power jet or a nuclear outflow. Higher-
resolution observations would be needed to distinguish between these two
options.

The host galaxy of J1242+3317 was classified as a possible S0 galaxy with a
strong bar (SB) (Ohta et al. 2007), but modelling the host with a bulge and a
disk yielded a Sérsic index of 3.45 for the bulge, suggesting it is a
classical bulge (Mathur et al. 2012), in which case the galaxy would not be a
pristine late-type galaxy. However, the presence of a bar was not taken into
account in the modelling, even if it is clearly visible in the residuals, and
might affect the final results.

J1242+3317 is also a bright mid-infrared emitter, even though it is quite
clear that the morphology is caused by an AGN. Its W3-W4 colour is 2.72, $q22$
is 1.59, the W3 flux density is 60.8 mJy higher than $S_{\textrm{1.4GHz,
JVLA}}$, and $S_{\textrm{5.2GHz, CDFS}}$ is 1.6 mJy higher than
$S_{\textrm{int}}$. Yang et al. (2020) estimated its star formation rate based
on the 5.2 GHz radio luminosity ad concluded that the radio emission cannot be
of star formation origin. However, they do not consider the possibility that
it can be a mixture of the two. Using BPT diagrams Stern & Laor (2013)
classified J1242+3317 as an AGN, a star-forming galaxy, and a Seyfert galaxy
in diagrams based on [N II], [S II], and [O I], respectively. This result
suggests that indeed probably also star formation is present in this source.
Also an [O III] wing is detectable in the spectrum, with a velocity of -213 km
s-1 and a FWHM of 888 km s-1, so additional contribution of the AGN-heated
dust to the mid-infrared emission cannot be ruled out either.

####  5.2.29 J1246+0222

J1246+0222 (PG 1244+026) is a radio-quiet, faint NLS1, at $z$ = 0.048, with a
peak flux density of 0.43 mJy beam-1, an integrated flux density of 0.71 mJy
and the flux densities of the tapered maps of 0.85 and 0.82 mJy for the
90k$\lambda$ and the 60k$\lambda$ taper, respectively. It is also a low-
luminosity source with log $\nu L_{\nu,\textrm{int}}$ = 38.31 erg s-1. In
Kellermann et al. (1994) its 5 GHz flux density was found to be 0.83 mJy, so
it seems like it has not undergone major changes during the last decades. Its
FIRST flux density is 2.23 mJy, giving a spectral index of -0.87 between 1.4
and 5.2 GHz. Upgraded Giant Metrewave Radio Telescope (uGMRT) observations at
685 MHz yielded a flux density of 3.1 $\pm$ 0.3 mJy, and they found a steep
spectral index of -0.7 (Silpa et al. 2020). Furthermore, they found the source
to be extended toward west - south-west on a kpc-scale. In the radio map in
Fig. 35(a) the source seems to be extended northward and west - southwestward.
The spectral index varies over the emitting region, and the extreme values
close to the edges are clearly a product of edge-effects and low S/N ratio
data. The core spectral index is flat with $\alpha$ = 0.05, consistent with
what was found in B18, although new observations are certainly required to
confirm our result. The tapered maps show extended emission toward south-east,
with an extent of 5.2 kpc in the 90k$\lambda$ tapered map. The host galaxy of
J1246+0222 is featureless (Fig. 35(e)), and the bulk of the radio emission is
confined within it. Its SDSS $r$ band 25 mag arcsec-2 isophotal semimajor axis
is 3.42 kpc, implying that the extended emission reaches outside the host.

Teng et al. (2013) studied the H I absorption properties of J1246+0222 and
concluded that the blueshifted and broad H I absorption features match an
outflow caused by the nuclear activity, probably driven by a jet. Moreover,
they found that the minima of the absorption troughs exceed the nuclear
continuum flux densities, meaning that there is also off-nuclear radio
emission contributing to the continuum in J1246+0222. Based on the resolution
of their observations they argue that the emitting region needs to be at least
3 kpc from the core. The south-eastward extended emission visible in our
tapered maps would fit the criterion of the required off-nuclear emission.
Also fitting of the timing data of J1246+0222 in X-rays supports the presence
of a jet or at least a jet base in this source. Chainakun & Young (2017) found
that the data is best fitted with two X-ray sources, with the source further
away from the black hole producing only small amounts of reflection X-rays.
They interpret this component as a jet whose emission is beamed away from the
accretion disk and thus does not get reflected. In any case, the radio
morphology would be highly unusual for a star-forming galaxy, and thus it is
probable that at least the extended emission is produced by nuclear activity,
which implies that a jet or a nuclear outflow is or has been present.
Interestingly, no [O III] shift is seen in the optical spectrum, but a wing is
present with a velocity of -503 km s-1 and a FWHM of 1245 km s-1. Also a
possible ultrafast outflow with a velocity of 0.08 c has been detected (Igo et
al. 2020).

According to the star formation diagnostics the bulk of the radio emission
could be from star formation as the W3-W4 colour is 2.56, $q22$ is 2.26, the
W3 flux density is 40.3 mJy higher than the $S_{\textrm{1.4GHz, JVLA}}$, and
$S_{\textrm{5.2GHz, CDFS}}$ is about the same as $S_{\textrm{int}}$. The star
formation rates derived using different methods are roughly in agreement,
ranging from 1 to 11 $M_{\odot}$ yr-1 (Young et al. 2014; Sargsyan et al.
2012; Shangguan et al. 2020). However, Yang et al. (2020) claim that producing
the amount of radio emission seen in J1246+0222 would require very improbable
amounts of star formation. Shi et al. (2007) claim to have detected PAH
features at 7.7 and 11.3$\mu$m in the infrared spectrum obtained with the
Infrared Spectrograph (IRS) onboard Spitzer Space Telescope, whereas Xie et
al. (2018) see only a flat, featureless spectrum in their data obtained with
the same instrument. They model the spectrum and conclude that it can be
reproduced by three continuum components with temperatures of about 394, 165,
and 77 K. If in absence of strong star formation the hot, as well as partly
the warm, dust emission could originate from the dust heated by the AGN, which
could explain the excess mid-infrared emission.

####  5.2.30 J1302+1624

J1302+1624 (Mrk 783) is a radio-loud NLS1 at $z$ = 0.067 with a peak flux
density of 3.28 mJy beam-1, an integrated flux density of 11.47 mJy and the
flux densities of the tapered maps about half a mJy higher. Its luminosity is
log $\nu L_{\nu,\textrm{int}}$ = 39.91 erg s-1. Its radio map in Fig. 36(a)
reveals a core and considerable extended emission toward east and south-east,
in a curved morphology. The tapered maps do not show any additional emission
(Figs. 36(c) and 36(d)). The extent of the emission from the core in the
normal map toward south-east is 7 kpc, toward north 3.9 kpc, and toward west 2
kpc. In the $\alpha$ map (Fig. 36(a)) the spectral index of the core is
$\sim$-0.83, and the total spectral index is -1.50. The extended emission is
mostly very steep, with spectral indices $<$ -1.5. An exception is a region
very south, at the tip of the extended emission: the spectral index at the
centre of the highest contour (RA 13:02:58.04,Dec +16:24:24.49) is -0.67 , and
thus considerably flatter than elsewhere. It is probable that these flatter
spectral indices are real, since they are not only seen at edges, and this
region is clearly brighter than the rest of the jet, resembling a hotspot, and
indicating interaction with the interstellar medium.

J1302+1624 was detected with FIRST with a flux density of 28.72 mJy, but this
value includes both the core and the extended emission, and thus deriving a
spectral index using it does not tell us much. Doi et al. (2013) observed the
source with the VLBA at 1.7 GHz finding it unresolved with a flux density of
1.3 mJy. It was not detected at 22 GHz using the JVN (Doi et al. 2016).
J1302+1624 was extensively studied in Congiu et al. (2017a), Congiu et al.
(2017b), and Congiu et al. (2020). Congiu et al. (2017b) found an extended
narrow-line region in this source. This region has the same position angle as
the radio emission extended south-east, but is much further away from the
nucleus, reaching as far as 35 kpc. They confirmed with a BPT diagram that
this region has been ionised by the AGN, and not star formation. Hints of gas
ionised by star formation are seen closer to the nucleus, and between these
there is a region where no line emission was detected at all. If the gas 35
kpc away from the nucleus was ionised by the AGN, we can use a simple light
travel time argument to estimate a lower limit for the age of the source: it
has taken the nuclear emission at least 35 kpc / $c$ = 105.05 yr to reach this
region.

Congiu et al. (2017a) thoroughly analysed the same JVLA data we are using.
Using tapering they estimated the extent of the emission to be 14 kpc south-
eastward and 12 kpc north-westward. They measured the spectral indices using
the same method as in B18 and estimate the core spectral index to be -0.67
$\pm$ 0.13, and the spectral index of the extended emission to be -2.02 $\pm$
0.74. We confirm these results with the $\alpha$ map. It should be noted that
their extended spectral index includes both the very steep and the flatter
region seen in our $\alpha$ map, thus being an average between them. Due to
the very steep spectral index and the lack of clear collimated structures they
hypothesise that the extended emission might be old, and is not being
replenished anymore. They also propose that the S-shaped extended emission in
the tapered map might suggest that the jets in J1302+1624 are precessing. They
tested this scenario further in Congiu et al. (2020) observing the source with
the VLBA and the enhanced Multi Element Remotely Linked Interferometer Network
(e-MERLIN). The e-MERLIN 1.6 GHz observations confirmed the diffuse, non-
collimated nature of the extended emission, as well as its steep spectral
index. Interestingly, the source is partially resolved in VLBA observations at
5 GHz, but the pc-scale jet seems to be pointing north-east. Fitting a
precession model to the JVLA 5.2 GHz data they confirm that it is a plausible
scenario for the origin of the detected emission. Another possibility they
bring up is intermittent activity, as J1302+1624 seems to be undergoing a
merger, and mergers are known to cause this kind of behaviour. Furthermore, a
merger can also be responsible for the change in the direction of the jet.
Current data are not enough to distinguish between these scenarios.
Interestingly, the flattened emission region we see in the $\alpha$ map
coincides with the southern tidal structure seen in Fig. 6 of Congiu et al.
(2020), possibly indicating that the jet is interacting with the interstellar
medium in the tidal tail. J1302+1624 is a blue outlier as its [O III] lines
are shifted by -191 km s-1, but it is unclear whether the shift is due to the
jet or possibly the merger.

The jet is clearly the main source of radio emission in this source, but as
indicated by the BPT diagram, there probably is also some star formation going
on in the host. The star formation diagnostics are dominated by the jet
emission, and only the W3-W4 colour of 2.62 reveals possible star formation.
$q22$ is -0.11, the W3 flux density is 63.0 mJy lower than the
$S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{int}}$ is 10.8 mJy higher than
$S_{\textrm{5.2GHz, CDFS}}$.

####  5.2.31 J1305+5116

J1305+5116 (SDSS J130522.74+511640.2) is a radio-loud, bright NLS1 at $z$ =
0.788 with a peak flux density of 32.10 mJy beam-1, an integrated flux density
of 53.41 mJy and the flux densities of the tapered maps approximately the
same. It is one of the most luminous sources in our sample with log $\nu
L_{\nu,\textrm{int}}$ = 42.91 erg s-1. J1305+5116 has also been detected in
gamma-rays, which - when taking into account the quite high redshift of this
source - strongly suggests it hosts a relativistic, beamed jet (Liao et al.
2017). Its radio map (Fig. 38(a)) shows two components with a separation of
8.1 kpc, and at a position angle of 168°. The northern component is
responsible for about 62% of the emission (33.1 mJy). The $\alpha$ map reveals
that this is the core, since it shows a flat spectral index ($\alpha$ =
-0.07), whereas the spectral index of the southern component is -0.73 (RA
13:05:22.76, Dec 51:16:39.10), and consistent with optically thin synchrotron
emission. The northern component is centred at the centre of the featureless
host galaxy (Fig. 38(e)). The SDSS $r$ band 25.0 mag arcsec-2 isophotal
semimajor axis of the host is 15.4 kpc, so also the jet is still confined
within the host galaxy.

Based on archival data J1305+5116 was previously classified as a steep
spectrum source: its flux densities at 150 MHz (6C/7C), 325 MHz (WENNS), and
1.4 GHz (FIRST) are 320, 211, and 86.9 mJy, respectively (Yuan et al. 2008).
It was later observed with the Allen Telescope Array (ATA) at 1.4 GHz with a
flux density of 92.7 mJy (Croft et al. 2010), indicating some degree of
variability. Higher-resolution observations revealed a north-south -oriented
core-jet structure (Petrov 2013). In VLBA observations at 5 GHz the core flux
density was found to be 15 mJy, and the jet flux density 8.9 mJy. These values
are less than 50% of what we observe at 5.2 GHz indicating that a substantial
amount of the emission is resolved-out at the VLBA scales. The source was not
detected at 22 GHz with the JVN with a detection limit of 9 mJy (Doi et al.
2016).

Komossa et al. (2018) investigated the optical spectrum of J1305+5116 and
found extreme emission line shifts, with the largest velocities derived from
them exceeding the escape velocity of the galaxy. They also found ionisation
stratification – the shifts of higher ionisation potential lines are larger –
and argue that this implies that the shifts are not caused by local
interaction, but concern the whole narrow-line region. They conclude that the
properties of the optical spectrum are most probably indicative of an early
stage of jet evolution, where a young jet is still advancing through the dense
interstellar medium and interacting with it. This scenario is in agreement
with the radio morphology of the source, since the jet still seems to reside
within the host. This source would be ideal for future observations about the
jet - interstellar medium interaction.

The radio emission in J1305+5116 seems to be produced by a powerful nucleus
and a relativistic, beamed jet. Nevertheless, for consistency the star
formation diagnostics of this source are: the W3-W4 colour is 2.34, $q22$ is
-0.61, $S_{\textrm{1.4GHz, JVLA}}$ is 72.3 mJy higher than $S_{\textrm{W3}}$,
$S_{\textrm{int}}$ is 5.03 mJy higher than $S_{\textrm{5.2GHz, CDFS}}$.

####  5.2.32 J1317+6010

J1317+6010 (SBS 1315+604) is a radio-quiet NLS1 ($z$ =0.137) with a peak flux
density of 0.51 mJy beam-1, an integrated flux density of 0.64 mJy, and the
tapered flux densities slightly higher, 0.80 and 0.79 mJy for the 90 and
60k$\lambda$ taper, respectively. Due to its redshift the low flux density
translates to a moderate luminosity of log $\nu L_{\nu,\textrm{int}}$ = 39.42
erg s-1. The only archival radio detection is 1.82 mJy from FIRST, yielding a
1.4-5.2 GHz spectral index of -0.80. Interestingly, according to the $\alpha$
map in Fig. 39(a) the core spectral index is quite flat with $\alpha$ = -0.34.
This is consistent with what was found in B18, but since they used the same
data set, this result should be checked. It can be an artefact due to the
faintness of the source. The morphology seems to be somewhat extended, and the
central region seems to be surrounded by patchy, diffuse emission. In the
tapered maps (Figs. 39(c) and 39(d)) these blobs form one, quite featureless,
emission region. The radio emission is confined within the host galaxy, that
does not show any distinguishable morphology either (Fig. 39(e)).

The star formation proxies indicate star formation in the host, W3-W4 is 2.6,
$q22$ is 1.60, $S_{\textrm{W3}}$ is 12.8 mJy higher than $S_{\textrm{1.4GHz,
JVLA}}$, and $S_{\textrm{int}}$ and $S_{\textrm{5.2GHz, CDFS}}$ are about the
same. However, the BPT diagrams using different emission line ratios all
indicate that the gas has been ionised by the AGN (Stern & Laor 2013). The [O
III] lines are not shifted, but the source does show an [O III] wing with a
velocity of -377 km s-1 and a FWHM of 1436 km s-1, so the dust heated by the
AGN could explain the mid-infrared emission.

####  5.2.33 J1337+2423

J1337+2423 (IRAS 13349+2438) is a radio-quiet source at $z$ = 0.108,
originally classified as an NLS1, but later found to be a broad-line Seyfert 1
(BLS1) with the FWHM(H$\beta$) $\sim$ 2500 km s-1 (Lee et al. 2013). However,
its optical spectrum exhibits very strong Fe II emission, unlike BLS1s in
general, so it might be an “overgrown” NLS1. We decided to include it in the
sample because it still shares a lot of properties with NLS1s, and might be
interesting from a evolutionary point of view. It is defined as radio-quiet,
mostly due to its strong optical/infrared emission. Its peak flux density is
9.64 mJy beam-1, the integrated flux density is 9.96 mJy, and the tapered flux
densities are about the same. It has a high luminosity of log $\nu
L_{\nu,\textrm{int}}$ = 40.26 erg s-1. Its FIRST flux density is 19.9 mJy,
giving a 1.4-5.2 GHz spectral index of -0.53. The $\alpha$ map in Fig. 40(a)
shows generally steep spectral indices: the total spectral index is -1.03, and
the core spectral index is -1.01. The source is slightly elongated toward
east. The tapered maps do not reveal significant new emission. The radio
emission is confined within the host galaxy, which shows hints of a bar and
possible spiral arms or disturbed morphology (Fig.40(e)).

The star formation diagnostics, except for the W3-W4 colour, indicate strong
star formation. The W3-W4 colour is 1.95, $q22$ is 1.29, the W3 flux density
is 424.1 mJy higher than the $S_{\textrm{1.4GHz, JVLA}}$, and
$S_{\textrm{5.2GHz, CDFS}}$ is 2.8 mJy higher than $S_{\textrm{int}}$.
However, this source has been studied in infrared in detail, and it has been
classified as a hot dust-obscured galaxy, where almost all the mid-infrared
emission originates from the AGN-heated dust (Lyu & Rieke 2018). The AGN-
fraction of the mid-infrared emission was studied also in (Alonso-Herrero et
al. 2016) and Gruppioni et al. (2016) and was found to be 0.97 and 0.98,
respectively, so it can be concluded that the very bright mid-infrared
emission is of AGN origin. On the other hand, this source also shows
considerable star formation, with the estimates ranging between $\sim$15 to 45
$M_{\odot}$ yr-1 (Sargsyan et al. 2012; Mahajan et al. 2019), so the
contribution of the star formation to the radio emission cannot be totally
ruled out.

J1337+2423 has also been found to host strong multiphase ultrafast outflows
with velocities of $\sim$0.14 c and $\sim$0.27 c (Parker et al. 2020), as well
as two warm absorbers with velocities around -600 km s-1 (Parker et al. 2018).
The optical spectrum does not show [O III] line shifts, but reveals a strong
[O III] wing component with a velocity of -598 km s-1 and a FWHM of 2241 km
s-1. Taking into account the proof of strong outflows in this source, it is
possible that also the eastward extended emission is rather caused by an
outflow than a jet. Also the steep spectral indices support this hypothesis.

####  5.2.34 J1348+2622

J1348+2622 (SDSS J134834.28+262205.9) is a radio-quiet NLS1 at $z$ = 0.917
with a very low integrated flux density of 0.38 mJy. However, this source has
the highest redshift among our sample, so the low $S_{\textrm{int}}$
translates to quite a high luminosity of log $\nu L_{\nu,\textrm{int}}$ =
40.84 erg s-1. The source does not exhibit any distinct features in the radio
maps (Fig. 41), and unsurprisingly also its host galaxy is unresolved due to
the redshift (Fig. 41(e)). The $\alpha$ map (Fig. 41(a)) shows that the
spectral index of the source is very steep, the total spectral index is -1.06,
and the core spectral index is -1.30.

The star formation diagnostics do not imply strong star formation in this
source, W3-W4 is 1.83, $q22$ is 0.32, and both $S_{\textrm{W3}}$ and
$S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{int}}$ and $S_{\textrm{5.2GHz,
CDFS}}$ are roughly the same. However, due to the high redshift of this
source, the PAH features around 11.3 $\mu$m are not within the W3 band
anymore, but actually fall into the W4 band, so the star formation proxies
used are not necessarily reliable for this source. Since the PAH features
should be within the W4 band we decided to use it instead of W3 in the
$S_{\textrm{W3}}$ \- $S_{\textrm{1.4GHz, JVLA}}$ relation. According to this
the W4 flux density is $\sim$1.6 mJy higher than $S_{\textrm{1.4GHz, JVLA}}$,
still not a very robust result in favour of star formation.

Based on these data it is not possible to determine the predominant source of
the radio emission in J1348+2622. However, star forming, or even starburst
galaxies rarely exceed the radio luminosity of log $\nu L_{\nu}$ = 40 erg s-1
(e.g., Sargsyan & Weedman 2009), which might suggest that the radio emission
is dominated by the AGN, and that J1348+2622 might be a CSS-like NLS1 galaxy.

####  5.2.35 J1358+2658

J1358+2658 (SDSS J135845.38+265808.4) is a radio-quiet source at $z$ = 0.331
with a peak flux density of 0.52 mJy beam-1, an integrated flux density of
0.61 mJy, with the flux densities of the tapered maps similar. Due to the
redshift of this source the low flux density results in a high luminosity of
log $\nu L_{\nu,\textrm{int}}$ = 40.05 erg s-1. The radio map in Fig. 42(a)
shows a core, and faint extended emission north-eastward. A similar structure
is seen also in the tapered maps (Figs. 42(c) and 42(d)). However, with these
data we cannot be certain the emission is real. The $\alpha$ map probably
suffers from the faintness of the source, and shows variable, possibly not
reliable, values. The general trend seems to be on the verge between flat and
steep, but based on these data nothing conclusive can be said. The FIRST flux
density is 1.17 mJy, and thus the 1.4-5.2 GHz spectral index -0.50, somewhat
in agreement with the $\alpha$ map, that has a total spectral index of -0.42.
The bulk of the radio emission comes from within the host galaxy, of whose
morphology nothing can be said, but the north-eastward extension seems to be
reaching outside it. The SDSS $r$ band 25 mag arcsec-2 isophotal semimajor
axis of the galaxy is 7.7 kpc, whereas the maximum extent of the radio
emission is 17 kpc.

Caccianiga et al. (2015) estimate the star formation activity in the host to
be at a starburst level with 129 $M_{\odot}$ yr-1. Strong star formation is
supported by the W3-W4 colour of 2.69, $q22$ of 1.30, the W3 flux density that
is 5.4 mJy higher than $S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{int}}$ and
$S_{\textrm{5.2GHz, CDFS}}$ that are about equal. Its optical spectrum also
shows an [O III] wing with a velocity of -253 km s-1 and a FWHM of 1366 km
s-1, so the contribution of the AGN-heated dust cannot be ruled out either.

The extended emission - if real - cannot easily be explained by star
formation, and it is most likely of nuclear origin, but its nature cannot be
determined with these data. Furthermore, the broad-band spectral index close
to the flat/step border is not in agreement with a pure star formation
scenario, which should result in a steeper spectral index. This indicates that
whereas a fraction on the radio emission originates from the starburst, there
needs to be also a flatter component present to account for the observed
spectral index. Moreover, the luminosity of this source would be very high for
a pure starburst galaxy, and thus the radio emission within the galaxy seems
to be a combination of starburst and nuclear activity.

####  5.2.36 J1402+2159

J1402+2159 (RX J1402.5+2159) is a radio-quiet NLS1 at $z$ = 0.066 with very
fait radio emission, the peak flux density is 0.28 mJy beam-1, the integrated
flux density is 0.33 mJy, and the 90k$\lambda$ tapered flux density is 0.46
mJy. J1402+2159 is a quite nearby source, so also the resulting luminosity is
low with log $\nu L_{\nu,\textrm{int}}$ = 38.34 erg s-1. The source was
detected in FIRST with a flux density of 0.87 mJy, giving a 1.4-5.2 GHz
spectral index of -0.74. In addition to these there are not other radio
observations of this source. The spectral index seen in the $\alpha$ map in
Fig. 43(a) is even steeper, the total spectral index is -0.95, and the core
spectral index is -1.54. However, the spectral index values might be affected
by the faintness and the low S/N of the data. The source morphology shows an
extension toward north-east. This feature is more pronounced in the tapered
maps (Figs. 43(c) and 43(d)). Overlaid with the host galaxy (Fig. 43(e)), it
looks like the radio emission reaches slightly outside of it, or at least to
the edges, even when considering that the elongated beam somewhat exaggerates
the extent of the emission.

The host galaxy is probably of E/S0 type (Ohta et al. 2007), and all the star
formation diagnostics indicate activity in this source, the W3-W4 colour is
2.75, $q22$ is 2.16, $S_{\textrm{W3}}$ is 41.2 mJy higher than
$S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{5.2GHz, CDFS}}$ is 2.1 mJy higher
than $S_{\textrm{int}}$. These diagnostics are supported by BPT diagrams based
on various line ratios where this source is consistently classified as a star-
forming galaxy, indicating that the main source of ionising emission in this
galaxy are stars. The data are consistent with the star formation activity
being the main source of radio emission in this source. However, the nature of
the extended radio emission cannot be determined with these data, more,
preferably higher S/N observations would be needed for that.

####  5.2.37 J1536+5433

J1536+5433 (Mrk 486) is a radio-quiet NLS1 at $z$ = 0.039 with very faint
radio emission: the peak flux density is 0.44 mJy beam-1 and the integrated
flux density is = 0.50 mJy. It also has a low luminosity of log $\nu
L_{\nu,\textrm{int}}$ = 37.98 erg s-1. It was detected at 5 GHz in the 1980’s
with a flux density of 0.47 mJy (Kellermann et al. 1989), indicating that it
does not show significant variability. The FIRST detection at 1.4 GHz is 1.23
mJy, giving a 1.4-5.2 GHz spectral index of -0.69. Our $\alpha$ map in Fig.
44(a) shows slightly flatter spectral indices, the total spectral index is
-0.35, and the core spectral index is -0.50. Due to the faintness of this
source we cannot be sure how reliable the spectral index range is but it seems
like the general spectral index is rather flat, which is in agreement with the
spectral index estimated in B18. The radio emission in confined within the
inner regions of the host galaxy (Fig. 44(e)), which appears to be late-type.

The PAH features of J1536+5433 at 11.3 $\mu$m were studied by Martínez-Paredes
et al. (2019), and whereas they were able to detect them they found no
evidence of enhanced star formation; in fact the estimated star formation rate
was only 0.24 $\pm$ 0.08 $M_{\odot}$ yr-1. Although it should be kept in mind
that an AGN can suppress the PAH features. Nevertheless, the mid-infrared
colour agrees with this (W3-W4 = 1.92), but the other diagnostics do not. The
$q22$ is 2.12, the W3 flux density is substantially higher, about 52.1 mJy,
than $S_{\textrm{1.4GHz, JVLA}}$, and also $S_{\textrm{5.2GHz, CDFS}}$ is
$\sim$1.0 mJy higher than $S_{\textrm{int}}$. Whereas $S_{\textrm{int}}$ and
$S_{\textrm{5.2GHz, CDFS}}$ are comparable and indicate that the radio
emission could be explained by modest star formation, the W3 flux density is
surprisingly high, and hard to explain by the low star formation seen in this
source. The [O III] lines of J1536+5433 are not shifted, and do not show a
wing either. This, of course, does not rule out the possibility that the
excess mid-infrared emission might originate from the equatorial dust heated
by the AGN. This would be in line with the flat spectral index – uncommon for
star formation emission – of the source.

####  5.2.38 J1555+1911

J1555+1911 (Mrk 291) is a faint, radio-quiet NLS1 ($z$ = 0.035) with a peak
flux density of 0.12 mJy beam-1 and an integrated flux density of 0.36 mJy.
The flux densities of the tapered maps are 0.42 and 0.44 mJy, for the 90 and
the 60k$\lambda$ taper, respectively. The normal map (Fig. 45(a)) exhibits
diffuse emission, and it is hard to say whether a proper core is even present.
The FIRST detection of this source is 1.85 mJy, giving a 1.4-5.2 GHz spectral
index of -1.25. Also our $\alpha$ map shows a very steep spectral index: the
core spectral index is -1.13, and the total spectral index is -0.82, albeit
the south-western part of the emission exhibits a flat spectral index, which
might be due to the edge effects and thus not reliable. The tapered maps (Fig.
45(c) and 45(d)) show a morphology slightly elongated toward north-west.
J1555+1911 is hosted by a spiral galaxy (Fig. 45(e)), and most of the radio
emission is confined within the bulge of the galaxy.

Based on its mid-infrared properties, the diffuse radio morphology, and the
low luminosity (log $\nu L_{\nu,\textrm{int}}$ = 37.85 erg s-1) it seems
rather evident that the predominant source of the radio emission is star
formation. W3-W4 colour is 2.62, $q22$ is 1.67, W3 flux density is $\sim$19.2
mJy higher than $S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{int}}$ and
$S_{\textrm{5.2GHz, CDFS}}$ are roughly in agreement.

####  5.2.39 J1559+3501

J1559+3501 (Mrk 493) is a radio-quiet NLS1 at $z$ = 0.031 with faint radio
emission, the peak flux density is 0.46 mJy beam-1, the integrated flux
density is 1.41 mJy, and the flux densities of the tapered maps are 1.47 and
1.48 mJy for the 90 and 60k$\lambda$, respectively. It is a nearby source so
the flux density results in a low luminosity of log $\nu L_{\nu,\textrm{int}}$
= 38.19 erg s-1. The radio map in Fig. 46(a) shows a core surrounded by
diffuse emission, slightly more elongated toward north-west and south-east.
The tapered maps in Figs. 46(c) and 46(d) do not show any additional
structures. The $\alpha$ map suffers from low S/N, as can be seen especially
close to the edges, but in general the map shows very steep spectral indices,
and the total spectral index is as steep as -1.16. The FIRST detection of this
source is 3.39 mJy, giving a 1.4-5.2 GHz spectral index of -0.67. Yang et al.
(2020) report a JVLA A-configuration 8.84 GHz flux density of 0.56 mJy,
resulting in a very steep 5.2-8.84 GHz spectral index of -1.74. They claim
that the 5.2-8.84 GHz spectral index is flat, but their reported value for the
integrated 5.2 GHz flux density is only 0.62 mJy, whereas their and ours peak
flux densities are about the same. They obtained the value using the same data
set we are using, but their rms is higher (11 vs. 15.64 $\mu$Jy beam-1).
However, the 1.4-5.2 GHz spectral index is steep, and it does not seem
probable that it would flatten above 5.2 GHz.

It is worth noting that the detected radio emission is confined within the
bulge of the host (Fig. 46(e)) and probably does not represent the whole radio
emission of the source. The host galaxy is a barred spiral that also exhibits
a nuclear ring and a nuclear dust spiral (Crenshaw et al. 2003; Deo et al.
2006; Ohta et al. 2007). Popović et al. (2009) studied the nuclear region in
detail and concluded that the nuclear ring can indeed be a place of enhanced
star formation. They estimated an unusually high star formation rate of 2
$M_{\odot}$ yr-1 for the innermost $<$ 1 kpc. In the BPT diagrams J1559+3501
is classified as a star-forming galaxy, indicating that the narrow-line region
lines as excited by star formation rather than the AGN. Sani et al. (2010)
detected strong 6.2 $\mu$m PAH features in this source, confirming the
presence of star formation. Our star formation diagnostics are somewhat in
agreement with this, but it should be noted that the WISE flux densities cover
the whole galaxy, whereas we likely are missing some faint, extended radio
emission from host. The W3-W4 colour is 2.34, $q22$ is 1.38, W3 flux density
is $\sim$64.3 mJy higher than $S_{\textrm{1.4GHz, JVLA}}$, and
$S_{\textrm{5.2GHz, CDFS}}$ is 1.3 mJy higher than $S_{\textrm{int}}$. The [O
III] lines of J1559+3501 show a prominent wing with a velocity of -349 km s-1
and a FWHM of 807 km s-1, so also nuclear winds are present.

The evidence seems to indicate that the star formation is a significant, if
not the major, contributor to the radio emission. The nuclear, possibly star-
forming ring found in previous studies falls within the radio emitting region
in our map, and also the radio morphology seems to be more consistent with
star formation than nuclear activity. This does not rule out the possibility
that the AGN contributes too, but higher resolution observations will be
needed to disentangle these two mechanisms.

####  5.2.40 J1633+4718

J1633+4718 (SDSS J163323.58+471858.9) is a radio-loud NLS1 ($z$ = 0.116) whose
radio map in Fig. 47(a) shows two distinct components. The integrated flux
density of the core component is 24.48 mJy, and the integrated flux density of
the north component is 0.79 mJy. The flux densities of the tapered maps are
very similar to these. J1633+4718 is a high-luminosity source with log $\nu
L_{\nu,\textrm{int}}$ = 40.67 erg s-1. J1633+4718 resides in an interacting
system (Olguín-Iglesias et al. 2020): the core component is associated with
the disk-like host galaxy of J1633+4718, whereas the north component is
associated with another disk galaxy (Fig. 47(e)). The nuclei of the two
galaxies are $\sim$8 kpc apart. The core spectral index of the component that
hosts the NLS1 is -0.54, and thus on the verge between flat and steep. The
$\alpha$ map of the companion galaxy suffers from low S/N and edge effects,
but its “core” spectral index is -0.69.

Past radio observations of J1633+4718 indicate variability in the core.
Neumann et al. (1994) found an inverted radio spectrum above 5 GHz with
simultaneous observations, whereas Doi et al. (2011) found the 1.7 to 8.5 GHz
spectral index to be steep with quasi-simultaneous observations. Also archival
data at $\sim$5 GHz suggest considerable variability: Gregory & Condon (1991)
reported a flux density of 34 mJy at 4.85 GHz obtained with the GBT, Laurent-
Muehleisen et al. (1997) obtained 30 mJy with the VLA at 5 GHz, and Gu & Chen
(2010) reported 55.2 mJy with the VLBA at 5 GHz. This is indicative of jet
activity, which has been confirmed by high brightness temperature of at least
1011.3 K (Doi et al. 2007), and a pc-scale core-jet radio morphology in high
resolution observations at 1.7 and 5 GHz (Gu & Chen 2010; Doi et al. 2011). At
8.4 GHz with the JVN the source remains unresolved with a flux density of 21.2
mJy.

Archival FIRST 1.4 GHz flux density for J1633+4718 is 65.02 mJy, whereas Doi
et al. (2011) were able to detect 55.5 mJy at 1.7 GHz with the VLBA. The
discrepancy might be due to resolved-out emission in the VLBA observations,
but also because of the variability. Several previous studies suggest that
this might be a CSS-like source due to its small size and steep spectral
incides, but taking into account that the source also shows flat or even
inverted radio spectra, it seems more probable that the occasional steep
indices have been due to the temporal variability in the source. Moreover, the
22 GHz flux density was found to be as high as 163 mJy observed with the JVN
(Doi et al. 2016), also not consistent with the CSS scenario.

The W3-W4 colour of 2.67 indicates that there might also be star formation
present in the galaxy, and Caccianiga et al. (2015) estimated the star
formation rate to be as high as 68 $M_{\odot}$ yr-1. The other indicators do
not favour star formation, since the radio emission is clearly dominated by
the nucleus: $q22$ is 0.22, the W3 flux density is $\sim$24.4 mJy lower than
$S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{5.2GHz, CDFS}}$ is 24.1 mJy lower
than $S_{\textrm{int}}$. However, also a strong nuclear wind seen in the [O
III] wing is present with a velocity of -279 km s-1 and a FWHM of 1186 km s-1
so AGN-heated dust can contribute to the observed mid-infrared emission. The
main source of the radio emission in J1633+4718 is the AGN, but the weak radio
emission of the companion galaxy might be produced by star formation.

J1633+4718 is also an interesting X-ray source, since an ultrasoft X-ray
component with a low blackbody temperature of $\sim$30 eV was detected in it
(Yuan et al. 2010; Mallick et al. 2016). This was intepreted as a direct
detection of the accretion disk, and so far similar components have not been
detected in other NLS1s. Mallick et al. (2016) also detect excess UV emission,
that requires the presence of a jet to be explained.

####  5.2.41 J1703+4540

J1703+4540 (SDSS J170330.38+454047.1) is a radio-loud NLS1 at $z$ = 0.060 that
was previously classified as a CSS with a turnover frequency below 150 MHz
(Snellen et al. 2004). Our $\alpha$ map, shown in Fig. 48(a), agrees with this
classification since the spectral index we find is steep: the total spectral
index is -0.87, and the core spectral index is -0.91. The peak flux density is
32.26 mJy beam-1, the integrated flux density of the normal map (Fig. 48(a))
is 34.68 mJy, with log $\nu L_{\nu,\textrm{int}}$ = 40.20 erg s-1, and it
exhibits very mild elongation toward south-east. The integrated flux densities
of the tapered maps (Figs. 48(c) and 48(d)) are about a mJy higher, and both
maps reveal extended, but faint radio emission toward south-east. The host of
J1703+4540 is clearly a spiral galaxy, as seen in Fig. 48(e), and the radio
emission is concentrated within the bulge of the galaxy, which is of a pseudo-
bulge type (Olguín-Iglesias et al. 2020). The extended radio emission seems to
somewhat trace the southern bar- or arm-like structure, possibly indicating
the presence of star formation in the galaxy. The star formation diagnostics
partly support this, as the mid-infrared colour W3-W4 is 2.59, and the W3 flux
density 15.9 mJy higher than the $S_{\textrm{1.4GHz, JVLA}}$.
$S_{\textrm{int}}$, on the other hand, is 28.8 mJy higher than
$S_{\textrm{5.2GHz, CDFS}}$, and $q22$ is only 0.52, suggesting that star
formation is not sufficient to explain all of the radio emission.

Indeed, analysing archival VLBA observations of J1703+4540 at 5 GHz Gu & Chen
(2010) found a one-sided core-jet structure directed toward south-west. They
also reported considerable flux density variability, which might also affect
our star formation diagnostics, since we do not know in which state the source
was during our observations. However, our integrated flux density is $\sim$20
mJy lower than the 5 GHz VLBA flux density (56.8 mJy), suggesting that the
source was in a rather low state during our observations. The presence of a
jet was confirmed in VLBA observations at 1.7 GHz by Doi et al. (2011), who
found a similar one-sided structure, consistent with the jet of Gu & Chen
(2010), and with a projected length of 35 pc. The jet emission is quite
diffuse and does not appear Doppler boosted, and they conclude that the jet is
either sub-relativistic or seen at a large angle. They also detect hints of a
counter-jet but cannot confirm whether it is real or not. Giroletti et al.
(2017) claim to have confirmed the presence of emission on the counter-side of
the core at 5 GHz, and they classify J1703+4540 as an asymmetric double, and a
low-power compact (LPC, Kunert-Bajraszewska et al. 2010) source, since it does
not show considerable radio morphology evolution over time, as would be
expected from a true CSS.

A powerful molecular outflow with a bulk velocity of -660 km s-1 (Longinotti
et al. 2018), and a multi-component ultra-fast X-ray outflow showing
velocities up to $\sim$0.1 c (Sanfrutos et al. 2018), consistent with each
other, have also been found in this source. Longinotti et al. (2018) state
that the outflow seen in J1703+4540 is consistent with an energy-conserving
outflow propagating through the host galaxy, an instance of probably the most
important and effective way of AGN feedback. In agreement with the outflows,
also an [O III] wing component is present, with a velocity of -336 km s-1, and
a FWHM of 646 km s-1. Interestingly, the [O III] emission line shows a
redshift of 210 km s-1(Berton & Järvelä 2021), the origin of which remains
unclear. Since the wing velocity is defined relative to the [O III] core, its
real velocity, compared to the restframe, is smaller than the previously
reported value. Spatially resolved spectra would be needed to study the
morphology and kinematics of [O III] in more detail.

####  5.2.42 J1713+3523

J1713+3523 (FBQS J1713+3523, $z$ = 0.083) is an NLS1 that would be classified
as radio-quiet based on the FIRST detection ($S_{\textrm{int, 1.4GHz}}$ =
11.24 mJy) or our data ($S_{\textrm{int}}$ = 3.81 mJy), but it was detected at
22 GHz with the JVN with a whopping flux density of 138 mJy (Doi et al. 2016),
implying extreme variability and jet activity. Interestingly the 1.4-5.2 GHz
spectral index is steep with $\alpha$ = -0.66, in agreement with our $\alpha$
map in Fig. 49(a) whose total spectral index is -0.97. Assuming the less
extreme spectral index of -0.66, the extrapolated 22 GHz flux density would be
1.47 mJy, requiring almost a hundred-fold increase in the flux density to
achieve the observed 138 mJy. The spectral index between 5.2 and 22 GHz is
2.49, very close to the theoretical synchrotron self-absorption (SSA) limit.
It should be noted that the observations are not simultaneous, the 22 GHz
observations were performed in April 2014, and the JVLA observations in
September 2015. However, radio flares usually propagate from higher to lower
frequencies within time-scales from several months to a few years, so if
anything the 22 GHz flare should have increased the 5.2 GHz flux density at
the time of our observation, if it would have propagated to such low
frequencies, so the estimated spectral index should be indicative of the real
spectral index.

The 22 GHz observation was done with the JVN so it probes mas scale emission
in the source. J1713+3523 lies at $z$ = 0.083, giving a scale of 1.562 pc
mas-1, thus the observed flux density must have come from a region of some
tens of pcs across. It is interesting that we do not see any signs of the
flattening of the spectral index in our $\alpha$ map, which would have been
expected to happen if the 22 GHz flare would have propagated to lower
frequencies. An alternative to a conventional AGN radio flare is that the 22
GHz flare was a sign of recently started jet activity in J1713+3523, as seen
in high-frequency peakers (HFP). In this case there could exist a quasi-
stationary feature at higher radio frequencies, with the inverted radio
spectrum produced by SSA (O’Dea & Saikia 2021) or free-free absorption by the
ionised gas produced in the bow shock of the jet (Bicknell et al. 1997). The
[O III] lines of J1713+3523 are extremely blueshifted with $v$ = -674 km s-1,
providing further support of the influence of a jet still confined within its
host galaxy and in interaction with the interstellar medium. Interestingly, no
[O III] wing is present in this source.

The whole radio emission seen in the JVLA observation of J1713+3523 could be
explained by star formation, as the W3 flux density is $\sim$17.8 mJy higher
than $S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{int}}$ and
$S_{\textrm{5.2GHz, CDFS}}$ are roughly the same. On the other hand, the near-
infrared colour, W3-W4 = 2.11, nor the $q22$ of 0.62, do not necessarily
indicate strong star formation. However, the source is not extraordinarily
bright, with log $\nu L_{\nu,\textrm{int}}$ = 39.54 erg s-1, lower than for
jetted sources in B18 in general. The higher frequency observations confirm
the presence of a jet, but it does not seem to dominate the 5.2 GHz radio
emission.

####  5.2.43 J2242+2943

J2242+2043 (Ark 564, $z$ = 0.025) is a radio-quiet NLS1 with a peak flux
density of 5.61 mJy beam-1, and an integrated flux density of 10.33 mJy that
gives quite a low luminosity of log $\nu L_{\nu,\textrm{int}}$ = 38.88 erg
s-1. It was detected in NVSS with a flux density of 29.1 mJy, yielding a
spectral index of -0.79 between 1.4 and 5.2 GHz, and at 95 GHz with a flux
density of 1.14 mJy (Behar et al. 2015), giving a spectral index of -0.76
between 5.2 and 95 GHz. These results are consistent with our $\alpha$ map
(Fig. 50(a)), which has a total spectral index of -0.81, and a core spectral
index of -0.85. The radio morphology of the source is clearly elongated toward
north, as also found in previous studies (e.g. Moran 2000), and the spectral
index seems to slightly flatten north-ward. The tapered maps (Figs. 50(c) and
50(d)) do not show any additional structures, and the radio emission is
confined within the bulge of a galaxy with a barred spiral morphology (Fig.
50(e)). This source exhibits complex X-ray properties as both relativistic and
sub-relativistic outflows have been detected (e.g., Gupta et al. 2013b, a;
Khanna et al. 2016).

Whereas the asymmetric radio morphology indicates that J2242+2943 possesses a
jet or a wind that produces radio emission, the star formation proxies suggest
that also star formation is present. The W3-W4 colour is 2.57, $q22$ is 1.16,
and the W3 flux density is as much as $\sim$123.8 mJy higher than
$S_{\textrm{1.4GHz, JVLA}}$, and $S_{\textrm{int}}$ and $S_{\textrm{5.2GHz,
CDFS}}$ are comparable. The low radio luminosity of the source also suggests
that a possible jet cannot be very powerful, supported by the fact that no
shift is seen in the [O III] lines. However, an [O III] wing is detected, with
the velocity and the FWHM of -564 and 1090 km s-1, respectively, so the mid-
infrared emission might be enchanced by AGN-heated dust emission. It seems
plausible that the radio emission is a combination of AGN and star formation
related processes, but nothing definite can be said before more detailed
observations.

####  5.2.44 J2314+2243

J2314+2243 (RX J2314.9+2243, $z$ = 0.169) is a radio-quiet source with a peak
flux density of 6.17 mJy beam-1, and an integrated flux density of 7.02 mJy.
It was also detected in NVSS with a flux density of 18.7 mJy, giving a
spectral index of -0.75 between 1.4 and 5.2 GHz. This value is in agreement
with our $\alpha$ map in Fig. 51(a): the total spectral index is -0.81, and
the core spectral index is -0.85. The radio maps (Fig. 51) do not show any
distinctive features, nor does the host galaxy in Fig. 51(e), within which the
radio emission is confined.

The W3-W4 colour (2.33) and $q22$ (0.85) do not indicate strong star
formation, and also $S_{\textrm{int}}$ is $\sim$4.5 mJy higher than
$S_{\textrm{5.2GHz, CDFS}}$. However, $S_{\textrm{W3}}$ is 43.8 mJy higher
than $S_{\textrm{1.4GHz, JVLA}}$. The origin of this excess emission can be
either star formation in the host galaxy, or AGN-heated dust. The existence of
a strong [O III] wing with a velocity of -1031 km s-1 and a FWHM of 1579 km
s-1 is in favour of the AGN-heated, possibly polar, dust being a major
contributor to the mid-infrared emission. J2314+2243 would be unusually bright
for a purely star forming galaxy, with log $\nu L_{\nu,\textrm{int}}$ = 40.48
erg s-1, so it seems possible that the AGN is the main source of the radio
emission. J2314+2243 could be another candidate for a CSS-like source.

##  6 Extraordinary diversity of NLS1s

The most notable result is the remarkable diversity of NLS1s: they show a host
of different components contributing to the radio, as well as the mid-infrared
emission, and no two sources are alike. Spatially resolved spectral index map
utilised in this paper offer a clear advantage in disentangling these
components when compared to conventional ways of determining the spectral
indices, that usually give one value for the whole emitting region. The
$\alpha$ maps allow us to examine the spectral index and its changes over the
whole emitting region, and determine the most plausible origin of the emission
in different parts of the source. In addition, the associated errors of the
$\alpha$ maps derived using the mt-mfs method are considerably smaller than
the errors for conventional estimates. Using this method, complemented by data
at other wavelengths, some NLS1 groups with similar properties can be
identified, and are discussed below.

###  6.1 AGN-dominated sources

In about one third of the sources in our sample the radio emission is clearly
dominated by the AGN. These sources are characterised by moderate to high
radio luminosities (all but one have log $\nu L_{\nu,\textrm{int}}>$ 39.00 erg
s-1), asymmetric radio morphologies, and/or variable radio emission, and flat
spectral indices in some cases. We can see considerable diversity even within
this group, probably due to different ages, inclination angles, and
environments of the sources.

####  6.1.1 Large-scale jets

A few of our sources exhibit radio jets with extents of the order of tens of
kpcs, with the jets reaching well outside the host galaxy. In some sources we
are also able to detect counter-jet emission, indicating that these sources
are seen at rather large inclination angles, and projection effects cannot be
responsible for the narrowness of their permitted lines.

Based on simple travel-time arguments these sources cannot be kinematically
very young, and must have maintained the jets for quite a long period of time
for them to propagate through the host interstellar medium. Czerny et al.
(2009) propose that intermittent activity due to radiation pressure
instabilities of the accretion disk could be common in young radio sources
accreting at high rates; the active, jetted phases would last for 103-104
years, separated by inactive, non-jetted periods of 104-106 years. It is
unclear whether such short activity periods can be responsible for the kpc-
scale emission seen in these impressive sources since the final extent of the
radio emission after the jets have turned off depends on several factors, for
example, the jet power, the length of the activity period, and the properties
of the medium it is expanding into.

Interestingly, one source in our sample, J1302+1624, shows possible relic
emission with a very steep spectral index. This seems to be a consequence of
its jets changing direction (Congiu et al. 2020), and with the current data we
cannot say whether the jets turned off at some point or not. On the other
hand, the most extended NLS1 discovered so far, J0814+5609, does not show
signs of intermittent activity, and based on its flat core and observations at
other radio frequencies the jets are currently active. However, this source
might represent a transition object between NLS1s and broad-line radio
galaxies, since its black hole mass is large for an NLS1 (108.44 M⊙, Berton et
al. 2016b), and its FWHM(H$\beta$) on the threshold between NLS1s and BLS1s
(2164 km s-1, Foschini et al. 2015). The other sources with large-scale jet
emission in our sample do not show signs of intermittency either, and their
black hole masses are well below 108 M⊙, implying that large-scale jets are
not exclusively a property of high black hole mass sources. Morphologically,
these sources resemble small broad-line radio galaxies, and are further proof
that eventually some NLS1s might evolve into fully developed radio galaxies
(Berton et al. 2016a).

####  6.1.2 Compact sources and small-scale jets

Another group whose radio emission is dominated by the AGN are the sources
with compact morphologies, showing no or moderate extended radio emission, or
small-scale jets. In these sources the high AGN activity - usually in form of
jets - was confirmed by, for example, flat spectrum, high luminosity,
variability, or previous high resolution radio observations. Interestingly,
many of these sources are nominally radio-quiet, casting yet another shadow
over the use of radio-loudness parameter as a proxy of AGN activity.

The compactness of these sources indicates that they are either less extended,
and thus possibly younger sources, or seen at smaller inclination angles. The
latter might be true especially in case of the most luminous sources, where
the emission is probably enhanced by considerable beaming. The lower
luminosity sources are probably either seen at large enough angles for the
beaming not to have considerable effects and/or host slower, sub-relativistic
jets.

An interesting individual is J1633+4718 since it resides in a late-type
galaxy, hosts most likely relativistic jets, and is interacting with another
late-type galaxy (Olguín-Iglesias et al. 2020). It has been hypothesised that
interaction might play a role in triggering the nuclear activity (Taniguchi
1999; Barth et al. 2008; Urrutia et al. 2008; Ellison et al. 2011), and seems
like jetted NLS1s might be a great example of this, since previous studies
have shown that interaction and mergers are more common among jetted than non-
jetted NLS1s (e.g., Olguín-Iglesias et al. 2017; Järvelä et al. 2018; Berton
et al. 2019). However, it should be remembered that in general the jetted
sources reside at higher redshifts than the non-jetted ones that have been
studied, and the impact of this difference has not been properly addressed
yet.

####  6.1.3 CSS-like sources

We identify five new NLS1s that exhibit properties consistent with the CSS
classification: J0806+7238, J1047+4725, J1138+3653, J1348+2622, and
J2314+2243. These sources show compact morphology, steep spectral index, and
luminosity that is too high to be explained by star formation or even a
starburst. Three out of five sources exceed the traditional luminosity
threshold for CSS sources at 5 GHz (log $\nu L_{\nu,\textrm{int}}$ = 40.60 erg
s-1), and the luminosities of the other two are log $\nu L_{\textrm{int}}>$
40.00 erg s-1.

A link between NLS1s and CSS sources has already been suggested several times,
and some NLS1s have been identified as CSS sources in the literature (e.g.,
Oshlack et al. 2001; Gallo et al. 2006; Caccianiga et al. 2014; Berton et al.
2016a). These results lead to a tentative model where some CSS sources,
especially the low-luminosity compact objects (LLCs) could be a part of the
parent population of jetted NLS1s (Berton et al. 2016a). According to Kunert-
Bajraszewska & Labiano (2010) some LLCs can be classified as high-excitation
sources, similarly to NLS1s, and their intrinsic properties, such as black
hole mass and Eddington ratio distributions, show significant overlap with
beamed NLS1s. Also their radio luminosity functions are consistent with each
other, when beaming is taken into account, and they show similar, simple radio
morphologies, without considerable diffuse emission.

What is not yet known is if their host galaxies, and larger, group- and
cluster-scale environments, are similar, and if so, to what extent (O’Dea &
Saikia 2021). The environments of high-excitation LCCs and most CSS-like NLS1s
have not been studied systematically, and also the knowledge of the
environments of beamed NLS1s is far from comprehensive. An additional issue
that must be taken into account in future studies is the selection criteria of
LLCs: in Kunert-Bajraszewska et al. (2010) the sources have 5 GHz flux
densities $>$ 150 mJy, and in Kunert-Bajraszewska et al. (2010) they have 1.4
GHz flux densities $>$ 70 mJy. Most CSS-like NLS1s are fainter than this, thus
in order to have comparable luminosities they must be at higher redshifts,
which might induce some selection biases in the samples. However, the
unification scenario of sub-populations of NLS1s and CSS sources should
definitely be studied further.

####  6.1.4 Large-scale outflows

A third group of AGN-dominated sources are the ones with large-scale outflows.
These sources are rather scarce in our sample, but it might partially be
because without higher-resolution observations they are hard to distinguish
from sources with slow, poorly collimated jets. Also their luminosities are
quite low, so unless they can be identified based on the morphology, or actual
observations of the outflow, they can easily go undetected. Outflows and jets
can also co-exist (for example, J1703+4540, Gu & Chen 2010; Longinotti et al.
2018), and since the radio emission of the jet is generally brighter than the
outflow radio emission, the jet dominates the radio band.

J0804+7248 is an example of a source than could host an outflow instead of a
jet: it is only moderately luminous, its radio morphology it clearly
asymmetric, but does not look collimated, and its [O III] lines are
blueshifted, implying that the NLR is experiencing bulk motion. Another
similar source, J0632+6340, shows a two-sided non-collimated radio morphology
at the same PA as [O III] ionisation cones, and has a very low luminosity
implying that the low-speed jets or outflows have very low power.

###  6.2 Host-dominated sources

About a dozen of our sources seem to be dominated by radio emission that
originates from the host galaxy, that is, from star formation related
processes. They do not show any clear signs of AGN activity in radio, and the
star formation activity is sufficient to explain their radio emission. Their
radio morphologies are generally featureless, and the emission is confined
within the host, or even the host bulge, hinting at circumnuclear star
formation. In some cases the radio morphology is scattered in patches around
the core, and in some cases it seems to trace the host galaxy features. The
sources for which BPT diagram data exists are consistently classified as star-
forming galaxies, suggesting that the stars are the main source of the
ionising continuum. Interestingly, more than half of these sources show a
prominent [O III] wing, indicating that nuclear winds reaching the inner NLR
are present.

###  6.3 Composite sources

Approximately one third of the sources can be classified as composite: in
these sources both the AGN and the host galaxy have a significant contribution
to the total radio emission. These sources show a mixed variety of properties
characteristic for AGN-dominated as well as for star formation dominated
sources. They can show, for example, a flat spectrum core, radio morphology
resembling a jet, or very high luminosity, and on the other hand, for example,
patchy radio emission reminiscent of circumnuclear star formation, PAH
features, or other positive star formation proxies, indicating star formation
even at a starburst level. The different BPT diagrams also often give mixed
results for these sources, indicating that both the AGN and the starlight
continuum are significant. In addition, different kinds of outflows have been
detected in these sources, including UFOs, BLR and NLR outflows, and molecular
outflows. As a cherry on top, one of these sources was identified as a water
maser source.

The ensemble of properties seen in these sources beautifully demonstrates why
NLS1s should not be treated as a class consisting of similar sources, but
should be treated as individuals. Caution should be exercised also when using
simple diagnostic tools, since they cannot properly take into account the
versatility seen in these sources, and might give misleading results when used
on their own. This matter is further discussed in Sect.6.4.

####  6.3.1 Peculiar sources

There are some unique sources that do not fit any of the categories described
earlier. An obvious one is J1337+2423 that in closer inspection turned out not
to be an NLS1, but a BLS1 galaxy (Lee et al. 2013). This source highlights a
possibly bigger issue in the identification on different types of Seyferts:
most of them have been classified based on quite low signal-to-noise (S/N)
ratio spectra, and often using automated algorithms. Subtle, but crucial,
details of the emission lines, such as the wings, cannot be properly modelled
using low S/N data, but can easily lead to misclassification of NLS1s and
BLS1s, and confusion with intermediate Seyferts (Järvelä et al. 2020; Peruzzi
et al. 2021). Mixed samples will add undesired noise to statistical population
studies, and hinder our efforts to understand the nature of the different
classes of AGN.

J1038+4227 is another extremely interesting object, since its radio morphology
does not resemble anything we see in other NLS1s. However, its optical
spectrum clearly identifies it as one: its FWHM(H$\beta$) = 1979 km s-1, and
it shows very strong Fe II multiplets. Were it a more evolved source, for
example, an FSRQ with a flattened BLR seen face-on, we would not expect to see
iron in the spectrum (Marziani et al. 2018b; Berton et al. 2020a). J1038+4227
is a luminous source, and since the core is responsible for most of its radio
emission, and shows a steep spectral index it would be consistent with a CSS
classification. The symmetrical, almost circular, faint radio emission
reaching well outside the host galaxy is what sets J1038+4227 apart from other
NLS1s. Due to its huge size its origin is probably the AGN, and it seems
plausible that it is a radio lobe. It resembles the kpc-scale morphologies
seen in some blazars (Kharb et al. 2010), but is considerably fainter.
Unfortunately, this structure is too faint to estimate its spectral index, and
derive any information about its age or whether its a relic or still
replenished. However, taking into account its circular shape it seems like the
jet must have pointed close to our line of sight, which does not seem to be
the case right now, otherwise we would expect to see a flatter spectral index
in the core. The nature of this source remains unclear for now, but it
definitely deserves to be studied in detail in the future.

The third source that stands out is J1713+3523. It has a steep spectral index
at least up to 5.2 GHz, and does not show strong signs of AGN activity at that
frequency. In fact, its 5.2 GHz radio emission could be explained with star
formation alone. However, a 22 GHz detection of 138 mJy is undeniable proof
that this source hosts at least moderately powerful jets, and it furthermore
yields a 5.2-22 GHz spectral index close to the theoretical SSA maximum. It is
interesting that no signs of the jets are seen at lower frequencies. This
source somewhat resembles an extraordinary group of previously radio-quiet or
totally radio-silent NLS1s, first discovered at 37 GHz due to their bright
(hundreds of mJy), and variable radio emission, characteristic for
relativistic jets (Lähteenmäki et al. 2018). These sources were later observed
with the JVLA, and all of them showed a steep 1.6-9.0 GHz radio spectrum, with
no signs of possible jets (Berton et al. 2020b). The nature of these sources
is still somewhat unclear, but absorption, either SSA or free-free has been
suggested as the culprit behind their unusual radio spectra. More observations
of J1713+2523 are certainly needed to explain the observed properties of this
source.

###  6.4 Implications and issues

The results of this paper strengthen the view of NLS1s as a wonderfully
diverse class of sources. However, this versatility might also cause
additional trouble, especially when studying larger samples of sources, rather
than individuals, and thus caution is required when dealing with large samples
and/or statistical studies. Below we discuss the biggest issues resulting from
this diversity.

####  6.4.1 NLS1s are _not_ a homogeneous class

It cannot be stressed enough that the variety of sources under the NLS1
umbrella is vast. After all, the classification is based solely on the
narrowness of the permitted lines, that indicates an undermassive black hole,
and the oxygen-to-hydrogen ratio, that ensures they are Type 1 sources. This
does not tell us anything about the activity state of the AGN, and it ranges
from radio-silent sources to rare individuals that are able to maintain fully
developed relativistic jets. It also does not reflect the properties of the
host galaxy, and in NLS1s the AGN rarely dominates the whole electromagnetic
spectrum over the host, unlike in quasars or blazars. Even if this might be
true also for some NLS1s, the majority does not fit one mould, as demonstrated
in this paper. It should also be kept in mind that the radio view of NLS1s is
very different than the optical view, and whereas their optical spectra are
similar, the radio properties absolutely are not.

####  6.4.2 Advice against using simple proxies

The amasing diversity of NLS1s, and the varying contributions of the AGN and
the host along the electromagnetic spectrum can have unexpected effects when
using relations and proxies possibly calibrated with more homogeneous or
simple samples of sources. This has been continuously demonstrated regarding
the use of the radio-loudness parameter, that might have originally served as
an acceptable first order estimate of the activity of the AGN, but fails to
perform when more complicated sources are included in the scenario (Järvelä et
al. 2017; Padovani 2017).

In this paper we used several star formation diagnostics based on mid-infrared
data, or relations between radio and mid-infrared. It soon became evident that
these proxies were not ideal tools to estimate the strength of star formation
in our sources. They often show results contradicting each other, as well as
results in the literature obtained with more reliable methods, for example,
PAH features. It is possibly the complexity of NLS1s that fool these proxies,
and especially in mid-infrared the presence of an emphasised dust component
that can mimic the properties of emission produced by star formation.
Nevertheless, we decided to include these diagnostics in our results to
highlight their behaviour in the case of NLS1s, and as a general caution for
future, especially statistical, studies using similar diagnostics.

##  7 Summary and conclusions

In this paper we analysed 44 NLS1 galaxies by means of spatially resolved
radio spectral index maps. In addition, we employed mid-infrared diagnostics
to estimate the star formation activity in their hosts. With these tools,
complemented by archival data, we estimated the predominant source or sources
of radio emission in them. Based on these analyses our main conclusions can be
summarised as:

  1. 1.

NLS1s are an exceptionally heterogeneous class of AGN, and in radio band they
range from individuals with blazar- and radio galaxy -like characteristics to
sources whose radio emission is dominated by the star formation activity in
the host galaxy. Furthermore, some of them could actually be classified as CSS
sources in addition to the NLS1 classification, thus providing further
confirmation of the scenario in which NLS1s are AGN in an early evolutionary
stage.

  2. 2.

Due to this diversity and the varying contributions of different sources of
radio as well as mid-infrared emission NLS1s are not ideal targets for simple
proxies, such as the radio loudness parameter, or star formation diagnostics.

  3. 3.

As the NLS1 classification is based solely on their optical spectrum, and
includes all Type 1 sources with narrow permitted lines, it does not reflect
their other physical properties, for example, in the radio band. As a result
the NLS1 classification does not tell us much as these sources do not
necessarily share similar physical properties. Also the threshold of
FWHM(H$\beta$) = 2000 km s-1is somewhat arbitrary, and should not be regarded
as a real physical limit. These remarks should be kept in mind especially when
studying them as a class, or using the general NLS1 population in wider AGN
studies.

NLS1s are complicated sources, and any conclusions regarding the whole class
cannot and should not be drawn lightly. Instead, these sources should be first
and foremost studied as individuals, at least until we understand their nature
and distinctive subclasses and their characteristics better. As demonstrated
by this study, the majority of them require very high-quality data and
detailed data analysis to disentangle the contribution of the different
elements. And even still, many of the sources in our sample will require
further observations, for example, in radio at higher resolution, or integral
field spectroscopy, to unquestionably determine their nature, and their
physical properties. Future facilities, such as the Cherenkov Telescope Array
(CTA), the Advanced Telescope for High Energy Astrophysics (ATHENA), and
Euclid, will also be essential to improve our understanding of this peculiar
class of objects. Despite the complexity, the NLS1s definitely are a class
worth investigating. The jetted NLS1s are an ideal laboratory for studying the
launching of the jets, and the first stages of the evolution of more powerful
FSRQs and high-excitation radio galaxies. In general in NLS1s, the interplay
of the host and several AGN-related phenomena, for example, jets and outflows,
offers a great opportunity to examine the AGN feedback in action.

###### Acknowledgements.

E.J. is a current ESA research fellow. E.C. acknowledges support from ANID
project Basal AFB-170002. The National Radio Astronomy Observatory is a
facility of the National Science Foundation operated under cooperative
agreement by Associated Universities, Inc. This publication makes use of data
products from the Wide-field Infrared Survey Explorer, which is a joint
project of the University of California, Los Angeles, and the Jet Propulsion
Laboratory/California Institute of Technology, funded by the National
Aeronautics and Space Administration. The Pan-STARRS1 Surveys (PS1) and the
PS1 public science archive have been made possible through contributions by
the Institute for Astronomy, the University of Hawaii, the Pan-STARRS Project
Office, the Max-Planck Society and its participating institutes, the Max
Planck Institute for Astronomy, Heidelberg and the Max Planck Institute for
Extraterrestrial Physics, Garching, The Johns Hopkins University, Durham
University, the University of Edinburgh, the Queen’s University Belfast, the
Harvard-Smithsonian Center for Astrophysics, the Las Cumbres Observatory
Global Telescope Network Incorporated, the National Central University of
Taiwan, the Space Telescope Science Institute, the National Aeronautics and
Space Administration under Grant No. NNX08AR22G issued through the Planetary
Science Division of the NASA Science Mission Directorate, the National Science
Foundation Grant No. AST-1238877, the University of Maryland, Eotvos Lorand
University (ELTE), the Los Alamos National Laboratory, and the Gordon and
Betty Moore Foundation. Funding for the Sloan Digital Sky Survey (SDSS) has
been provided by the Alfred P. Sloan Foundation, the Participating
Institutions, the National Aeronautics and Space Administration, the National
Science Foundation, the U.S. Department of Energy, the Japanese
Monbukagakusho, and the Max Planck Society. The SDSS Web site is
http://www.sdss.org/. The SDSS is managed by the Astrophysical Research
Consortium (ARC) for the Participating Institutions. The Participating
Institutions are The University of Chicago, Fermilab, the Institute for
Advanced Study, the Japan Participation Group, The Johns Hopkins University,
the Korean Scientist Group, Los Alamos National Laboratory, the Max-Planck-
Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics
(MPA), New Mexico State University, University of Pittsburgh, University of
Portsmouth, Princeton University, the United States Naval Observatory, and the
University of Washington. Thank you coffee and sparkling wine. This research
has made use of the NASA/IPAC Extragalactic Database (NED), which is operated
by the Jet Propulsion Laboratory, California Institute of Technology, under
contract with the National Aeronautics and Space Administration. This research
has made use of the SIMBAD database, operated at CDS, Strasbourg, France.

## References

  * Alonso-Herrero et al. (2016) Alonso-Herrero, A., Poulton, R., Roche, P. F., et al. 2016, MNRAS, 463, 2405
  * Antón et al. (2008) Antón, S., Browne, I. W. A., & Marchã, M. J. 2008, A&A, 490, 583
  * Asmus (2019) Asmus, D. 2019, MNRAS, 489, 2177
  * Asmus et al. (2016) Asmus, D., Hönig, S. F., & Gandhi, P. 2016, ApJ, 822, 109
  * Baldi et al. (2018) Baldi, R. D., Williams, D. R. A., McHardy, I. M., et al. 2018, MNRAS, 476, 3478
  * Baldwin et al. (1981) Baldwin, J. A., Phillips, M. M., & Terlevich, R. 1981, PASP, 93, 5
  * Barth et al. (2008) Barth, A. J., Bentz, M. C., Greene, J. E., & Ho, L. C. 2008, ApJ, 683, L119
  * Baum et al. (1993) Baum, S. A., O’Dea, C. P., Dallacassa, D., de Bruyn, A. G., & Pedlar, A. 1993, ApJ, 419, 553
  * Becker et al. (1991) Becker, R. H., White, R. L., & Edwards, A. L. 1991, ApJS, 75, 1
  * Becker et al. (1995) Becker, R. H., White, R. L., & Helfand, D. J. 1995, ApJ, 450, 559
  * Behar et al. (2015) Behar, E., Baldi, R. D., Laor, A., et al. 2015, MNRAS, 451, 517
  * 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 & Järvelä (2021) Berton, M. & Järvelä, E. 2021, Universe, 7, 188
  * Berton et al. (2020b) Berton, M., Järvelä, E., Crepaldi, L., et al. 2020b, A&A, 636, A64
  * Bhatnagar et al. (2013) Bhatnagar, S., Rau, U., & Golap, K. 2013, ApJ, 770, 91
  * Bian & Huang (2010) Bian, W.-H. & Huang, K. 2010, MNRAS, 401, 507
  * Bicay et al. (1995) Bicay, M. D., Kojoian, G., Seal, J., Dickinson, D. F., & Malkan, M. A. 1995, ApJS, 98, 369
  * Bicknell et al. (1997) Bicknell, G. V., Dopita, M. A., & O’Dea, C. P. O. 1997, ApJ, 485, 112
  * Boller et al. (2001) Boller, T., Keil, R., Trümper, J., et al. 2001, A&A, 365, L146
  * Boroson & Green (1992) Boroson, T. A. & Green, R. F. 1992, ApJS, 80, 109
  * Boyle et al. (2007) Boyle, B. J., Cornwell, T. J., Middelberg, E., et al. 2007, MNRAS, 376, 1182
  * Brinkmann et al. (1995) Brinkmann, W., Siebert, J., Reich, W., et al. 1995, A&AS, 109, 147
  * Buhariwalla et al. (2020) Buhariwalla, M. Z., Waddell, S. G. H., Gallo, L. C., Grupe, D., & Komossa, S. 2020, ApJ, 901, 118
  * Burtscher et al. (2009) Burtscher, L., Jaffe, W., Raban, D., et al. 2009, ApJ, 705, L53
  * Buta (2017) Buta, R. J. 2017, MNRAS, 471, 4027
  * 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
  * Chainakun & Young (2017) Chainakun, P. & Young, A. J. 2017, MNRAS, 465, 3965
  * Chamani et al. (2021) Chamani, W., Savolainen, T., Hada, K., & Xu, M. H. 2021, A&A, 652, A14
  * Chen et al. (2020) Chen, S., Järvelä, E., Crepaldi, L., et al. 2020, MNRAS, 498, 1278
  * Christopoulou et al. (1997) Christopoulou, P. E., Holloway, A. J., Steffen, W., et al. 1997, MNRAS, 284, 385
  * Cohen et al. (2007) Cohen, A. S., Lane, W. M., Cotton, W. D., et al. 2007, AJ, 134, 1245
  * Condon (1992) Condon, J. J. 1992, ARA&A, 30, 575
  * 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
  * Crenshaw et al. (2003) Crenshaw, D. M., Kraemer, S. B., & Gabel, J. R. 2003, AJ, 126, 1690
  * Croft et al. (2010) Croft, S., Bower, G. C., Ackermann, R., et al. 2010, ApJ, 719, 45
  * Czerny et al. (2009) Czerny, B., Siemiginowska, A., Janiuk, A., Nikiel-Wroczyński, B., & Stawarz, Ł. 2009, ApJ, 698, 840
  * da Cunha et al. (2008) da Cunha, E., Charlot, S., & Elbaz, D. 2008, MNRAS, 388, 1595
  * Dalla Bontà et al. (2020) Dalla Bontà, E., Peterson, B. M., Bentz, M. C., et al. 2020, ApJ, 903, 112
  * D’Ammando et al. (2018) D’Ammando, F., Acosta-Pulido, J. A., Capetti, A., et al. 2018, MNRAS, 478, L66
  * D’Ammando et al. (2017) D’Ammando, F., Acosta-Pulido, J. A., Capetti, A., et al. 2017, MNRAS, 469, L11
  * Decarli et al. (2008) Decarli, R., Dotti, M., Fontana, M., & Haardt, F. 2008, MNRAS, 386, L15
  * Deller & Middelberg (2014) Deller, A. T. & Middelberg, E. 2014, AJ, 147, 14
  * Deo et al. (2006) Deo, R. P., Crenshaw, D. M., & Kraemer, S. B. 2006, AJ, 132, 321
  * Dicken et al. (2009) Dicken, D., Tadhunter, C., Axon, D., et al. 2009, ApJ, 694, 268
  * Doi et al. (2013) Doi, A., Asada, K., Fujisawa, K., et al. 2013, ApJ, 765, 69
  * Doi et al. (2011) Doi, A., Asada, K., & Nagai, H. 2011, ApJ, 738, 126
  * Doi et al. (2007) Doi, A., Fujisawa, K., Inoue, M., et al. 2007, PASJ, 59, 703
  * Doi et al. (2016) Doi, A., Oyama, T., Kono, Y., et al. 2016, PASJ, 68, 73
  * Doi et al. (2015) Doi, A., Wajima, K., Hagiwara, Y., & Inoue, M. 2015, ApJ, 798, L30
  * Dressel & Condon (1978) Dressel, L. L. & Condon, J. J. 1978, ApJS, 36, 53
  * Du & Wang (2019) Du, P. & Wang, J.-M. 2019, ApJ, 886, 42
  * Du et al. (2018) Du, P., Zhang, Z.-X., Wang, K., et al. 2018, ApJ, 856, 6
  * Edelson (1987) Edelson, R. A. 1987, ApJ, 313, 651
  * Edelson & Malkan (1986) Edelson, R. A. & Malkan, M. A. 1986, ApJ, 308, 59
  * Ellison et al. (2011) Ellison, S. L., Patton, D. R., Mendel, J. T., & Scudder, J. M. 2011, MNRAS, 418, 2043
  * Esparza-Arredondo et al. (2020) Esparza-Arredondo, D., Osorio-Clavijo, N., González-Martín, O., et al. 2020, ApJ, 905, 29
  * Esquej et al. (2014) Esquej, P., Alonso-Herrero, A., González-Martín, O., et al. 2014, ApJ, 780, 86
  * Faucher-Giguère & Quataert (2012) Faucher-Giguère, C.-A. & Quataert, E. 2012, MNRAS, 425, 605
  * Fischer et al. (2013) Fischer, T. C., Crenshaw, D. M., Kraemer, S. B., & Schmitt, H. R. 2013, ApJS, 209, 1
  * Flewelling et al. (2020) Flewelling, H. A., Magnier, E. A., Chambers, K. C., et al. 2020, ApJS, 251, 7
  * Foschini (2011) Foschini, L. 2011, in Narrow-Line Seyfert 1 Galaxies and their Place in the Universe, https://pos.sissa.it/126/024/pdf
  * Foschini (2017) Foschini, L. 2017, Frontiers in Astronomy and Space Sciences, 4, 6
  * 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
  * Fraix-Burnet et al. (2017) Fraix-Burnet, D., Marziani, P., D’Onofrio, M., & Dultzin, D. 2017, Frontiers in Astronomy and Space Sciences, 4, 1
  * Fritz et al. (2006) Fritz, J., Franceschini, A., & Hatziminaoglou, E. 2006, MNRAS, 366, 767
  * Gallimore et al. (2006) Gallimore, J. F., Axon, D. J., O’Dea, C. P., Baum, S. A., & Pedlar, A. 2006, AJ, 132, 546
  * Gallo et al. (2006) Gallo, L. C., Edwards, P. G., Ferrero, E., et al. 2006, MNRAS, 370, 245
  * Ganci et al. (2019) Ganci, V., Marziani, P., D’Onofrio, M., et al. 2019, A&A, 630, A110
  * Giroletti & Panessa (2009) Giroletti, M. & Panessa, F. 2009, ApJ, 706, L260
  * Giroletti et al. (2017) Giroletti, M., Panessa, F., Longinotti, A. L., et al. 2017, A&A, 600, A87
  * Goad et al. (2012) Goad, M. R., Korista, K. T., & Ruff, A. J. 2012, MNRAS, 426, 3086
  * Gonzalez Delgado & Perez (1996) Gonzalez Delgado, R. M. & Perez, E. 1996, MNRAS, 278, 737
  * González-Martín et al. (2019) González-Martín, O., Masegosa, J., García-Bernete, I., et al. 2019, ApJ, 884, 11
  * Goodrich (1989) Goodrich, R. W. 1989, ApJ, 342, 224
  * Gregory & Condon (1991) Gregory, P. C. & Condon, J. J. 1991, ApJS, 75, 1011
  * Gruppioni et al. (2016) Gruppioni, C., Berta, S., Spinoglio, L., et al. 2016, MNRAS, 458, 4297
  * Gu & Chen (2010) Gu, M. & Chen, Y. 2010, AJ, 139, 2612
  * Gu et al. (2015) Gu, M., Chen, Y., Komossa, S., et al. 2015, ApJS, 221, 3
  * Gupta et al. (2013a) Gupta, A., Mathur, S., Krongold, Y., & Nicastro, F. 2013a, ApJ, 768, 141
  * Gupta et al. (2013b) Gupta, A., Mathur, S., Krongold, Y., & Nicastro, F. 2013b, ApJ, 772, 66
  * Gürkan et al. (2019) Gürkan, G., Hardcastle, M. J., Best, P. N., et al. 2019, A&A, 622, A11
  * Hagiwara et al. (2018) Hagiwara, Y., Doi, A., Hachisuka, K., & Horiuchi, S. 2018, PASJ, 70, 54
  * Hamilton et al. (2021) Hamilton, T. S., Berton, M., Antón, S., et al. 2021, MNRAS, 504, 5188
  * Harrison et al. (2012) Harrison, C. M., Alexander, D. M., Swinbank, A. M., et al. 2012, MNRAS, 426, 1073
  * Heinz & Sunyaev (2003) Heinz, S. & Sunyaev, R. A. 2003, MNRAS, 343, L59
  * Heisler & Vader (1994) Heisler, C. A. & Vader, J. P. 1994, AJ, 107, 35
  * Hönig et al. (2013) Hönig, S. F., Kishimoto, M., Tristram, K. R. W., et al. 2013, ApJ, 771, 87
  * Hutchings & Craven (1988) Hutchings, J. B. & Craven, S. E. 1988, AJ, 95, 677
  * Ibar et al. (2008) Ibar, E., Cirasuolo, M., Ivison, R., et al. 2008, MNRAS, 386, 953
  * Igo et al. (2020) Igo, Z., Parker, M. L., Matzeu, G. A., et al. 2020, MNRAS, 493, 1088
  * Jackson et al. (2020) Jackson, T. M., Rosario, D. J., Alexander, D. M., et al. 2020, MNRAS, 498, 2323
  * Jarrett et al. (2011) Jarrett, T. H., Cohen, M., Masci, F., et al. 2011, ApJ, 735, 112
  * Järvelä et al. (2020) Järvelä, E., Berton, M., Ciroi, S., et al. 2020, A&A, 636, L12
  * Järvelä et al. (2021) Järvelä, E., Berton, M., & Crepaldi, L. 2021, arXiv e-prints, arXiv:2108.08521
  * 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
  * Jiang et al. (2010) Jiang, Y.-F., Ciotti, L., Ostriker, J. P., & Spitkovsky, A. 2010, ApJ, 711, 125
  * Jin et al. (2021) Jin, C., Done, C., & Ward, M. 2021, MNRAS, 500, 2475
  * Jones et al. (2017) Jones, S., McHardy, I., & Maccarone, T. J. 2017, MNRAS, 465, 1336
  * Kellermann et al. (1989) Kellermann, K. I., Sramek, R., Schmidt, M., Shaffer, D. B., & Green, R. 1989, AJ, 98, 1195
  * Kellermann et al. (1994) Kellermann, K. I., Sramek, R. A., Schmidt, M., Green, R. F., & Shaffer, D. B. 1994, AJ, 108, 1163
  * Khanna et al. (2016) Khanna, S., Kaastra, J. S., & Mehdipour, M. 2016, A&A, 586, A2
  * Kharb et al. (2010) Kharb, P., Lister, M. L., & Cooper, N. J. 2010, ApJ, 710, 764
  * King et al. (2011) King, A. L., Miller, J. M., Cackett, E. M., et al. 2011, ApJ, 729, 19
  * Kinney et al. (2000) Kinney, A. L., Schmitt, H. R., Clarke, C. J., et al. 2000, ApJ, 537, 152
  * Kollatschny (2003) Kollatschny, W. 2003, A&A, 407, 461
  * Kollatschny & Zetzl (2011) Kollatschny, W. & Zetzl, M. 2011, Nature, 470, 366
  * 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. (2018) Komossa, S., Xu, D. W., & Wagner, A. Y. 2018, MNRAS, 477, 5115
  * Kotilainen et al. (2016) Kotilainen, J. K., León-Tavares, J., Olguín-Iglesias, A., et al. 2016, ApJ, 832, 157
  * Kozieł-Wierzbowska et al. (2021) Kozieł-Wierzbowska, D., Vale Asari, N., Stasińska, G., et al. 2021, ApJ, 910, 64
  * Krongold et al. (2001) Krongold, Y., Dultzin-Hacyan, D., & Marziani, P. 2001, AJ, 121, 702
  * Kukula et al. (1998) Kukula, M. J., Dunlop, J. S., Hughes, D. H., & Rawlings, S. 1998, MNRAS, 297, 366
  * Kukula et al. (1995) Kukula, M. J., Pedlar, A., Baum, S. A., & O’Dea, C. P. 1995, MNRAS, 276, 1262
  * Kunert et al. (2002) Kunert, M., Marecki, A., Spencer, R. E., Kus, A. J., & Niezgoda, J. 2002, A&A, 391, 47
  * Kunert-Bajraszewska et al. (2010) Kunert-Bajraszewska, M., Gawroński, M. P., Labiano, A., & Siemiginowska, A. 2010, MNRAS, 408, 2261
  * Kunert-Bajraszewska & Labiano (2010) Kunert-Bajraszewska, M. & Labiano, A. 2010, MNRAS, 408, 2279
  * Lähteenmäki et al. (2018) Lähteenmäki, A., Järvelä, E., Ramakrishnan, V., et al. 2018, A&A, 614, L1
  * LaMassa et al. (2012) LaMassa, S. M., Heckman, T. M., Ptak, A., et al. 2012, ApJ, 758, 1
  * LaMassa et al. (2016) LaMassa, S. M., Ricarte, A., Glikman, E., et al. 2016, ApJ, 820, 70
  * Lamperti et al. (2020) Lamperti, I., Saintonge, A., Koss, M., et al. 2020, ApJ, 889, 103
  * Laor (2000) Laor, A. 2000, ApJ, 543, L111
  * Laurent-Muehleisen et al. (1997) Laurent-Muehleisen, S. A., Kollgaard, R. I., Ryan, P. J., et al. 1997, A&AS, 122, 235
  * Lee et al. (2013) Lee, J. C., Kriss, G. A., Chakravorty, S., et al. 2013, MNRAS, 430, 2650
  * Leftley et al. (2019) Leftley, J. H., Hönig, S. F., Asmus, D., et al. 2019, ApJ, 886, 55
  * Lianou et al. (2019) Lianou, S., Barmby, P., Mosenkov, A. A., Lehnert, M., & Karczewski, O. 2019, A&A, 631, A38
  * Liao et al. (2017) Liao, N.-H., Liang, Y.-F., Weng, S.-S., et al. 2017, Discovery of $\gamma$-ray emission from steep radio spectrum NLS1s
  * Lister (2018) Lister, M. 2018, in Revisiting Narrow-Line Seyfert 1 Galaxies and their Place in the Universe, 22
  * Longinotti et al. (2018) Longinotti, A. L., Vega, O., Krongold, Y., et al. 2018, ApJ, 867, L11
  * Lyu & Rieke (2018) Lyu, J. & Rieke, G. H. 2018, ApJ, 866, 92
  * Mahajan et al. (2019) Mahajan, S., Ashby, M. L. N., Willner, S. P., et al. 2019, MNRAS, 482, 560
  * Maitra et al. (2011) Maitra, D., Miller, J. M., Markoff, S., & King, A. 2011, ApJ, 735, 107
  * Mallick et al. (2016) Mallick, L., Dewangan, G. C., Gandhi, P., Misra, R., & Kembhavi, A. K. 2016, MNRAS, 460, 1705
  * Martínez-Paredes et al. (2019) Martínez-Paredes, M., Aretxaga, I., González-Martín, O., et al. 2019, ApJ, 871, 190
  * 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. (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
  * Meena et al. (2021) Meena, B., Crenshaw, D. M., Schmitt, H. R., et al. 2021, ApJ, 916, 31
  * Miller et al. (1993) Miller, P., Rawlings, S., & Saunders, R. 1993, MNRAS, 263, 425
  * Moran (2000) Moran, E. C. 2000, NewAR, 44, 527
  * Muñoz Marín et al. (2007) Muñoz Marín, V. M., González Delgado, R. M., Schmitt, H. R., et al. 2007, AJ, 134, 648
  * Mulchaey et al. (1996) Mulchaey, J. S., Wilson, A. S., & Tsvetanov, Z. 1996, ApJS, 102, 309
  * Nagar et al. (1999) Nagar, N. M., Wilson, A. S., Mulchaey, J. S., & Gallimore, J. F. 1999, ApJS, 120, 209
  * Neumann et al. (1994) Neumann, M., Reich, W., Fuerst, E., et al. 1994, A&AS, 106, 303
  * Norris et al. (1990) Norris, R. P., Allen, D. A., Sramek, R. A., Kesteven, M. J., & Troup, E. R. 1990, ApJ, 359, 291
  * O’Dea (1998) O’Dea, C. P. 1998, PASP, 110, 493
  * O’Dea & Saikia (2021) O’Dea, C. P. & Saikia, D. J. 2021, A&A Rev., 29, 3
  * Ohta et al. (2007) Ohta, K., Aoki, K., Kawaguchi, T., & Kiuchi, G. 2007, ApJS, 169, 1
  * Olguín-Iglesias et al. (2020) Olguín-Iglesias, A., Kotilainen, J., & Chavushyan, V. 2020, MNRAS, 492, 1450
  * Olguín-Iglesias et al. (2017) Olguín-Iglesias, A., Kotilainen, J. K., León Tavares, J., Chavushyan, V., & Añorve, C. 2017, MNRAS, 467, 3712
  * Orienti & Prieto (2010) Orienti, M. & Prieto, M. A. 2010, MNRAS, 401, 2599
  * 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
  * Panda et al. (2018) Panda, S., Czerny, B., Adhikari, T. P., et al. 2018, ApJ, 866, 115
  * Panessa et al. (2019) Panessa, F., Baldi, R. D., Laor, A., et al. 2019, Nature Astronomy, 3, 387
  * Parker et al. (2020) Parker, M. L., Matzeu, G. A., Alston, W. N., et al. 2020, MNRAS, 498, L140
  * Parker et al. (2018) Parker, M. L., Matzeu, G. A., Guainazzi, M., et al. 2018, MNRAS, 480, 2365
  * Parra et al. (2010) Parra, R., Conway, J. E., Aalto, S., et al. 2010, ApJ, 720, 555
  * Peng et al. (2014) Peng, Z.-X., Chen, Y.-M., Gu, Q.-S., & Zhang, K. 2014, Research in Astronomy and Astrophysics, 14, 913
  * Peruzzi et al. (2021) Peruzzi, T., Pasquato, M., Ciroi, S., et al. 2021, A&A, 652, A19
  * Peterson (2011) Peterson, B. M. 2011, ArXiv e-prints [arXiv:1109.4181]
  * Petrov (2013) Petrov, L. 2013, AJ, 146, 5
  * Planck Collaboration et al. (2013) Planck Collaboration, Ade, P. A. R., Aghanim, N., et al. 2013, A&A, 550, A133
  * Popescu et al. (2011) Popescu, C. C., Tuffs, R. J., Dopita, M. A., et al. 2011, A&A, 527, A109
  * Popović et al. (2009) Popović, L. Č., Smirnova, A. A., Kovačević, J., Moiseev, A. V., & Afanasiev, V. L. 2009, AJ, 137, 3548
  * Pounds et al. (1995) Pounds, K. A., Done, C., & Osborne, J. P. 1995, MNRAS, 277, L5
  * Pounds & King (2013) Pounds, K. A. & King, A. R. 2013, MNRAS, 433, 1369
  * Proga (2007) Proga, D. 2007, in Astronomical Society of the Pacific Conference Series, Vol. 373, The Central Engine of Active Galactic Nuclei, ed. L. C. Ho & J. W. Wang, 267
  * Rakshit et al. (2021) Rakshit, S., Schramm, M., Stalin, C. S., et al. 2021, MNRAS[arXiv:2103.16521]
  * Rakshit et al. (2018) Rakshit, S., Stalin, C. S., Hota, A., & Konar, C. 2018, ApJ, 869, 173
  * Rau & Cornwell (2011) Rau, U. & Cornwell, T. J. 2011, A&A, 532, A71
  * Rodríguez-Ardila et al. (2005) Rodríguez-Ardila, A., Contini, M., & Viegas, S. M. 2005, MNRAS, 357, 220
  * Rodríguez-Ardila & Mazzalay (2006) Rodríguez-Ardila, A. & Mazzalay, X. 2006, MNRAS, 367, L57
  * Rodríguez-Ardila & Viegas (2003) Rodríguez-Ardila, A. & Viegas, S. M. 2003, MNRAS, 340, L33
  * Romano et al. (2018) Romano, P., Vercellone, S., Foschini, L., et al. 2018, MNRAS, 481, 5046
  * Rosario et al. (2013) Rosario, D. J., Burtscher, L., Davies, R., et al. 2013, ApJ, 778, 94
  * Ruschel-Dutra et al. (2017) Ruschel-Dutra, D., Rodríguez Espinosa, J. M., González Martín, O., Pastoriza, M., & Riffel, R. 2017, MNRAS, 466, 3353
  * Saikia et al. (2018) Saikia, P., Körding, E., Coppejans, D. L., et al. 2018, A&A, 616, A152
  * Sanfrutos et al. (2018) Sanfrutos, M., Longinotti, A. L., Krongold, Y., Guainazzi, M., & Panessa, F. 2018, ApJ, 868, 111
  * Sani et al. (2010) Sani, E., Lutz, D., Risaliti, G., et al. 2010, MNRAS, 403, 1246
  * Sargsyan et al. (2012) Sargsyan, L., Lebouteiller, V., Weedman, D., et al. 2012, ApJ, 755, 171
  * Sargsyan et al. (2014) Sargsyan, L., Samsonyan, A., Lebouteiller, V., et al. 2014, ApJ, 790, 15
  * Sargsyan & Weedman (2009) Sargsyan, L. A. & Weedman, D. W. 2009, ApJ, 701, 1398
  * Scheuer & Readhead (1979) Scheuer, P. A. G. & Readhead, A. C. S. 1979, Nature, 277, 182
  * Schmitt et al. (2001a) Schmitt, H. R., Antonucci, R. R. J., Ulvestad, J. S., et al. 2001a, ApJ, 555, 663
  * Schmitt et al. (2003) Schmitt, H. R., Donley, J. L., Antonucci, R. R. J., Hutchings, J. B., & Kinney, A. L. 2003, ApJS, 148, 327
  * Schmitt et al. (2001b) Schmitt, H. R., Ulvestad, J. S., Antonucci, R. R. J., & Kinney, A. L. 2001b, ApJS, 132, 199
  * Schönell et al. (2014) Schönell, A. J., Riffel, R. A., Storchi-Bergmann, T., & Winge, C. 2014, MNRAS, 445, 414
  * Seyfert (1943) Seyfert, C. K. 1943, ApJ, 97, 28
  * Shangguan et al. (2020) Shangguan, J., Ho, L. C., Bauer, F. E., Wang, R., & Treister, E. 2020, ApJ, 899, 112
  * Shen & Ho (2014) Shen, Y. & Ho, L. C. 2014, Nature, 513, 210
  * Shi et al. (2007) Shi, Y., Ogle, P., Rieke, G. H., et al. 2007, ApJ, 669, 841
  * Shirazi & Brinchmann (2012) Shirazi, M. & Brinchmann, J. 2012, MNRAS, 421, 1043
  * Shuder & Osterbrock (1981) Shuder, J. M. & Osterbrock, D. E. 1981, ApJ, 250, 55
  * Silpa et al. (2020) Silpa, S., Kharb, P., Ho, L. C., et al. 2020, MNRAS, 499, 5826
  * Smethurst et al. (2019) Smethurst, R. J., Simmons, B. D., Lintott, C. J., & Shanahan, J. 2019, MNRAS, 489, 4016
  * Snellen et al. (2004) Snellen, I. A. G., Mack, K. H., Schilizzi, R. T., & Tschager, W. 2004, MNRAS, 348, 227
  * Stern & Laor (2013) Stern, J. & Laor, A. 2013, MNRAS, 431, 836
  * Sulentic et al. (2000a) Sulentic, J. W., Marziani, P., & Dultzin-Hacyan, D. 2000a, ARA&A, 38, 521
  * Sulentic et al. (2000b) Sulentic, J. W., Zwitter, T., Marziani, P., & Dultzin-Hacyan, D. 2000b, ApJ, 536, L5
  * Tan et al. (2019) Tan, Q.-H., Gao, Y., Kohno, K., et al. 2019, ApJ, 887, 24
  * Taniguchi (1999) Taniguchi, Y. 1999, ApJ, 524, 65
  * Tarchi et al. (2011) Tarchi, A., Castangia, P., Columbano, A., Panessa, F., & Braatz, J. A. 2011, A&A, 532, A125
  * Teng et al. (2013) Teng, S. H., Veilleux, S., & Baker, A. J. 2013, ApJ, 765, 95
  * Terrazas et al. (2017) Terrazas, B. A., Bell, E. F., Woo, J., & Henriques, B. M. B. 2017, ApJ, 844, 170
  * Thean et al. (2000) Thean, A., Pedlar, A., Kukula, M. J., Baum, S. A., & O’Dea, C. P. 2000, MNRAS, 314, 573
  * Tombesi et al. (2010) Tombesi, F., Cappi, M., Reeves, J. N., et al. 2010, A&A, 521, A57
  * Ulvestad et al. (1995) Ulvestad, J. S., Antonucci, R. R. J., & Goodrich, R. W. 1995, AJ, 109, 81
  * Ulvestad & Wilson (1984) Ulvestad, J. S. & Wilson, A. S. 1984, ApJ, 285, 439
  * Urrutia et al. (2008) Urrutia, T., Lacy, M., & Becker, R. H. 2008, ApJ, 674, 80
  * Wang et al. (2016) Wang, F., Du, P., Hu, C., et al. 2016, ApJ, 824, 149
  * Whittle (1992) Whittle, M. 1992, ApJS, 79, 49
  * Wiegert et al. (2015) Wiegert, T., Irwin, J., Miskolczi, A., et al. 2015, AJ, 150, 81
  * Wright et al. (2010) Wright, E. L., Eisenhardt, P. R. M., Mainzer, A. K., et al. 2010, AJ, 140, 1868
  * Xie et al. (2018) Xie, Y., Ho, L. C., Li, A., & Shangguan, J. 2018, ApJ, 860, 154
  * Yang et al. (2020) Yang, X., Yao, S., Yang, J., et al. 2020, ApJ, 904, 200
  * Yao & Komossa (2021) Yao, S. & Komossa, S. 2021, MNRAS, 501, 1384
  * Young et al. (2014) Young, J. E., Eracleous, M., Shemmer, O., et al. 2014, MNRAS, 438, 217
  * Yuan et al. (2010) Yuan, W., Liu, B. F., Zhou, H., & Wang, T. G. 2010, ApJ, 723, 508
  * Yuan et al. (2008) Yuan, W., Zhou, H. Y., Komossa, S., et al. 2008, ApJ, 685, 801
  * Zhang et al. (2013) Zhang, K., Wang, T.-G., Yan, L., & Dong, X.-B. 2013, ApJ, 768, 22
  * Zhang et al. (2017a) Zhang, P., Zhang, P.-f., Yan, J.-z., Fan, Y.-z., & Liu, Q.-z. 2017a, ApJ, 849, 9
  * Zhang et al. (2017b) Zhang, S., Zhou, H., Shi, X., et al. 2017b, ApJ, 845, 126
  * Zheng et al. (2002) Zheng, X. Z., Xia, X. Y., Mao, S., Wu, H., & Deng, Z. G. 2002, AJ, 124, 18
  * Zhuang et al. (2018) Zhuang, M.-Y., Ho, L. C., & Shangguan, J. 2018, ApJ, 862, 118

##  Appendix A Tables

Table 3: Basic data of the sample.

Short name | NED Alias | RA | Dec | $z$ | Scale | Morph.
---|---|---|---|---|---|---
|  | (hh mm ss.ss) | (dd mm ss.ss) |  | (kpc/”) | type
J0347+0105 | IRAS 03450+0055 | 03 47 40.19 | 01 05 14.04 | 0.031 | 0.620 | E
J0629-0545 | IRAS 06269-0543 | 06 29 24.77 | -05 45 26.34 | 0.117 | 2.116 | I
J0632+6340 | UGC 3478 | 06 32 47.17 | 63 40 25.29 | 0.013 | 0.266 | E
J0706+3901 | FBQS J0706+3901 | 07 06 25.14 | 39 01 51.74 | 0.086 | 1.612 | I
J0713+3820 | FBQS J0713+3820 | 07 13 40.28 | 38 20 39.92 | 0.123 | 2.210 | E
J0804+3853 | FBQS J0804+3853 | 08 04 09.25 | 38 53 48.80 | 0.212 | 3.453 | E
J0806+7248 | RGB J0806+728 | 08 06 38.96 | 72 48 20.45 | 0.098 | 1.812 | I
J0814+5609 | SDSS J081432.11+560956.6 | 08 14 32.12 | 56 09 56.68 | 0.510 | 6.168 | I
J0913+3658 | RX J0913.2+3658 | 09 13 13.73 | 36 58 17.25 | 0.107 | 1.958 | I
J0925+5217 | Mrk 110 | 09 25 12.85 | 52 17 10.40 | 0.035 | 0.697 | E
J0926+1244 | Mrk 705 | 09 26 03.27 | 12 44 03.75 | 0.029 | 0.581 | I
J0937+3615 | SDSS J093703.03+361537.2 | 09 37 03.03 | 36 15 37.24 | 0.180 | 3.035 | E
J0952-0136 | Mrk 1239 | 09 52 19.09 | -01 36 43.57 | 0.020 | 0.405 | I
J1034+3938 | KUG 1031+398 | 10 34 38.59 | 39 38 28.16 | 0.042 | 0.829 | I
J1038+4227 | SDSS J103859.58+422742.3 | 10 38 59.58 | 42 27 42.22 | 0.220 | 3.552 | E
J1047+4725 | SDSS J104732.68+472532.0 | 10 47 32.67 | 47 25 32.04 | 0.799 | 7.505 | E
J1048+2222 | SDSS J104816.57+222238.9 | 10 48 16.58 | 22 22 38.99 | 0.330 | 4.752 | I
J1102+2239 | SDSS J110223.38+223920.7 | 11 02 23.38 | 22 39 20.70 | 0.453 | 5.781 | I
J1110+3653 | SDSS J111005.03+365336.3 | 11 10 05.04 | 36 53 36.27 | 0.629 | 6.830 | I
J1121+5351 | SBS 1118+541 | 11 21 08.58 | 53 51 21.03 | 0.103 | 1.893 | I
J1138+3653 | SDSS J113824.54+365327.1 | 11 38 24.54 | 36 53 27.16 | 0.356 | 4.994 | I
J1159+2838 | SDSS J115917.32+283814.5 | 11 59 17.32 | 28 38 14.58 | 0.210 | 3.427 | I
J1203+4431 | NGC 4051 | 12 03 09.61 | 44 31 52.72 | 0.002 | 0.041 | E
J1209+3217 | RX J1209.7+3217 | 12 09 45.21 | 32 17 01.13 | 0.144 | 2.527 | I
J1215+5442 | SBS 1213+549A | 12 15 49.44 | 54 42 23.97 | 0.150 | 2.614 | I
J1218+2948 | Mrk 766 | 12 18 26.52 | 29 48 46.51 | 0.013 | 0.266 | E
J1227+3214 | SDSS J122749.14+321458.9 | 12 27 49.14 | 32 14 58.96 | 0.136 | 2.408 | I
J1242+3317 | WAS 61 | 12 42 10.62 | 33 17 02.73 | 0.044 | 0.866 | E
J1246+0222 | PG 1244+026 | 12 46 35.25 | 02 22 08.72 | 0.048 | 0.941 | E
J1302+1624 | Mrk 783 | 13 02 58.84 | 16 24 27.81 | 0.067 | 1.284 | E
J1305+5116 | SDSS J130522.74+511640.2 | 13 05 22.74 | 51 16 40.15 | 0.788 | 7.469 | E
J1317+6010 | SBS 1315+604 | 13 17 50.33 | 60 10 40.70 | 0.137 | 2.423 | E
J1337+2423 | IRAS 13349+2438 | 13 37 18.71 | 24 23 03.30 | 0.108 | 1.974 | I
J1348+2622 | SDSS J134834.28+262205.9 | 13 48 34.28 | 26 22 05.95 | 0.917 | 7.832 | I
J1358+2658 | SDSS J135845.38+265808.4 | 13 58 45.38 | 26 58 08.49 | 0.331 | 4.762 | I
J1402+2159 | RX J1402.5+2159 | 14 02 34.47 | 21 59 51.65 | 0.066 | 1.266 | I
J1536+5433 | Mrk 486 | 15 36 38.40 | 54 33 33.23 | 0.039 | 0.772 | I
J1555+1911 | Mrk 291 | 15 55 07.92 | 19 11 32.82 | 0.035 | 0.697 | E
J1559+3501 | Mrk 493 | 15 59 09.62 | 35 01 47.54 | 0.031 | 0.620 | E
J1633+4718 | SDSS J163323.58+471858.9 | 16 33 23.58 | 47 18 58.96 | 0.116 | 2.101 | I
J1703+4540 | SDSS J170330.38+454047.1 | 17 03 30.39 | 45 40 47.19 | 0.060 | 1.159 | I
J1713+3523 | FBQS J1713+3523 | 17 13 04.46 | 35 23 33.55 | 0.083 | 1.561 | I
J2242+2943 | Ark 564 | 22 42 39.34 | 29 43 31.09 | 0.025 | 0.504 | E
J2314+2243 | RX J2314.9+2243 | 23 14 55.89 | 22 43 22.79 | 0.169 | 2.884 | I

777(1) Short name, (2) alias, (3), (4) right ascension and declination (J2000)
from the JVLA radio maps, (5) redshift, (6) scale, (7) radio morphology type:
I = intermediate, E = extended.

Table 4: Basic data of the sample.

Short name | log $M_{\textrm{BH}}$ | $L/L_{\textrm{Edd}}$ | log $\nu L_{\nu,\textrm{int}}$
---|---|---|---
| ($M_{\odot}$) | (erg s-1) | (erg s-1)
J0347+0105 | 7.26 | 0.07 | 39.17 $\pm$ 0.01
J0629-0545 | 7.80 | 0.11 | 40.43 $\pm$ 0.01
J0632+6340 | 6.48 | 0.02 | 37.76 $\pm$ 0.01
J0706+3901 | 7.04 | 0.04 | 39.33 $\pm$ 0.01
J0713+3820 | 8.20 | 0.17 | 39.82 $\pm$ 0.01
J0804+3853 | 7.72 | 0.14 | 39.69 $\pm$ 0.02
J0806+7248 | 6.94 | 0.08 | 40.17 $\pm$ 0.01
J0814+5609 | 8.44 | 0.11 | 42.19 $\pm$ 0.01
J0913+3658 | 7.01 | 0.06 | 38.66 $\pm$ 0.06
J0925+5217 | 7.51 | 0.10 | 38.72 $\pm$ 0.01
J0926+1244 | 7.31 | 0.04 | 38.52 $\pm$ 0.01
J0937+3615 | 7.58 | 0.05 | 39.71 $\pm$ 0.02
J0952-0136 | 7.13 | 0.05 | 39.04 $\pm$ 0.01
J1034+3938 | 6.03 | 0.17 | 39.24 $\pm$ 0.01
J1038+4227 | 7.88 | 0.08 | 40.67 $\pm$ 0.01
J1047+4725 | 8.39 | 0.08 | 43.77 $\pm$ 0.01
J1048+2222 | 7.53 | 0.11 | 39.75 $\pm$ 0.05
J1102+2239 | 8.17 | 0.06 | 40.47 $\pm$ 0.04
J1110+3653 | 7.09 | 0.23 | 41.89 $\pm$ 0.01
J1121+5351 | 7.64 | 0.13 | 39.25 $\pm$ 0.01
J1138+3653 | 7.61 | 0.09 | 41.02 $\pm$ 0.01
J1159+2838 | 7.54 | 0.09 | 39.76 $\pm$ 0.02
J1203+4431 | 6.30 | 0.01 | 36.46 $\pm$ 0.01
J1209+3217 | 7.48 | 0.10 | 39.31 $\pm$ 0.02
J1215+5442 | 7.51 | 0.10 | 39.30 $\pm$ 0.01
J1218+2948 | 6.43 | 0.05 | 38.47 $\pm$ 0.01
J1227+3214 | 6.84 | 0.16 | 39.94 $\pm$ 0.01
J1242+3317 | 7.06 | 0.06 | 38.74 $\pm$ 0.01
J1246+0222 | 7.09 | 0.05 | 38.31 $\pm$ 0.02
J1302+1624 | 7.36 | 0.12 | 39.91 $\pm$ 0.01
J1305+5116 | 8.20 | 0.67 | 42.91 $\pm$ 0.01
J1317+6010 | 6.25 | 0.15 | 39.42 $\pm$ 0.02
J1337+2423 | 8.04 | 0.22 | 40.26 $\pm$ 0.01
J1348+2622a | – | – | 40.84 $\pm$ 0.02
J1358+2658 | 7.84 | 0.12 | 40.05 $\pm$ 0.02
J1402+2159 | 7.15 | 0.11 | 38.34 $\pm$ 0.02
J1536+5433 | 7.34 | 0.09 | 37.98 $\pm$ 0.03
J1555+1911 | 6.50 | 0.03 | 37.85 $\pm$ 0.03
J1559+3501 | 6.86 | 0.06 | 38.19 $\pm$ 0.01
J1633+4718 | 6.91 | 0.11 | 40.67 $\pm$ 0.01
J1703+4540 | 7.73 | 0.01 | 40.20 $\pm$ 0.01
J1713+3523 | 6.69 | 0.06 | 39.54 $\pm$ 0.01
J2242+2943 | 6.67 | 0.08 | 38.88 $\pm$ 0.01
J2314+2243 | 8.00 | 0.17 | 40.48 $\pm$ 0.01
888(1) Short name, (2) logarithmic black hole mass estimate, (3), Eddington
ratio estimate, (4) logarithm of the 5.2 GHz integrated radio luminosity. a
The redshift of this source is too high to measure the H$\beta$ line and
derive a black hole mass estimate, or an Eddington ratio.

Table 5: Measurements from the normal radio map, and 90 and 60k$\lambda$
tapered maps, and associated rms for our sample.

Short name | rms | $S_{\textrm{peak}}$ | $S_{\textrm{int}}$ | rms${}_{90\textrm{k}\lambda}$ | $S_{\textrm{90k}\lambda,\textrm{int}}$ | rms${}_{60\textrm{k}\lambda}$ | $S_{\textrm{60k}\lambda,\textrm{int}}$
---|---|---|---|---|---|---|---
| ($\mu$Jy beam-1) | (mJy beam-1) | (mJy) | ($\mu$Jy beam-1) | (mJy) | ($\mu$Jy beam-1) | (mJy)
J0347+0105 | 14 | 7.54 $\pm$ 0.07 | 11.65 $\pm$ 0.05 | 25 | 12.10 $\pm$ 0.07 | 30 | 12.04 $\pm$ 0.08
J0629-0545 | 14 | 12.52 $\pm$ 0.06 | 14.58 $\pm$ 0.05 | 22 | 14.81 $\pm$ 0.06 | 26 | 14.82 $\pm$ 0.06
J0632+6340 | 10 | 1.69 $\pm$ 0.03 | 2.83 $\pm$ 0.04 | 14 | 3.02 $\pm$ 0.04 | 16 | 3.02 $\pm$ 0.04
J0706+3901 | 13 | 1.67 $\pm$ 0.02 | 2.14 $\pm$ 0.03 | 17 | 2.18 $\pm$ 0.04 | 21 | 2.21 $\pm$ 0.04
J0713+3820 | 13 | 2.32 $\pm$ 0.03 | 3.39 $\pm$ 0.04 | 19 | 3.71 $\pm$ 0.06 | 23 | 3.80 $\pm$ 0.06
J0804+3853 | 10 | 0.37 $\pm$ 0.02 | 0.81 $\pm$ 0.03 | 18 | 0.74 $\pm$ 0.03 | 23 | 0.72 $\pm$ 0.04
J0806+7248 | 10 | 10.92 $\pm$ 0.02 | 11.46 $\pm$ 0.03 | 11 | 11.33 $\pm$ 0.03 | 12 | 11.33 $\pm$ 0.03
J0814+5609 | 10 | 25.85 $\pm$ 0.02 | 29.27 $\pm$0.06 | 16 | 29.38 $\pm$ 0.06 | 19 | 30.09 $\pm$ 0.08
J0913+3658 | 11 | 0.27 $\pm$ 0.01 | 0.30 $\pm$ 0.02 | 17 | 0.41 $\pm$ 0.03 | 26 | 0.36 $\pm$ 0.04
J0925+5217 | 12 | 1.70 $\pm$ 0.02 | 2.25 $\pm$ 0.03 | 16 | 3.10 $\pm$ 0.06 | 21 | 3.11 $\pm$ 0.06
J0926+1244 | 11 | 2.81 $\pm$ 0.03 | 3.36 $\pm$ 0.04 | 17 | 3.46 $\pm$ 0.05 | 22 | 3.37 $\pm$ 0.05
J0937+3615 | 12 | 0.79 $\pm$ 0.01 | 1.03 $\pm$ 0.03 | 15 | 0.94 $\pm$ 0.03 | 19 | 0.90 $\pm$ 0.03
J0952-0136 | 10 | 17.49 $\pm$ 0.03 | 18.54 $\pm$ 0.04 | 12 | 18.92 $\pm$ 0.04 | 14 | 18.97 $\pm$ 0.04
J1034+3938 | 10 | 7.05 $\pm$ 0.02 | 7.37 $\pm$ 0.03 | 12 | 7.63 $\pm$ 0.04 | 14 | 7.71 $\pm$ 0.04
J1038+4227 | 10 | 2.06 $\pm$ 0.03 | 3.75 $\pm$ 0.07 | 13 | 5.36 $\pm$ 0.08 | 17 | 5.22 $\pm$ 0.07
J1047+4725 | 35 | 190.90 $\pm$ 5.70 | 381.75 $\pm$ 0.17 | 50 | 381.47 $\pm$ 0.15 | 80 | 381.34 $\pm$ 0.23
J1048+2222 | 8 | 0.27 $\pm$ 0.01 | 0.24 $\pm$ 0.01 | 9 | 0.32 $\pm$ 0.02 | 10 | 0.47 $\pm$ 0.03
J1102+2239 | 13 | 0.70 $\pm$ 0.01 | 0.70 $\pm$ 0.02 | 36 | 0.58 $\pm$ 0.05 | 45 | 1.22 $\pm$ 0.10
J1110+3653 | 10 | 8.12 $\pm$ 0.03 | 9.34 $\pm$ 0.05 | 11 | 9.82 $\pm$ 0.04 | 12 | 9.93 $\pm$ 0.04
J1121+5351 | 11 | 1.02 $\pm$ 0.02 | 1.22 $\pm$ 0.03 | 12 | 1.25 $\pm$ 0.03 | 14 | 1.24 $\pm$ 0.03
J1138+3653 | 10 | 4.37 $\pm$ 0.01 | 4.64 $\pm$ 0.02 | 12 | 4.63 $\pm$ 0.03 | 12 | 4.62 $\pm$ 0.03
J1159+2838 | 10 | 0.79 $\pm$ 0.01 | 0.84 $\pm$ 0.02 | 13 | 0.82 $\pm$ 0.02 | 14 | 0.79 $\pm$ 0.02
J1203+4431 | 12 | 0.78 $\pm$ 0.04 | 5.47 $\pm$ 0.10 | 29 | 5.82 $\pm$ 0.12 | 38 | 5.64 $\pm$ 0.13
J1209+3217 | 13 | 0.60 $\pm$ 0.03 | 0.72 $\pm$ 0.03 | 14 | 0.61 $\pm$ 0.02 | 15 | 0.63 $\pm$ 0.03
J1215+5442 | 10 | 0.56 $\pm$ 0.02 | 0.64 $\pm$ 0.02 | 12 | 0.65 $\pm$ 0.02 | 13 | 0.67 $\pm$ 0.02
J1218+2948 | 9 | 11.17 $\pm$ 0.08 | 15.46 $\pm$ 0.04 | 10 | 15.56 $\pm$ 0.03 | 12 | 15.55 $\pm$ 0.03
J1227+3214 | 11 | 2.96 $\pm$ 0.03 | 3.34 $\pm$ 0.03 | 17 | 3.63 $\pm$ 0.05 | 23 | 3.60 $\pm$ 0.06
J1242+3317 | 12 | 1.32 $\pm$ 0.03 | 2.23 $\pm$ 0.04 | 16 | 2.21 $\pm$ 0.04 | 21 | 2.09 $\pm$ 0.04
J1246+0222 | 10 | 0.43 $\pm$ 0.02 | 0.71 $\pm$ 0.03 | 11 | 0.85 $\pm$ 0.03 | 15 | 0.82 $\pm$ 0.03
J1302+1624 | 12 | 3.28 $\pm$ 0.05 | 11.47 $\pm$ 0.13 | 28 | 11.87 $\pm$ 0.13 | 35 | 11.99 $\pm$ 0.13
J1305+5116 | 14 | 32.10 $\pm$ 1.60 | 53.41 $\pm$ 0.05 | 16 | 53.37 $\pm$ 0.05 | 20 | 53.35 $\pm$ 0.06
J1317+6010 | 16 | 0.51 $\pm$ 0.03 | 0.64 $\pm$ 0.03 | 24 | 0.80 $\pm$ 0.04 | 29 | 0.79 $\pm$ 0.05
J1337+2423 | 13 | 9.64 $\pm$ 0.03 | 9.96 $\pm$ 0.04 | 14 | 9.84 $\pm$ 0.04 | 16 | 9.83 $\pm$ 0.04
J1348+2622 | 12 | 0.38 $\pm$ 0.01 | 0.38 $\pm$ 0.02 | 14 | 0.38 $\pm$ 0.02 | 15 | 0.37 $\pm$ 0.02
J1358+2658 | 11 | 0.52 $\pm$ 0.02 | 0.61 $\pm$ 0.02 | 13 | 0.62 $\pm$ 0.03 | 15 | 0.60 $\pm$ 0.03
J1402+2159 | 16 | 0.28 $\pm$ 0.01 | 0.33 $\pm$ 0.03 | 22 | 0.46 $\pm$ 0.04 | 22 | 0.28 $\pm$ 0.03
J1536+5433 | 10 | 0.44 $\pm$ 0.01 | 0.50 $\pm$ 0.02 | 18 | 0.45 $\pm$ 0.03 | 23 | 0.40 $\pm$ 0.03
J1555+1911 | 11 | 0.12 $\pm$ 0.01 | 0.36 $\pm$ 0.03 | 15 | 0.42 $\pm$ 0.03 | 17 | 0.44 $\pm$ 0.03
J1559+3501 | 11 | 0.46 $\pm$ 0.04 | 1.41 $\pm$ 0.04 | 14 | 1.47 $\pm$ 0.03 | 16 | 1.48 $\pm$ 0.04
J1633+4718 | 10 | 24.10 $\pm$ 0.03 | 25.27 $\pm$ 0.04 | 12 | 25.53 $\pm$ 0.04 | 14 | 25.59 $\pm$ 0.05
J1703+4540 | 11 | 32.26 $\pm$ 0.06 | 34.68 $\pm$ 0.04 | 17 | 35.97 $\pm$ 0.06 | 22 | 35.01 $\pm$ 0.07
J1713+3523 | 10 | 3.71 $\pm$ 0.01 | 3.81 $\pm$ 0.02 | 13 | 3.95 $\pm$ 0.03 | 15 | 3.98 $\pm$ 0.04
J2242+2943 | 11 | 5.61 $\pm$ 0.18 | 10.33 $\pm$ 0.04 | 15 | 10.43 $\pm$ 0.04 | 25 | 10.41 $\pm$ 0.06
J2314+2243 | 10 | 6.17 $\pm$ 0.02 | 7.02 $\pm$ 0.03 | 14 | 6.96 $\pm$ 0.03 | 15 | 6.99 $\pm$ 0.03

999(1) Short name, (2) map rms, (3), peak flux density, (4) integrated flux
density, (5) rms of the 90k$\lambda$ tapered map, (6) integrated flux density
of the 90k$\lambda$ tapered map, (7) rms of the 60k$\lambda$ tapered map, (8)
integrated flux density of the 60k$\lambda$ tapered map.

Table 6: Weighted total spectral indices, core spectral indices, spectral
indices of regions of interest for our sample, and the non-simultaneous 1.4 -
5.2 GHz spectral index.

Short name | $\alpha_{\textrm{total}}$ | $\alpha_{\textrm{core}}$ | $\alpha_{\textrm{region}}$ | $\alpha_{\textrm{1.4-5.2\leavevmode\nobreak\ GHz}}$
---|---|---|---|---
J0347+0105 | -0.731 $\pm$ 0.001 | -0.806 $\pm$ 0.001 |  | -0.87 $\pm$ 0.01
J0629-0545 | -0.923 $\pm$ 0.001 | -0.921 $\pm$ 0.001 |  | -0.62 $\pm$ 0.04a
J0632+6340 | -0.509 $\pm$ 0.001 | -0.626 $\pm$ 0.003 |  | -1.16 $\pm$ 0.07a
J0706+3901 | -0.807 $\pm$ 0.002 | -0.795 $\pm$ 0.007 |  | -0.71 $\pm$ 0.03
J0713+3820 | -0.990 $\pm$ 0.001 | -1.004 $\pm$ 0.004 |  | -0.88 $\pm$ 0.02
J0804+3853 | -0.979 $\pm$ 0.005 | -1.091 $\pm$ 0.020 |  | -0.91 $\pm$ 0.07
J0806+7248 | -0.961 $\pm$ 0.001 | -0.981 $\pm$ 0.001 |  | -1.12 $\pm$ 0.03a
J0814+5609 | -0.104 $\pm$ 0.001 | -0.002 $\pm$ 0.001 | -0.448 $\pm$ 0.009 | -0.77 $\pm$ 0.01
J0913+3658 | -0.767 $\pm$ 0.006 | -0.992 $\pm$ 0.036 |  | -0.78 $\pm$ 0.17
J0925+5217 | -0.341 $\pm$ 0.001 | -0.327 $\pm$ 0.002 |  | -0.98 $\pm$ 0.03
J0926+1244 | -0.592 $\pm$ 0.001 | -0.746 $\pm$ 0.004 |  | -0.71 $\pm$ 0.02
J0937+3615 | -0.765 $\pm$ 0.002 | -0.946 $\pm$ 0.010 |  | -0.87 $\pm$ 0.05
J0952-0136 | -0.983 $\pm$ 0.001 | -1.073 $\pm$ 0.002 |  | -0.89 $\pm$ 0.01
J1034+3938 | -1.170 $\pm$ 0.001 | -1.130 $\pm$ 0.002 |  | -0.96 $\pm$ 0.01
J1038+4227 | -0.893 $\pm$ 0.002 | -0.767 $\pm$ 0.004 |  | 0.33 $\pm$ 0.06
J1047+4725 | -0.747 $\pm$ 0.001 | -0.629 $\pm$ 0.000 |  | -0.53 $\pm$ 0.01
J1048+2222 | -0.559 $\pm$ 0.004 | -0.710 $\pm$ 0.015 |  | -1.39 $\pm$ 0.11
J1102+2239 | -0.761 $\pm$ 0.004 | -0.955 $\pm$ 0.016 |  | -0.72 $\pm$ 0.09
J1110+3653 | -0.182 $\pm$ 0.001 | -0.099 $\pm$ 0.001 |  | -0.61 $\pm$ 0.01
J1121+5351 | -0.612 $\pm$ 0.002 | -0.760 $\pm$ 0.008 |  | -0.43 $\pm$ 0.06
J1138+3653 | -0.588 $\pm$ 0.001 | -0.680 $\pm$ 0.001 |  | -0.76 $\pm$ 0.01
J1159+2838 | -0.711 $\pm$ 0.004 | -0.926 $\pm$ 0.015 |  | -0.66 $\pm$ 0.07
J1203+4431 | -0.616 $\pm$ 0.002 | -0.853 $\pm$ 0.008 |  | -0.96 $\pm$ 0.02
J1209+3217 | -0.548 $\pm$ 0.006 | -0.621 $\pm$ 0.016 |  | -0.76 $\pm$ 0.09
J1215+5442 | -1.012 $\pm$ 0.005 | -1.160 $\pm$ 0.017 |  | -0.99 $\pm$ 0.07
J1218+2948 | -0.830 $\pm$ 0.001 | -0.835 $\pm$ 0.001 |  | -0.73 $\pm$ 0.01
J1227+3214 | -0.671 $\pm$ 0.001 | -0.737 $\pm$ 0.002 |  | -0.50 $\pm$ 0.02
J1242+3317 | -0.739 $\pm$ 0.003 | -0.883 $\pm$ 0.007 |  | -0.79 $\pm$ 0.03
J1246+0222 | 0.199 $\pm$ 0.001 | 0.053 $\pm$ 0.004 |  | -0.87 $\pm$ 0.09
J1302+1624 | -1.500 $\pm$ 0.003 | -0.822 $\pm$ 0.003 | -0.666 $\pm$ 0.015 | -0.70 $\pm$ 0.01
J1305+5116 | -0.318 $\pm$ 0.001 | -0.072 $\pm$ 0.000 | -0.727 $\pm$ 0.001 | -0.37 $\pm$ 0.01
J1317+6010 | -0.191 $\pm$ 0.005 | -0.339 $\pm$ 0.018 |  | -0.80 $\pm$ 0.10
J1337+2423 | -1.030 $\pm$ 0.001 | -1.009 $\pm$ 0.003 |  | -0.53 $\pm$ 0.01
J1348+2622 | -1.055 $\pm$ 0.004 | -1.298 $\pm$ 0.027 |  | -1.00 $\pm$ 0.11
J1358+2658 | -0.424 $\pm$ 0.001 | -0.509 $\pm$ 0.008 |  | -0.50 $\pm$ 0.11
J1402+2159 | -0.954 $\pm$ 0.004 | -1.543 $\pm$ 0.027 |  | -0.74 $\pm$ 0.17
J1536+5433 | -0.346 $\pm$ 0.003 | -0.497 $\pm$ 0.011 |  | -0.69 $\pm$ 0.12
J1555+1911 | -0.824 $\pm$ 0.004 | -1.128 $\pm$ 0.018 |  | -1.25 $\pm$ 0.12
J1559+3501 | -1.157 $\pm$ 0.001 | -1.103 $\pm$ 0.008 |  | -0.67 $\pm$ 0.05
J1633+4718 | -0.543 $\pm$ 0.001 | -0.537 $\pm$ 0.001 | -0.694 $\pm$ 0.006 | -0.72 $\pm$ 0.01
J1703+4540 | -0.868 $\pm$ 0.001 | -0.912 $\pm$ 0.001 |  | -0.94 $\pm$ 0.01
J1713+3523 | -0.974 $\pm$ 0.001 | -1.050 $\pm$ 0.003 |  | -0.82 $\pm$ 0.01
J2242+2943 | -0.812 $\pm$ 0.001 | -0.850 $\pm$ 0.002 |  | -0.79 $\pm$ 0.03a
J2314+2243 | -0.812 $\pm$ 0.001 | -0.850 $\pm$ 0.003 |  | -0.75 $\pm$ 0.04a
101010(1) Short name, (2) weighted average spectral index of the total
emitting region, (3), weighted average spectral index of the centremost region
within a circle with a radius of 2 pixels, (4) weighted average spectral index
of a region of interest, (5) Non-simultaneous 1.4-5.2 GHz spectral index. 1.4
GHz data are from FIRST, except for the sources marked with a from NVSS. 5.2
GHz data are the intergrated flux densities in this paper.

Table 7: Parameters related to the star formation estimation for our sample.

Short name | W3 | W4 | W3-W4 | $S_{\textrm{W3}}$ |  $S_{\textrm{20cm}}$ CDFS |  $S_{\textrm{20\leavevmode\nobreak\ cm}}$ ELAIS | $q22$
---|---|---|---|---|---|---|---
| (mag) | (mag) | (mag) | (mJy) | (mJy) | (mJy) |
J0347+0105 | 5.32 $\pm$ 0.02 | 3.03 $\pm$ 0.02 | 2.29 $\pm$ 0.05 | 234.80 $\pm$ 3.46 | 21.03 $\pm$ 1.43 | 20.01 $\pm$ 2.44 | 1.21
J0629-0545 | 5.48 $\pm$ 0.02 | 2.60 $\pm$ 0.02 | 2.88 $\pm$ 0.08 | 203.38 $\pm$ 2.44 | 31.36 $\pm$ 2.11 | 29.82 $\pm$ 3.61 | 1.19
J0632+6340 | 6.28 $\pm$ 0.02 | 3.76 $\pm$ 0.02 | 2.52 $\pm$ 0.06 | 97.43 $\pm$ 1.35 | 10.74 $\pm$ 0.71 | 10.23 $\pm$ 1.23 | 1.65
J0706+3901 | 8.60 $\pm$ 0.03 | 6.05 $\pm$ 0.05 | 2.55 $\pm$ 0.14 | 11.51 $\pm$ 0.28 | 1.31 $\pm$ 0.13 | 1.25 $\pm$ 0.19 | 0.70
J0713+3820 | 6.26 $\pm$ 0.02 | 3.98 $\pm$ 0.02 | 2.28 $\pm$ 0.05 | 99.15 $\pm$ 1.37 | 8.74 $\pm$ 0.60 | 8.32 $\pm$ 1.02 | 1.19
J0804+3853 | 8.08 $\pm$ 0.02 | 5.37 $\pm$ 0.03 | 2.70 $\pm$ 0.10 | 18.62 $\pm$ 0.38 | 2.44 $\pm$ 0.19 | 2.32 $\pm$ 0.31 | 1.36
J0806+7248 | 8.11 $\pm$ 0.02 | 5.54 $\pm$ 0.04 | 2.57 $\pm$ 0.11 | 17.99 $\pm$ 0.40 | 2.08 $\pm$ 0.18 | 1.99 $\pm$ 0.28 | 0.11
J0814+5609 | 10.41 $\pm$ 0.07 | 8.43 $\pm$ 0.45 | 1.98 $\pm$ 0.66 | 2.17 $\pm$ 0.15 | 0.15 $\pm$ 0.07 | 0.15 $\pm$ 0.08 | -0.99
J0913+3658 | 8.20 $\pm$ 0.02 | 5.65 $\pm$ 0.04 | 2.55 $\pm$ 0.11 | 16.67 $\pm$ 0.32 | 1.89 $\pm$ 0.16 | 1.80 $\pm$ 0.25 | 1.67
J0925+5217 | 6.72 $\pm$ 0.02 | 4.71 $\pm$ 0.03 | 2.01 $\pm$ 0.05 | 64.67 $\pm$ 1.01 | 4.46 $\pm$ 0.34 | 4.25 $\pm$ 0.56 | 1.35
J0926+1244 | 6.16 $\pm$ 0.02 | 3.80 $\pm$ 0.02 | 2.36 $\pm$ 0.05 | 108.92 $\pm$ 1.50 | 10.40 $\pm$ 0.75 | 9.90 $\pm$ 1.25 | 1.54
J0937+3615 | 8.14 $\pm$ 0.02 | 5.47 $\pm$ 0.04 | 2.66 $\pm$ 0.12 | 17.65 $\pm$ 0.34 | 2.22 $\pm$ 0.18 | 2.12 $\pm$ 0.29 | 1.34
J0952-0136 | 4.33 $\pm$ 0.01 | 2.31 $\pm$ 0.02 | 2.02 $\pm$ 0.04 | 589.25 $\pm$ 7.60 | 40.81 $\pm$ 2.56 | 38.83 $\pm$ 4.52 | 1.16
J1034+3938 | 6.65 $\pm$ 0.02 | 4.51 $\pm$ 0.03 | 2.14 $\pm$ 0.06 | 69.49 $\pm$ 0.96 | 5.40 $\pm$ 0.40 | 5.14 $\pm$ 0.66 | 0.57
J1038+4227 | 8.87 $\pm$ 0.03 | 6.58 $\pm$ 0.06 | 2.29 $\pm$ 0.13 | 8.97 $\pm$ 0.23 | 0.80 $\pm$ 0.08 | 0.77 $\pm$ 0.12 | 0.07
J1047+4725 | 10.44 $\pm$ 0.09 | 7.75 $\pm$ 0.19 | 2.69 $\pm$ 0.52 | 2.12 $\pm$ 0.17 | 0.27 $\pm$ 0.06 | 0.27 $\pm$ 0.08 | -2.18
J1048+2222 | 9.29 $\pm$ 0.04 | 7.08 $\pm$ 0.11 | 2.20 $\pm$ 0.21 | 6.08 $\pm$ 0.25 | 0.50 $\pm$ 0.08 | 0.48 $\pm$ 0.10 | 1.09
J1102+2239 | 9.15 $\pm$ 0.03 | 6.49 $\pm$ 0.07 | 2.66 $\pm$ 0.20 | 6.96 $\pm$ 0.22 | 0.87 $\pm$ 0.10 | 0.83 $\pm$ 0.14 | 0.81
J1110+3653 | nd | nd | N/A | N/A | N/A | N/A | N/A
J1121+5351 | 7.32 $\pm$ 0.02 | 4.99 $\pm$ 0.03 | 2.33 $\pm$ 0.07 | 37.39 $\pm$ 0.62 | 3.46 $\pm$ 0.27 | 3.30 $\pm$ 0.44 | 1.48
J1138+3653 | 10.53 $\pm$ 0.08 | 8.01 $\pm$ 0.21 | 2.52 $\pm$ 0.50 | 1.94 $\pm$ 0.15 | 0.22 $\pm$ 0.05 | 0.21 $\pm$ 0.06 | -0.28
J1159+2838 | 9.02 $\pm$ 0.03 | 5.98 $\pm$ 0.04 | 3.03 $\pm$ 0.17 | 7.85 $\pm$ 0.25 | 1.39 $\pm$ 0.12 | 1.33 $\pm$ 0.19 | 1.22
J1203+4431 | 4.58 $\pm$ 0.02 | 2.16 $\pm$ 0.02 | 2.41 $\pm$ 0.05 | 468.06 $\pm$ 6.47 | 46.81 $\pm$ 2.97 | 44.53 $\pm$ 5.23 | 1.96
J1209+3217 | 8.18 $\pm$ 0.02 | 5.30 $\pm$ 0.03 | 2.88 $\pm$ 0.11 | 16.93 $\pm$ 0.33 | 2.60 $\pm$ 0.20 | 2.48 $\pm$ 0.32 | 1.75
J1215+5442 | 7.48 $\pm$ 0.02 | 5.01 $\pm$ 0.04 | 2.47 $\pm$ 0.10 | 32.17 $\pm$ 0.53 | 3.40 $\pm$ 0.28 | 3.24 $\pm$ 0.44 | 1.52
J1218+2948 | 5.00 $\pm$ 0.01 | 2.13 $\pm$ 0.02 | 2.87 $\pm$ 0.07 | 316.16 $\pm$ 4.08 | 48.08 $\pm$ 3.23 | 45.74 $\pm$ 5.54 | 1.40
J1227+3214 | 7.42 $\pm$ 0.02 | 4.80 $\pm$ 0.02 | 2.62 $\pm$ 0.06 | 34.10 $\pm$ 0.57 | 4.12 $\pm$ 0.29 | 3.92 $\pm$ 0.49 | 1.06
J1242+3317 | 6.70 $\pm$ 0.02 | 3.98 $\pm$ 0.02 | 2.72 $\pm$ 0.07 | 66.30 $\pm$ 0.98 | 8.82 $\pm$ 0.62 | 8.39 $\pm$ 1.04 | 1.59
J1246+0222 | 7.22 $\pm$ 0.02 | 4.66 $\pm$ 0.03 | 2.56 $\pm$ 0.08 | 40.95 $\pm$ 0.64 | 4.69 $\pm$ 0.34 | 4.46 $\pm$ 0.57 | 2.26
J1302+1624 | 7.88 $\pm$ 0.02 | 5.26 $\pm$ 0.04 | 2.62 $\pm$ 0.12 | 22.28 $\pm$ 0.41 | 2.70 $\pm$ 0.22 | 2.58 $\pm$ 0.35 | -0.11
J1305+5116 | 8.91 $\pm$ 0.03 | 6.57 $\pm$ 0.07 | 2.34 $\pm$ 0.16 | 8.63 $\pm$ 0.21 | 0.81 $\pm$ 0.09 | 0.78 $\pm$ 0.13 | -0.61
J1317+6010 | 8.40 $\pm$ 0.02 | 5.80 $\pm$ 0.04 | 2.60 $\pm$ 0.11 | 13.79 $\pm$ 0.27 | 1.64 $\pm$ 0.14 | 1.57 $\pm$ 0.22 | 1.60
J1337+2423 | 4.59 $\pm$ 0.01 | 2.66 $\pm$ 0.02 | 1.95 $\pm$ 0.03 | 462.06 $\pm$ 5.96 | 30.03 $\pm$ 1.91 | 28.57 $\pm$ 3.35 | 1.29
J1348+2622 | 10.40 $\pm$ 0.06 | 8.57 $\pm$ 0.29 | 1.83 $\pm$ 0.40 | 2.19 $\pm$ 0.13 | 0.13 $\pm$ 0.04 | 0.13 $\pm$ 0.05 | 0.32
J1358+2658 | 9.22 $\pm$ 0.03 | 6.53 $\pm$ 0.06 | 2.69 $\pm$ 0.18 | 6.47 $\pm$ 0.18 | 0.84 $\pm$ 0.09 | 0.80 $\pm$ 0.13 | 1.30
J1402+2159 | 7.18 $\pm$ 0.02 | 4.34 $\pm$ 0.03 | 2.75 $\pm$ 0.11 | 42.41 $\pm$ 0.63 | 5.78 $\pm$ 0.43 | 5.50 $\pm$ 0.70 | 2.16
J1536+5433 | 6.95 $\pm$ 0.02 | 5.02 $\pm$ 0.02 | 1.92 $\pm$ 0.04 | 52.76 $\pm$ 0.78 | 3.36 $\pm$ 0.23 | 3.20 $\pm$ 0.40 | 2.12
J1555+1911 | 7.97 $\pm$ 0.02 | 5.35 $\pm$ 0.04 | 2.62 $\pm$ 0.12 | 20.47 $\pm$ 0.38 | 2.49 $\pm$ 0.20 | 2.37 $\pm$ 0.32 | 1.67
J1559+3501 | 6.62 $\pm$ 0.02 | 4.28 $\pm$ 0.02 | 2.34 $\pm$ 0.05 | 71.04 $\pm$ 0.98 | 6.64 $\pm$ 0.47 | 6.33 $\pm$ 0.78 | 1.38