# The Relative Importance of Thermal Gas, Radiation, and Magnetic Pressures
Around Star-Forming Regions in Normal Galaxies and Dusty Starbursts

Eric J. Murphy<EMAIL_ADDRESS>National Radio Astronomy Observatory
520 Edgemont Road
Charlottesville, VA 22903, USA

###### Abstract

In this paper, an investigation on the relative importance of the thermal gas,
radiation, and (minimum-energy) magnetic pressures around $\approx$200 star-
forming regions in a sample of nearby normal and luminous infrared galaxies is
presented. Given the range of galaxy distances, pressure estimates are made on
spatial scales spanning $\sim$0.1$-3$ kpc. The ratio of thermal gas-to-
radiation pressures does not appear to significantly depend on star formation
rate surface density ($\Sigma_{\rm SFR}$), but exhibits a steady decrease with
increasing physical size of the aperture over which the quantities are
measured. The ratio of magnetic-to-radiation pressures appears to be
relatively flat as a function of $\Sigma_{\rm SFR}$ and similar in value for
both nuclear and extranuclear regions, but unlike the ratio of thermal gas-to-
radiation pressures, exhibits a steady increase with increasing aperture size.
Furthermore, it seems that the magnetic pressure is typically weaker than the
radiation pressure on sub-kpc scales, and only starts to play a significant
role on few-kpc scales. When the internal pressure terms are summed, their
ratio to the ($\Sigma_{\rm SFR}$-inferred) kpc-scale dynamical equilibrium
pressure estimates is roughly constant. Consequently, it appears that the
physical area of the galaxy disk, and not necessarily environment (e.g.,
nuclear vs. extranuclear regions) or star formation activity, may play the
dominant role in determining which pressure term is most active around star-
forming regions. These results are consistent with a scenario in which a
combination of processes acting primarily on different physical scales work
collectively to regulate the star formation process in galaxy disks.

Disk galaxies; Dust continuum emission; Galaxy evolution; Luminous infrared
galaxies; Magnetic fields; Stellar feedback; Galaxies; Extragalactic magnetic
fields; Radio continuum emission; Star forming regions; Star formation

††facilities: VLA, Spitzer, GALEX

##  1 Introduction

Star formation is a slow, inefficient process that appears to be largely
invariant with galaxy type or the physical scales it is averaged over (e.g.,
Kennicutt, 1998; Kennicutt et al., 2007; Krumholz & Tan, 2007; Bigiel et al.,
2008; Leroy et al., 2008; Lada et al., 2012; Utomo et al., 2018). Globally,
observations suggest that galaxies are typically able to convert only a few
percent of their available gas into stars per unit dynamical time, while
within galaxies, it is found that $\lesssim$1% of the available gas is able to
be converted into stars per unit gravitational free-fall time. To date, a
number of scenarios have been proposed to help explain the inefficiency of
star formation, though a dominant process has yet to clearly emerge.

Magnetic fields, cosmic-rays and/or externally-driven turbulence that can
pervade the entire disks of galaxies (e.g., Parker, 1969; Sellwood & Balbus,
1999; Ostriker et al., 2001) have all been proposed as a means to prohibit gas
from collapsing into newly-formed stars. The dynamical importance of magnetic
fields and cosmic rays have also been demonstrated observationally by showing
that their energy densities can be comparable to the thermal and turbulent gas
energy densities on large scales (e.g., Beck, 2007; Tabatabaei et al., 2008).
However, stellar feedback processes that can operate on the scales of
individual giant molecular clouds almost certainly play a substantial role
driving the inefficiency of star formation (e.g. Krumholz et al., 2014).

Massive ($\gtrsim 8\,M_{\odot}$) stars are able to inject large amounts of
energy and momentum into the interstellar medium (ISM) of galaxies through
numerous mechanisms. Models for feedback in the ISM have largely included
energy and momentum injection by supernova, stellar winds, and radiation
pressure on dust (McKee & Ostriker, 1977; Silk, 1997; Scoville, 2003; Thompson
et al., 2005; Draine, 2011; Murray et al., 2010). Furthermore, observations
have additionally provided support for thermal gas and radiation pressure on
dust being potentially important regulators of star formation both for
individual massive star clusters (e.g., Scoville et al., 2001; Murray, 2011),
as well as for entire starbust galaxies (e.g., Andrews & Thompson, 2011).

In this paper, a combination of mid-infrared and multi-frequency radio data
are used to investigate the relative strengths of thermal gas, radiation, and
magnetic pressures on $\sim$100 pc to $\sim$3 kpc scales around star-forming
regions within a sample of nearby galaxies included in the Star Formation in
Radio Survey (SFRS; Murphy et al., 2012, 2018) and luminous infrared galaxies
(LIRGs; $L_{\rm IR}>10^{11}\,L_{\sun}$) in the Great Observatories All-Sky
LIRG Survey (GOALS; Armus et al., 2009). The paper is organized as follows:
The data and analysis procedures are presented in Section 2. The results are
then presented in Section 3. A brief discussion of the findings and main
conclusions are given in Section 4.

##  2 Data and Analysis

To investigate the relative roles played by thermal gas, radiation and
magnetic pressures within galaxy disks, a combination of mid-infrared and
multi-frequency radio data are used for a sample of relatively nearby
galaxies. The SFRS galaxy sample is derived from the Spitzer Infrared Nearby
Galaxies Survey (SINGS; Kennicutt et al., 2003) and Key Insights on Nearby
Galaxies: a Far-Infrared Survey with Herschel (KINGFISH; Kennicutt et al.,
2011) legacy programs. Consequently, all sources are well studied and have a
wealth of ancillary data available. Radio data were collected for a total of
112 star-forming complexes (50 nuclei and 62 extranuclear regions) easily
observable with the VLA (i.e., having $\delta>-35\arcdeg$).

The GOALS galaxy sample includes over 200 LIRGs included in the flux density-
limited ($S_{\rm 60\,\mu m}>5.24$ Jy) IRAS Revised Bright Galaxy Sample (RBGS;
Sanders et al., 2003). A subset of 68 GOALS galaxies within the declination
range spanning $-20\arcdeg<\delta<20\arcdeg$ were observed with the VLA as
part of the GOALS “equatorial” survey and included in the present study.

###  2.1 SFRS Data

Radio data were obtained as part of multiple campaigns using the VLA covering
the S- ($2-4$ GHz), Ku- ($12-18$ GHz), and Ka- ($26.5-40$ GHz) bands. Details
on the data reduction and imaging procedures can be found in Murphy et al.
(2018); Linden et al. (2020). To summarize, standard calibration procedures
are followed, using the VLA calibration pipeline built on the Common Astronomy
Software Applications (CASA; McMullin et al., 2007) versions 4.6.0 and 4.7.0.
After each initial pipeline run, the calibration tables and visibilities are
manually inspected and flagged for signs of instrumental problems (e.g., bad
deformatters) and radio frequency interference (RFI). After flagging, the
pipeline is re-run, with this process repeated until any further signs of bad
data could not be detected. As with the data calibration, a detailed
description of the imaging procedure used here can be found in Murphy et al.
(2018); Linden et al. (2020). Calibrated VLA measurement sets for each source
were then imaged using the task tclean in CASA version 4.7.0.

We also make use of archival infrared and UV data. Spitzer 24 $\mu$m data were
taken from the SINGS and Local Volume Legacy (LVL) legacy programs. Details on
the associated observation strategies and data reduction steps can be found in
Dale et al. (2007) and Dale et al. (2009), respectively. Two galaxies, IC 342
and NGC 2146, were not a part of SINGS or LVL; their 24 $\mu$m imaging comes
from Engelbracht et al. (2008). GALEX far-UV (FUV; 1528Å) data were taken from
the GALEX Large Galaxy Atlas (Seibert, 2007).

###  2.2 GOALS Data

Similar to the SFRS observations, radio data for the GOALS equatorial sample
were obtained using the VLA covering the S- ($2-4$ GHz), Ku- ($12-18$ GHz),
and Ka- ($26.5-40$ GHz) bands. Calibration and imaging procedures are
identical to those described above for the SFRS radio data. More details on
the data reduction and imaging procedures, which follows those described in
Murphy et al. (2018), can be found in Linden et al. (2019).

Archival Spitzer 8 $\mu$m data taken as part of GOALS are also used in this
analysis as the basis for total infrared luminosity ($L_{\rm 8-1000\,\mu m}$)
estimates of each region. Details on the associated observation strategies and
data reduction steps are available in J. M. Mazzarella et al. (2022, in
preparation). The conversion between the 8 $\mu$m and total infrared
luminosities rely on the study of Elbaz et al. (2011) and assume a fixed ratio
${\rm IR8}=L_{\rm IR}/L_{\rm 8\mu m}=8.1\pm 2$ (see Linden et al., 2019, for
more details).

Figure 1: Left: Minimum-energy magnetic field strength distribution for all
sources, along with the distributions for nuclear and extranuclear regions.
The field strengths for the extranuclear regions are, on average, smaller and
more narrowly distributed than those for the nuclear regions. Right: Minimum-
energy magnetic field strength distributions for regions within the SFRS and
GOALS galaxies. The field strengths of regions within the SFRS galaxies are,
on average, smaller and more narrowly distributed than those within the GOALS
sources.

Figure 2: Left: Radiation pressure distribution for all sources, along with
the distributions for nuclear and extranuclear regions. Similar to the
minimum-energy magnetic field strength distribution, the radiation pressure
estimates for the extranuclear regions are, on average, smaller and more
narrowly distributed than those for the nuclear regions. Right: Radiation
pressure distributions for regions within the SFRS and GOALS galaxies. The
radiation pressures for regions within the SFRS galaxies are, on average,
smaller than those within the GOALS sources. The vertical lines in both
figures indicate the minimum and maximum radiation pressure estimates averaged
over a sample of Hii regions measured on ${\sim}50{-}100$ pc scales by Barnes
et al. (2021), indicating that our measurements largely fall between their
extreme average values.

Figure 3: Left: Thermal gas pressure distribution for all sources, along with
the distributions for nuclear and extranuclear regions. Unlike what is seen
for the radiation pressure distributions, the thermal gas pressure estimates
for the nuclear and extranculear regions are, on average, quite similar.
Right: Thermal gas pressure distributions for regions within the SFRS and
GOALS galaxies. The thermal gas pressure distributions for regions within the
SFRS and GOALS galaxies are quite similar. The vertical lines in both figures
indicate the minimum and maximum thermal gas pressure estimates averaged over
a sample of Hii regions measured on ${\sim}50{-}100$ pc scales by (Barnes et
al., 2021), indicating that our measurements largely fall between their
extreme average values.

###  2.3 Photometry

Details on the SFRS photometry can be found in Linden et al. (2020). To
summarize, photometry on the full VLA and Spitzer/MIPS 24 $\mu$m data sets
were carried out after matching their resolution (7″ FWHM), cropping each
image to a common field-of-view, and regridding to a common pixel scale.
Apertures having a diameter of 7″ were used to measure the emission from each
region within the galaxies. Given the distance of the SFRS galaxies, this
angular diameter corresponds to linear values ranging from $\sim$100$-1000$
pc, with a median of $\sim$260 pc. Flux densities and associated uncertainties
for all star-forming regions used in the present analysis are given in Linden
et al. (2020). Sources identified as likely background galaxies, possibly
being associated with supernovae, or potential anomalous microwave emission
candidates in Linden et al. (2020), are not included in this study.

Details on the GOALS photometry can be found in Linden et al. (2019). Similar
to the handling of the SFRS data, photometry on the GOALS radio and 8 $\mu$m
imaging was carried out after first matching their resolution (2$\farcs$5
FWHM), cropping each image to a common field-of-view, and regridding to a
common pixel scale. Apertures having a diameter of 4″ were used to measure the
emission from each region with the galaxies. Given the distance of the GOALS
galaxies, this angular diameter corresponds to linear values ranging from
$\sim$1$-2.8$ kpc, with a median of $\sim$1.7 kpc. Flux densities and
associated uncertainties for all regions are given in Linden et al. (2019).

Given that the photometry comes from two catalogs having used photometric
apertures of differing sizes relative to the angular resolution of the maps,
aperture corrections are applied assuming that the source emission profiles
are largely Gaussian. The SFRS photometry is therefore multiplied by a factor
of 1.72 and the GOALS photometry is multiplied by a factor of 1.13. The final
sample includes a total of 204 sources: 31 nuclear (i.e., those at a
galactocentric radius $r_{\rm g}<250$ pc) and 107 extranuclear star forming
regions from the SFRS sample, along with 22 nuclear and 44 extranuclear star-
forming regions from the GOALS sample.

###  2.4 Separation of Non-Thermal and Free-Free Emission

Using the multi-frequency radio data, the non-thermal and free-free emission
is separated to independently estimate the minimum-energy magnetic field
strengths and star formation rates, as described below. The method used to
separate the non-thermal and free-free components is described in detail in
Murphy et al. (2018); Linden et al. (2019, 2020). However, rather than
measuring the relative fractions at 33 GHz, measurements at 3 GHz are used as
this frequency is more sensitive to the amount of non-thermal emission
necessary to estimate the minimum-energy magnetic field strengths.

Briefly, the fraction of free-free emission at 3 GHz was estimated using he
prescription in Klein et al. (1984), such that

$f_{\rm
T}^{\nu_{1}}=\frac{(\frac{\nu_{2}}{\nu_{1}})^{\alpha}-(\frac{\nu_{2}}{\nu_{1}})^{\alpha^{\rm
NT}}}{(\frac{\nu_{2}}{\nu_{1}})^{-0.1}-(\frac{\nu_{2}}{\nu_{1}})^{\alpha^{\rm
NT}}}$ (1)

where $\nu_{1}$ is the target frequency (3 GHz), $\alpha$ is the observed
spectral index between 3 and 33 GHz, $\alpha^{\rm NT}$ is the non-thermal
spectral index, and a single power-law exponent for the free-free emission
($\sim-0.1$) is assumed. The non-thermal spectral index is fixed at
$\alpha^{\rm NT}=-0.83$ except when $\alpha<-0.83$ or $\alpha_{3-15\,{\rm
GHz}}<-0.83$, in which case the lesser (steeper) values are used to set the
lower limit on $\alpha^{\rm NT}$. A constant non-thermal radio spectral index
of $-0.83$ is assumed based on the average non-thermal spectral index found
among the 10 star-forming regions studied in NGC 6946 by Murphy et al. (2011).
This value is additionally consistent with the results of Niklas & Beck (1997,
i.e., $\alpha^{\rm NT}=-0.83$ with a scatter of $\sigma_{\alpha^{\rm
NT}}=0.13$) for a sample of 74 nearby galaxies, although quite a bit steeper
than the mean value of $-0.6$ measured by Klein et al. (2018) for a sample of
14 galaxies when employing more sophisticated models for radio spectra than
simple power laws.

While the calculated thermal fractions are clearly sensitive to the non-
thermal spectral index, Linden et al. (2020) showed via a Monte Carlo Markov
Chain (MCMC) analysis that adopting a fixed value does not introduce any
systematic biases over a reasonable set of physically-motivated input
conditions for a simple two-component power-law model for radio spectra. In
order to test the sensitivity of the assumed non-thermal spectral index of
$-0.83$ on the present analysis, the significantly flatter value of
$\alpha_{\rm NT}=-0.6$ was instead assumed for all sources. This resulted in a
median 3 GHz thermal fraction of 0.37, which is only 20% lower than the value
when assuming $\alpha_{\rm NT}=-0.83$ (i.e., 0.47), and thus would not impact
the main conclusions of this paper.

###  2.5 Magnetic Pressure Estimates

Taking the non-thermal flux densities, the revised equipartition and minimum-
energy formulas of Beck & Krause (2005) are used to calculate the minimum-
energy magnetic field strength for each region. This calculation assumes a
proton-to-electron number density ratio for particles in the energy range
corresponding to GHz synchrotron emission of 100, a fixed non-thermal spectral
index (as assumed above for the estimate of the thermal fractions), and a path
length through the emitting medium equal to 1 kpc. The distribution of
minimum-energy magnetic fields for all regions are shown in Figure 1.

The values used for both the non-thermal spectral index and proton-to-electron
number density ratio do not strongly affect the resulting minimum-energy
magnetic field strength estimates. For example, choosing the rather extreme
value of $\alpha_{\rm NT}=-0.6$ for all regions only increases the median
minimum-energy magnetic field strength by $\approx$10% (i.e., from 21.4 to
23.6 $\mu$G), while doubling the proton-to-electron number density ratio to
200 for all regions, in the case of extreme energy losses to CR electrons,
only increases the median minimum-energy magnetic field strength by
$\approx$20% (i.e., from 21.4 to 25.6 $\mu$G). However, even for extreme
starbursts where where hadronic losses may dominate, the assumption of a
proton-to-electron number density ratio of $\sim$100 and the usage of the
equipartition and minimum-energy formulas are expected to be valid (Lacki &
Beck, 2013).

The choice of 1 kpc for the path length should be appropriate for most
regions, given that this is the value for the vertical synchrotron scaleheight
in typical galaxy disks (e.g., Krause et al., 2018). This is also likely
appropriate for nuclear regions that may drive outflows on similar scales,
though such path lengths could in fact be larger. Where this assumption may
not be valid is for Hii regions resolved on $\sim$100 pc scales, where shocks
are able to drive enhancements to the local magnetic field, resulting in
synchrotron emission in excess of the local medium. In such cases, smaller
values of the path length may be more appropriate, which would increase the
minimum-energy magnetic field strength estimates. However, as shown later in
§3.1, scaling the path length with the physical diameter for those
extranuclear star-forming regions resolved at $<$1 kpc scales after removing a
local background does not qualitatively affect the results.

For a relativistic gas the magnetic pressure, $P_{\rm B}$, is related to the
magnetic field energy density such that

$P_{\rm B}=u_{\rm B}=\frac{B^{2}}{8\pi},$ (2)

where, in cgs units, $P_{\rm B}$ is in units of ${\rm dyn\,cm^{-2}}$, $u_{\rm
B}$ is in units of ${\rm erg\,cm^{-3}}$, and $B$ is in units of Gauss.

Under the same minimum-energy assumptions, the cosmic-ray energy density is
related to the magnetic field energy density, such that

$u_{\rm CR}=\left(\frac{2}{1-\alpha^{\rm NT}}\right)u_{\rm B}$ (3)

(Beck & Krause, 2005), and the cosmic-ray pressure of the relativistic gas is
related to the cosmic-ray energy density such that

$P_{\rm CR}=\frac{u_{\rm CR}}{3}.$ (4)

Assuming that the true magnetic field strength is near the minimum-energy
value, the cosmic-ray pressure will be sub-dominant relative to the magnetic
pressure. If this were not the case, Parker instabilities would occur,
inflating bubbles that expel cosmic rays from the galaxy disks (Parker, 1965).

###  2.6 Radiation Pressure Estimates

To get an estimate for the radiation pressure of each star-forming region, an
estimate of the bolometric surface brightness is required. For the SFRS
galaxies, the bolometric surface brightness is estimated from each region by
first calculating observed (i.e., non-extinction corrected) far-UV and
infrared based star formation rates using the $GALEX$ FUV and Spitzer 24
$\mu$m photometry along with Equations 2 and 14 in Murphy et al. (2012),
respectively, where

$\Biggl{(}\frac{\rm\psi_{FUV}}{M_{\sun}\,{\rm yr^{-1}}}\Biggr{)}=4.42\times
10^{-44}\Biggl{(}\frac{L_{\rm FUV}}{\rm erg~{}s^{-1}}\Biggr{)}$ (5)

and

$\Biggl{(}\frac{\rm\psi_{24\,\micron}}{M_{\sun}\,{\rm
yr^{-1}}}\Biggr{)}=2.45\times 10^{-43}\Biggl{[}\frac{\nu
L_{\nu}(24\,\micron)}{\rm erg~{}s^{-1}}\Biggr{]}.$ (6)

These star formation rates are then summed and converted back into an estimate
of the bolometric surface brightness by first estimating the bolometric flux,
$F_{\rm bol}$, using Equation 15 in Murphy et al. (2012) and then dividing by
the area of the region, such that

$I_{\rm bol}=\frac{F_{\rm bol}}{\Omega}\sim\Biggl{(}\frac{2.65\times
10^{-7}}{\Omega d_{\rm L}^{2}}\Biggr{)}\Biggl{(}\frac{\psi_{\rm
FUV}+\psi_{24\,\mu{\rm m}}}{M_{\sun}~{}{\rm yr^{-1}}}\Biggr{)},$ (7)

where $\Omega$ is in units of steradians, $d_{\rm L}$ is the luminosity
distance to the source in Mpc, and $I_{\rm bol}$ is in cgs units of ${\rm
erg\,s^{-1}\,cm^{-2}\,sr^{-1}}$. For the GOALS galaxies, the Spitzer 8
$\mu$m-based total infrared luminosities given in Linden et al. (2019) are
used as an estimate for the bolometric fluxes for these extremely dusty
systems. With an estimate of the bolometric surface brightness for each
region, the radiation pressure, $P_{\rm rad}$ is related to the radiation
energy density for radiation emitted near the surface of a semitransparent
body such that

$P_{\rm rad}=u_{\rm rad}\approx\left(\frac{2\pi}{c}\right)I_{\rm bol}.$ (8)

The distribution of radiation pressures for all regions are show in Figure 2.

###  2.7 Thermal Gas Pressure Estimates

At radio frequencies, where $\tau\ll 1$, the ionizing photon rate is directly
proportional to the free-free spectral luminosity, $L_{\nu}^{\rm ff}$, varying
only weakly with electron temperature $T_{\rm e}$ (Rubin, 1968), such that

$\begin{split}\biggl{[}\frac{Q(H^{0})}{\rm s^{-1}}\Biggr{]}&=6.3\times
10^{25}\\\ &\Biggl{(}\frac{T_{\rm e}}{10^{4}~{}{\rm
K}}\Biggr{)}^{-0.45}\Biggl{(}\frac{\nu}{\rm
GHz}\Biggr{)}^{0.1}\Biggl{(}\frac{L_{\nu}^{\rm T}}{\rm
erg~{}s^{-1}~{}Hz^{-1}}\Biggr{)},\end{split}$ (9)

where an electron temperature of $T_{\rm e}=10^{4}$ K is assumed.

Assuming that each region being investigated can be roughly approximated by an
ionization bounded Hii region in equilibrium, the ionizing photon rates can be
used to estimate the corresponding electron densities such that,

$n_{\rm e}\gtrsim\sqrt{\frac{3Q(H^{0})}{4\pi\alpha_{\rm H}R_{\rm s}^{3}}},$
(10)

where, in cgs units, $n_{\rm e}$ is in units of ${\rm cm^{-3}}$, $R_{s}$ is
Strömgren radius in units of cm, taken as half of the photometric aperture
diameter, and $\alpha_{\rm H}\approx 3\times 10^{-13}\,{\rm cm^{3}\,s^{-1}}$
is the case B recombination coefficient for hydrogen. The $\gtrsim$ sign in
the above equation comes from the fact that the photometric apertures are
almost certainly upper limits to the actual Hii region sizes.

With an estimate for the electron density, the thermal gas pressure can also
be estimated such that,

$P_{\rm th}=2\times 10^{4}n_{\rm e}k_{\rm B}\Biggl{(}\frac{T_{\rm e}}{\rm
10^{4}\,K}\Biggl{)}$ (11)

where, in cgs units, $P_{\rm th}$ is in units of ${\rm dyn\,cm^{-2}}$ and
$k_{\rm B}=1.38\times 10^{-16}\,{\rm cm^{2}\,g\,s^{-2}\,K^{-1}}$ is
Boltzmann’s constant. Again, an electron temperature of $T_{\rm e}=10^{4}$ K
is assumed. The factor of 2 is included under the additional assumption that
hydrogen is fully ionized and helium is singly ionized, such that $n_{\rm
e}+n_{\rm H}+n_{\rm He}\approx 2n_{\rm e}$. The distribution of thermal gas
pressures for all regions are shown in Figure 3.

###  2.8 Star Formation Rate Estimates

Being directly proportional to the ionizing photon rate, optically-thin free-
free emission can additionally be used to estimate star formation rates of
each region using Equation 11 in Murphy et al. (2011) such that

$\begin{split}\Biggl{(}\frac{\psi_{\rm
ff}}{M_{\odot}\mathrm{\,yr}^{-1}}\Biggr{)}&=4.6\times 10^{-28}\\\
&\Biggl{(}\frac{T_{\rm e}}{10^{4}~{}{\rm
K}}\Biggr{)}^{-0.45}\Biggl{(}\frac{\nu}{\rm
GHz}\Biggr{)}^{0.1}\Biggl{(}\frac{L_{\nu}^{\rm
ff}}{\mathrm{erg\,s^{-1}\,Hz}^{-1}}\Biggr{)}.\end{split}$ (12)

Star formation rates are converted into star formation rate surface densities
($\Sigma_{\rm SFR}$) by dividing SFRs by the physical area of each aperture.
By using the free-free emission, this estimate of the star formation rate is
likely more robust than the infrared-based estimates as it does not depend on
the fraction of photons that are reprocessed by dust relative to those that
are able to escape each star-forming region or the fraction of dust that is
heated by older stellar populations. Even so, the free-free to infrared-based
star formation rates have a median ratio of unity with a median absolute
deviation of $\approx$0.65. It is likely worth noting that if a significantly
flatter value for the non-thermal spectral index (i.e., $-0.6$) was assumed
when carrying out the thermal/non-thermal decomposition of the radio emission
as described in §2.4, the free-free and infrared based star formation rates
become increasingly discrepant, decreasing from a median ratio of unity to a
value of 0.89.

Figure 4: Logarithmic thermal gas-to-radiation pressure ratios plotted as a
function of star formation rate surface density. Subsets of the sample are
identified using different plotting symbols. The dashed line indicates a ratio
of unity. The thin solid line illustrates the expected trend given the
correlated nature of the axes, as they both rely on the photometric aperture
diameters. The two sources having thermal gas-to-radiation pressure ratios
$\gtrsim$5 are labeled: the nucleus of NGC 4594 hosts an AGN (Moustakas et
al., 2010); NGC 2403 Enuc. 1 A is likely an optically-faint SNR (Turner & Ho,
1994; Schlegel & Pannuti, 2003). A least-squares fit to the data is shown by a
dot-dashed line (along with 1$\sigma$ uncertainties by the dotted-lines)
indicating that the decrease in thermal gas-to-radiation pressure ratios with
increasing values of $\Sigma_{\rm SFR}$ is only marginally steeper than what
is expected given the correlated nature of the axes.  Figure 5: Logarithmic
thermal gas-to-radiation pressure ratio plotted as a function of aperture
diameter in physical units. Subsets of the sample are identified using
different plotting symbols: SFRS nuclear regions (open diamonds); SFRS
extranuclear regions (filled circles); GOALS nuclear regions (open triangles);
and GOALS extranuclear regions (open squares). Running medians for all
regions, as well as all nuclear and extranuclear regions, are shown by solid
lines using large filled circles, triangles, and squares, respectively. The
thermal gas-to-radiation pressure ratios appear to decrease with increasing
physical size of each region, indicating that the thermal pressure (i.e.,
free-free emission) is more centrally peaked on individual star-forming
regions than the radiation pressure (i.e., infrared emission).  Figure 6:
Logarithmic magnetic-to-radiation pressure ratios plotted as a function of
star formation rate surface density. Subsets of the sample are identified
using different plotting symbols. The dashed line indicates a ratio of unity.
The two sources having magnetic-to-radiation pressure ratios $\gtrsim$5 are
labeled: the nucleus of NGC 5866 hosts a known AGN (Moustakas et al., 2010);
NGC 5194 Enuc. 11 C may be an optically-faint SNR (Maddox et al., 2007). The
distribution of magnetic-to-radiation pressure ratios appears relatively flat
as a function of star formation rate surface density.

##  3 Results

In the following section, the results from the pressure estimates described
above (i.e., thermal gas, radiation, and magnetic/CR) are presented. This is
first done by comparing each of these locally-measured internal pressure terms
with one another, followed by looking at how their total sum compares to the
required pressure for the interstellar gas to stay in equilibrium within the
gravitational potential of a galaxy on kpc scales. Estimated pressures and
other physical quantities for each region are provided in Table 1 of the
Appendix.

###  3.1 Individual Pressure Terms

In Figure 4, thermal gas-to-radiation pressure ratios are plotted as a
function of $\Sigma_{\rm SFR}$. Extranuclear and nuclear regions from the SFRS
and GOALS galaxies are identified. There are two sources with thermal gas-to-
radiation pressure ratios $\gtrsim$5\. One is the nucleus of NGC 4594, which
is known to host an active galactic nuclei (AGN) (Moustakas et al., 2010) that
is most likely driving this large ratio due to a significant amount of radio
emission associated with accretion relative to the thermal dust emission. The
second is NGC 2403 Enuc. 1 a, which has a counterpart in the Open Supernova
Catalog (OSC; Guillochon et al., 2017), but was not categorized by Linden et
al. (2020) as a potential supernova remnant (SNR) due to the lack of an
optical counterpart. NGC 2403 Enuc. 1 A was previously identified as a
possible radio supernova candidate by Turner & Ho (TH-2; 1994) and was also
detected in X-rays by Schlegel & Pannuti (SP-12; 2003), thus indeed appearing
consistent with being an optically-faint SNR.

From Figure 4, the radiation pressure is found to be a factor of $\sim$2.2
larger, on average when taking a median, compared to the thermal gas pressure
among all regions. Also, while there appears to be a clear trend of decreasing
thermal gas-to-radiation pressure ratios with increasing values of
$\Sigma_{\rm SFR}$, a least-squares fit to the data (dot-dashed line)
indicates that the decrease is only marginally steeper than what is expected
given the correlated nature of the axes (i.e., thin solid line). This result
suggests that the radiation pressure (i.e., infrared dust emission) does not
increase that much more rapidly with star formation activity than the thermal
gas pressure (i.e., free-free emission).

In Figure 5, thermal gas-to-radiation pressure ratios are plotted as a
function of aperture diameter in linear units. Various subsets of the sample
are identified using different plotting symbols. Additionally, running medians
for all regions, as well as all nuclear and extranuclear regions, are shown by
solid lines using large filled circles, triangles, and squares, respectively.
What is clearly shown is that the ratio of thermal gas-to-radiation pressures
decrease with increasing physical size of each region, likely indicating that
the thermal pressure (i.e., free-free emission) is more centrally peaked on
individual star-forming regions than the radiation pressure (i.e., infrared
dust emission).

In Figure 6, magnetic-to-radiation pressure ratios are instead plotted as a
function of $\Sigma_{\rm SFR}$. Again, extranuclear and nuclear regions from
the SFRS and GOALS galaxies are identified. Overall, the distribution of
magnetic-to-radiation pressure ratios appears to be relatively flat as a
function of $\Sigma_{\rm SFR}$. Given that the magnetic pressure estimates
were made under the minimum energy assumption, any comparison with the CR and
radiation pressures will yield similar results, and therefore are not shown.

There are two sources that have magnetic-to-radiation pressure ratios
$\gtrsim$5: NGC 5866 and NGC 5194 Enuc. 11 C. NGC 5866 is an S0 galaxy that
hosts a known AGN (Moustakas et al., 2010), which is most likely driving this
large ratio. NGC 5194 Enuc. 11 C has a counterpart in the OSC (Guillochon et
al., 2017), but was not categorized by Linden et al. (2020) as a potential SNR
due to the lack of an optical or X-ray counterpart. However, as pointed out by
Maddox et al. (2007), this source may in fact be an optically-faint SNR.

In Figure 7, a box and whiskers plot is shown indicating the minimum, lower
quartile, median, upper quartile, and maximum values for the ratio of
magnetic-to-radiation pressures for the SFRS and GOALS nuclear and
extranuclear regions. This is done to see if there is any significant
differences in the pressure ratios among the galaxy samples and nuclear and
extranuclear environments.

Figure 7: A box and whiskers plot illustrating the minimum, lower quartile,
median, upper quartile, and maximum values for the logarithmic ratio of
magnetic-to-radiation pressures for the SFRS and GOALS nuclear and
extranuclear regions. The median ratios are quite similar for the nuclear
regions in both the SFRS and GOALS galaxies, however, the ratio for the
extranuclear regions in the GOALS galaxies is a factor of $\sim$2 larger than
that for the SFRS extranuclear regions.  Figure 8: Logarithmic magnetic-to-
radiation pressure ratio plotted as a function of aperture diameter in
physical units. Subsets of the sample are identified using different plotting
symbols: SFRS nuclear regions (open diamonds); SFRS extranuclear regions
(filled circles); GOALS nuclear regions (open triangles); and GOALS
extranuclear regions (open squares). Running medians for all regions, as well
as all nuclear and extranuclear regions, are shown by solid lines using large
filled circles, triangles, and squares, respectively. The red dashed-line and
open symbols indicate how the magnetic pressure estimates change by letting
the path length of the synchrotron emitting medium scale with the physical
diameter of the photometric apertures for extranuclear regions resolved on
scales $<$1 kpc after subtracting a local background. The magnetic-to-
radiation pressure ratios appear to be sensitive to the physical size of each
region, showing a steady increase with increasing physical size of the
aperture.

The median ratios appear to be quite similar for the nuclear regions in both
the SFRS and GOALS galaxies. However, there does appear to be a measurable
difference in the pressure ratios for the extranculear regions. The magnetic-
to-radiation pressure ratio for the extragalactic regions in the GOALS
galaxies is a factor of $\sim$2 larger than that for the SFRS extranuclear
regions.

Given that the GOALS galaxies are at much larger distances than the SFRS
galaxy sample, the physical sizes of the regions being studied in the GOALS
galaxies are larger by a factor of $\sim$6.5, on average (i.e., a median
aperture diameter of 1.7 kpc for the GOALS galaxies compared to a median
aperture diameter of 260 pc for the SFRS galaxies). To see how this affects
the magnetic-to-radiation pressure ratios, they are plotted as a function of
aperture diameter in linear units in Figure 8.

In Figure 8, subsets of the sample are identified using different plotting
symbols, which include the SFRS nuclear regions, SFRS extranuclear regions,
GOALS nuclear regions, and GOALS extranuclear regions. Running medians for all
regions, along with the subset of nuclear and extranuclear regions, are also
shown by solid lines indicating an increase in the magnetic-to-radiation
pressure ratios with increasing physical size of the apertures. Red dashed-
lines indicate how the the running medians change when using magnetic pressure
estimates that let the path length of the synchrotron emitting medium scale
with the physical diameter of the photometric apertures for extranuclear
regions resolved on scales $<$1 kpc after subtracting a local background (see
§2.5). As is clearly shown, such a change does not significantly change the
magnetic-to-radiation pressure ratios for such regions. Consequently, the
magnetic-to-radiation pressure ratios appear to be sensitive to the physical
size of each region, showing a steady increase with increasing physical size
of the aperture.

Figure 9: Logarithmic ratios of the total internal-to-dynamical equilibrium
pressures plotted as a function of aperture diameter in physical units. The
dynamical equilibrium pressures are estimated via a linear interpolation using
the measured star formation rate surface densities and the correlation between
$\Sigma_{\rm SFR}$ and $P_{\rm DE}$ found by Sun et al. (2020). Subsets of the
sample are identified using different plotting symbols. The dashed line
indicates a ratio of unity, while the solid line illustrated the median ratio.
Neglecting any additional pressure term associated with turbulent motions of
the gas, only $\approx$5% of the star-forming regions appear to be over-
pressurized relative to the dynamical equilibrium pressure, each of which are
measured on $\lesssim$1 kpc scales.

###  3.2 Total Internal Pressure vs. Dynamical Equilibrium Pressure

Having estimates for a number of internal pressure terms, their net effect can
be used to determine what fraction of sources appear to be over-pressurized
relative to what is required for the interstellar gas to stay in equilibrium
within the gravitational potential of a galaxy. Unfortunately, the current
analysis lacks an independent estimate of the dynamical equilibrium pressure
to actually determine what fraction of sources may be over-pressurized.

The dynamical equilibrium pressure typically includes a term to account for
the weight associated with self-gravity of the ISM disk (Spitzer, 1942;
Elmegreen, 1989) and another term to account for the weight of the ISM due to
stellar gravity in the potential well of the stars (Spitzer, 1942; Blitz &
Rosolowsky, 2004). Following this approach, the dynamical equilibrium pressure
has recently been expressed by Sun et al. (2020) such that

$P_{\rm DE}=\frac{\pi G}{2}\Sigma_{\rm gas}^{2}+\Sigma_{\rm
gas}\sqrt{2G\rho_{*}}\sigma_{\rm gas,z},$ (13)

where, $\Sigma_{\rm gas}=\Sigma_{\rm mol}+\Sigma_{\rm atom}$ is the total gas
surface density, $\rho_{*}$ is the stellar mass volume density near the disk
mid-plane, and $\sigma_{\rm gas,z}$ is the velocity dispersion of the gas
perpendicular to the disk. Since such independent estimates of $P_{\rm DE}$
are not available for the regions being studied here, we assign values based
on the work of Sun et al. (2020), who present kpc-scale dynamical equilibrium
pressures for 1762 regions in 28 nearby galaxies.

Total internal pressures are calculated as $P_{\rm int}=P_{\rm th}+P_{\rm
rad}+P_{\rm CR}+P_{\rm B}$ and are compared with the dynamical equilibrium
pressures estimated via linear interpolation using the measured star formation
rate surface densities and the correlation between $\Sigma_{\rm SFR}$ and
$P_{\rm DE}$ found by Sun et al. (2020), which has a scatter of $\approx$0.2
dex. While this correlation was strictly established on kpc-scale measurements
of $\Sigma_{\rm SFR}$ and $P_{\rm DE}$, it should still provide a
qualitatively meaningful comparison for the present investigation even though
$\Sigma_{\rm SFR}$ is measured on scales spanning $\sim 100$ pc to $\sim$3
kpc. For reference, the median dynamical equilibrium pressures for the SFRS
and GOALS nuclear and extranuclear regions are: $\langle P_{\rm
DE}\rangle=(188\pm 53)\times 10^{-12}$, $(143\pm 15)\times 10^{-12}$,
$(1773\pm 287)\times 10^{-12}$, and $(466\pm 72)\times 10^{-12}\,{\rm
dyn\,cm^{-2}}$, respectively.

Using this relation to estimate the dynamical equilibrium pressures for each
region, the median ratio of $P_{\rm int}/P_{\rm DE}$ is found to be
$\approx$0.3, with very few sources (i.e., $\approx$5%) appearing to be over-
pressurized. In Figure 9, the ratio of the total internal-to-dynamical
equilibrium pressures are plotted as a function of aperture diameter. Of the
sources that are found to be over-pressurized, each of them were measured in
sub-kpc apertures, again suggesting that some combination of the the internal
pressure terms are likely stronger on smaller spatial scales.

It is worth noting that there is a missing internal pressure term in this
analysis associated with the turbulent motions in the gas. However, given that
Sun et al. (2020) report that the turbulent gas pressure both correlates with,
and almost always exceeds, the dynamical equilibrium pressure on kpc scales,
its inclusion in the calculation above will simply result in every region
being over-pressurized. Again, taking the results from Sun et al. (2020), who
measures turbulent gas pressures on 120 pc scales, and linearly interpolating
from our estimates of the star formation rate surface densities, the median
turbulent gas pressures for the SFRS and GOALS nuclear and extranuclear
regions are: $\langle P_{\rm turb}\rangle=(200\pm 57)\times 10^{-12}$,
$(152\pm 16)\times 10^{-12}$, $(1874\pm 304)\times 10^{-12}$, and $(494\pm
76)\times 10^{-12}\,{\rm dyn\,cm^{-2}}$, respectively. Indeed, including the
estimates of $P_{\rm turb}$ in the calculation of $P_{\rm int}$ results in
every region being marginally over-pressurized, with a median ratio of $P_{\rm
int}/P_{\rm DE}\approx 1.3$.

The above result arises from that fact that in feedback models, expanding Hii
regions, radiation pressure, and supernovae all drive the observed turbulence,
which in turn supports the dynamical equilibrium pressure of the ISM (e.g.,
Thompson et al., 2005; Ostriker & Shetty, 2011; Faucher-Giguère et al., 2013;
Ostriker & Kim, 2022). Consequently, in this situation, it is not appropriate
to add the stellar feedback pressures with the turbulent pressure. Though,
this picture may be slightly more complicated given that recent models by
Krumholz et al. (2018) argue that gas accretion should be included as an
additional mechanism driving turbulence.

##  4 Discussion and Conclusions

Observational studies, either for entire galaxies or smaller regions within
them defined by a specific volume or density threshold, find that the gas
depletion time $t_{\rm dep}=\Sigma_{\rm gas}/\Sigma_{\rm SFR}$, where
$\Sigma_{\rm gas}$ is the gas surface density, is always $\sim 1-3$ orders of
magnitude longer than the free-fall time $t_{\rm ff}\propto 1/\sqrt{G\rho}$,
where $\rho$ is the volume density of the gas (e.g., Krumholz & Tan, 2007;
Evans et al., 2009; Krumholz et al., 2012; Federrath, 2013). In contrast,
numerical simulations of star cluster formation that do not include any form
of feedback generally produce $t_{\rm dep}\sim t_{\rm ff}$ (e.g., Klessen &
Burkert, 2000, 2001; Bate et al., 2003; Bonnell et al., 2003), clearly
indicating its importance in regulating the star formation process. In the
present analysis, a number of primary candidates for regulating star formation
(i.e., thermal gas, radiation, and magnetic pressure estimates) are compared
as a function of galaxy environment, star formation activity, and physical
areas over which the quantities are measured.

In comparing the thermal gas and radiation pressure estimates presented here
with those in the literature, they appear highly consistent to values measured
via other methods. For instance, Barnes et al. (2021) recently looked into the
various stellar (pre-supernova) feedback mechanisms (i.e., photoionized gas,
radiation, winds) from $\approx$6000 Hii regions within a sample of 19 nearby
galaxies on $\sim 50-100$ pc scales included in Physics at High Angular
resolution in Nearby GalaxieS (PHANGS) sample.

In their analysis, they calculated upper and lower limits for each of the
pressure estimates using minimum (i.e., clumpy structure: $\sim$10 pc) and
maximum (i.e., smooth structure: $\sim$100 pc) sizes for each region,
respectively. Thermal gas pressure estimates were made by using the [Sii]
$\lambda\lambda$6716,31 doublet to obtain electron densities and additionally
assuming a fixed electron temperature of 8000 K. Radiation pressures were
estimated by converting H$\alpha$ measurements into bolometric luminosities.
The mean values of their lower and upper limit pressure estimates for their
entire sample are shown as vertical lines in Figures 2 and 3, which clearly
illustrate that our radiation and thermal-gas pressure estimates are largely
consistent with theirs.

Similarly, minimum-energy magnetic field strengths, and therefore pressures,
are also consistent with other estimates in the literature. Typical values of
the isotropic turbulent field strengths in star-forming galaxy arms and bars
can range between $20-30\,\mu$G, while central starbursting regions have field
strengths as high as $50-100\,\mu$G (e.g., Beck, 2015, and reference therein).
As shown in Figure 1, these values are in good agreement with the minimum-
energy magnetic field strength distributions for extranuclar and nuclear
regions, as well as for normal galaxies and the starbursting GOALS galaxies.

The results from this analysis appear to indicate that the physical area, and
not necessarily environmental location (e.g., nuclear vs. extranuclear
regions) is the dominant factor in determining the relative importance of
individual pressure terms around star-forming regions. For instance, the
thermal gas pressure appears to be similar to, and in some cases exceeds, the
radiation pressure on sub-kpc scales. On the other hand, it seems that the
magnetic pressure is typically sub-dominant on sub-kpc scales and only starts
to play a significant role on few-kpc scales relative to radiation pressure.

The result of finding the magnetic pressure to start to dominate over the
radiation pressure on galaxy-disk scales is consistent with expectations based
on the well-known far-infrared–radio correlation for galaxies (e.g., de Jong
et al., 1985; Helou et al., 1985; Yun et al., 2001), whereby the ratio of
magnetic-to-radiation energy densities must be $\gtrsim$1 to ensure that
synchrotron and inverse Compton losses to radio continuum emitting cosmic-ray
electrons is fixed (e.g., Condon, 1992; Murphy, 2009). It is also worth
pointing out that even if the minimum-energy magnetic field is significantly
underestimating the true magnetic field strength on small scales, as has been
suggested for starburst galaxies (e.g., Thompson et al., 2006) and their
central molecular zones (e.g., Yoast-Hull et al., 2016), such that it is
similar to what is required for equipartition with the hydrostatic pressure of
the molecular clouds, this should not affect our overall results. Observations
suggest that molecular clouds are magnetically supercritical by a factor of
$\sim$2 (Crutcher, 2012, and references therein) while simulations also
indicate that a magnetic field of this strength is only likely to suppress the
overall star formation activity on such scales by factor of a few compared to
the purely hydrodynamic case (Price & Bate, 2009; Padoan & Nordlund, 2011).

These results likely suggest that it is the combination of all feedback
mechanisms that work together to regulate star formation on a range of
physical scales, keeping the process inefficient both in individual clouds and
galaxy disks. This conclusion is also supported by the relatively constant
ratio of total internal-to-dynamical equilibrium pressures as a function of
physical aperture area. These findings are in agreement with Krumholz et al.
(2014), who suggested that magnetic fields alone are unlikely to be the single
driver that is able to keep star formation rates regulated, but rather a
combination of magnetic fields with stellar feedback as the prime candidate
that regulate star formation. Accordingly, these observations support a
picture where stellar feedback processes in mature galaxy disks (e.g., thermal
gas and radiation pressure on dust) are able to be highly effective on the
scales of individual star-forming complexes. However, in the transition to
larger (i.e., few-kpc) scales, other processes (e.g., a galaxy’s large-scale
magnetic field) become the dominant mechanism regulating star formation. This
scenario is consistent with the magnetohydrodynamical calculations of Price &
Bate (2009), who found that magnetic fields primarily provide support on large
scales to low-density gas, whereas radiation is found to strongly suppress
small-scale fragmentation of clouds by increasing the temperature in the high-
density material near the protostars.

While only a handful of stellar feedback process have been looked at in this
study, which is additionally limited by the fixed angular resolution of the
data in hand, the distance of the sample galaxies, and the lack of independent
estimates for the dynamical equilibrium pressure, the general conclusion of
the importance of multiple processes being effective on different physical
scales to collectively regulate the star formation efficiency of galaxies is
most likely robust. In considering the effects of wind pressure in addition to
the pressure terms from stellar feedback also considered here, Barnes et al.
(2021) found that each of these processes were generally comparable, and their
sum was in excess of the kpc-scale dynamical equilibrium pressure for 99% of
their sample when a clumpy ($\sim$10 pc) structure was assumed. When a smooth
($\sim$100 pc) structure was assumed, the fraction of sources with a total
internal pressure larger than the kpc-scale dynamical equilibrium pressure
decreases. This is in general agreement with the results of this analysis,
which is limited to $\gtrsim$100 pc scales. Consequently, it appears that each
of these physical feedback processes exhibit differing degrees of relative
importance on sub-kpc scales, and that their total sum is less than the
dynamical equilibrium pressure on kpc scales unless an additional contribution
from the turbulent gas pressure is included. However, to definitively
determine the exact pressure balance among each of these processes, and
properly assess the physical scale(s) over which each term may dominate,
requires substantially more data across the electromagnetic spectrum and at a
range of angular resolutions.

EJM would like the thank the anonymous referee for a very careful reading of
the text and substantial comments that helped improve the content of this
paper. EJM would like to thank Sean Linden for making all of his catalogs
available for this study. EJM is grateful to Shane Davis for a close reading
of the text and for a number of very useful comments that helped improved the
paper. EJM is also grateful to Rainer Beck for a number of stimulating
discussions that improved this work. The National Radio Astronomy Observatory
is a facility of the National Science Foundation operated under cooperative
agreement by Associated Universities, Inc.

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In Table 1, estimated physical quantities for all regions are provided.

Table 1: Estimated Physical Quantities for Each Region

Source ID | R.A. | Decl. | $D_{\mathrm{Ap}}$ | $\Sigma_{\mathrm{SFR}}$ | $P_{\mathrm{B}}$ | $P_{\mathrm{CR}}$ | $P_{\mathrm{rad}}$ | $P_{\mathrm{th}}$ | $P_{\mathrm{DE}}^{\dagger}$ | $P_{\mathrm{turb}}^{\dagger}$
---|---|---|---|---|---|---|---|---|---|---
| (J2000) | (J2000) | (kpc) | ($M_{\odot}$ yr-1 kpc-2) | (10-12 dyn cm-2)
SFRS Nuclear Regions
NGC 0855 | $02~{}14~{}03.68$ | $+27~{}52~{}37.9$ | 0.330 | $0.29\pm 0.02$ | $13.94\pm 3.90$ | $5.08\pm 1.42$ | $8.89\pm 0.44$ | $10.76\pm 3.18$ | $167.18\pm 14.60$ | $178.22\pm 15.57$
NGC 0925 | $02~{}27~{}17.00$ | $+33~{}34~{}43.0$ | 0.310 | $0.04\pm 0.02$ | $4.55\pm 3.20$ | $1.66\pm 1.17$ | $3.10\pm 0.31$ | $3.95\pm 2.82$ | $19.15\pm 9.77$ | $21.84\pm 11.14$
NGC 1097 | $02~{}46~{}18.98$ | $-30~{}16~{}28.8$ | 0.482 | $2.50\pm 0.20$ | $57.60\pm 15.60$ | $20.98\pm 5.68$ | $68.83\pm 3.39$ | $26.36\pm 7.55$ | $1481.86\pm 121.42$ | $1567.14\pm 128.41$
NGC 1482 | $03~{}54~{}38.97$ | $-20~{}30~{}07.8$ | 0.767 | $3.40\pm 0.51$ | $180.71\pm 66.84$ | $65.83\pm 24.35$ | $497.38\pm 24.86$ | $24.35\pm 9.40$ | $2013.84\pm 299.83$ | $2129.16\pm 317.00$
NGC 2798 | $09~{}17~{}22.85$ | $+42~{}00~{}00.4$ | 0.876 | $0.28\pm 0.17$ | $124.44\pm 95.67$ | $45.33\pm 34.85$ | $423.48\pm 21.08$ | $6.55\pm 5.09$ | $164.28\pm 99.28$ | $175.17\pm 105.86$
NGC 2841 | $09~{}22~{}02.67$ | $+50~{}58~{}35.7$ | 0.479 | $0.84\pm 0.05$ | $13.08\pm 3.11$ | $4.76\pm 1.13$ | $6.51\pm 0.34$ | $15.37\pm 3.89$ | $498.74\pm 31.93$ | $528.51\pm 33.84$
NGC 3049 | $09~{}54~{}49.56$ | $+09~{}16~{}16.1$ | 0.652 | $0.66\pm 0.05$ | $26.74\pm 6.92$ | $9.74\pm 2.52$ | $74.90\pm 3.58$ | $11.63\pm 3.19$ | $388.72\pm 29.19$ | $412.27\pm 30.96$
NGC 3077 | $10~{}03~{}19.15$ | $+68~{}43~{}59.9$ | 0.130 | $2.50\pm 0.16$ | $37.29\pm 8.96$ | $13.58\pm 3.27$ |  $\cdots$ | $50.71\pm 12.95$ | $1479.37\pm 96.52$ | $1564.51\pm 102.07$
NGC 3190 | $10~{}18~{}05.64$ | $+21~{}49~{}55.9$ | 0.655 | $0.25\pm 0.03$ | $27.96\pm 8.60$ | $10.19\pm 3.13$ | $12.05\pm 0.60$ | $7.20\pm 2.33$ | $148.45\pm 15.52$ | $158.44\pm 16.57$
NGC 3184 | $10~{}18~{}16.94$ | $+41~{}25~{}27.0$ | 0.397 | $0.16\pm 0.02$ | $13.62\pm 4.33$ | $4.96\pm 1.58$ | $22.91\pm 1.10$ | $7.28\pm 2.43$ | $91.14\pm 10.15$ | $97.90\pm 10.90$
NGC 3198 | $10~{}19~{}54.99$ | $+45~{}32~{}59.3$ | 0.479 | $0.18\pm 0.02$ | $20.64\pm 7.06$ | $7.52\pm 2.57$ | $68.88\pm 3.44$ | $7.06\pm 2.53$ | $103.56\pm 13.26$ | $111.01\pm 14.22$
NGC 3265 | $10~{}31~{}06.77$ | $+28~{}47~{}48.0$ | 0.665 | $0.32\pm 0.04$ | $39.83\pm 12.91$ | $14.51\pm 4.70$ | $56.88\pm 2.77$ | $8.04\pm 2.73$ | $188.47\pm 21.78$ | $200.72\pm 23.20$
NGC 3351 a | $10~{}43~{}57.68$ | $+11~{}42~{}08.0$ | 0.317 | $0.69\pm 0.06$ | $46.16\pm 13.35$ | $16.82\pm 4.87$ | $122.64\pm 5.99$ | $17.11\pm 5.22$ | $408.71\pm 38.02$ | $433.39\pm 40.31$
NGC 3627 | $11~{}20~{}15.00$ | $+12~{}59~{}29.4$ | 0.318 | $0.78\pm 0.09$ | $63.02\pm 20.10$ | $22.96\pm 7.32$ | $96.00\pm 4.79$ | $18.12\pm 6.07$ | $461.11\pm 51.69$ | $488.76\pm 54.79$
NGC 3773 | $11~{}38~{}13.02$ | $+12~{}06~{}43.8$ | 0.421 | $0.47\pm 0.04$ | $19.86\pm 5.24$ | $7.23\pm 1.91$ | $35.35\pm 2.23$ | $12.21\pm 3.41$ | $275.80\pm 21.50$ | $292.98\pm 22.84$
NGC 4254 b | $12~{}18~{}49.67$ | $+14~{}24~{}59.0$ | 0.489 | $0.15\pm 0.03$ | $2.68\pm 1.08$ | $0.98\pm 0.39$ | $21.28\pm 1.04$ | $6.37\pm 2.67$ | $85.56\pm 15.04$ | $92.00\pm 16.17$
NGC 4536 | $12~{}34~{}27.06$ | $+02~{}11~{}18.2$ | 0.492 | $3.52\pm 0.38$ | $138.49\pm 43.43$ | $50.45\pm 15.82$ | $379.83\pm 18.98$ | $30.93\pm 10.19$ | $2083.77\pm 226.12$ | $2203.04\pm 239.06$
NGC 4569 | $12~{}36~{}49.80$ | $+13~{}09~{}46.6$ | 0.335 | $0.20\pm 0.04$ | $51.93\pm 23.40$ | $18.92\pm 8.52$ | $125.24\pm 6.10$ | $8.93\pm 4.16$ | $116.13\pm 25.23$ | $124.30\pm 27.00$
NGC 4579 | $12~{}37~{}43.52$ | $+11~{}49~{}05.6$ | 0.557 | $24.47\pm 1.37$ | $20.87\pm 4.63$ | $7.60\pm 1.69$ | $38.62\pm 1.89$ | $76.74\pm 18.15$ | $14522.11\pm 812.14$ | $15343.71\pm 858.09$
NGC 4594 | $12~{}39~{}59.42$ | $-11~{}37~{}23.0$ | 0.308 | $44.58\pm 2.58$ | $29.28\pm 6.61$ | $10.67\pm 2.41$ | $14.73\pm 0.70$ | $139.18\pm 33.49$ | $26450.69\pm 1531.74$ | $27945.84\pm 1618.33$
NGC 4725 a | $12~{}50~{}26.56$ | $+25~{}30~{}03.0$ | 0.404 | $0.05\pm 0.01$ | $1.80\pm 0.93$ | $0.66\pm 0.34$ | $8.71\pm 0.42$ | $4.15\pm 2.21$ | $28.69\pm 8.15$ | $31.92\pm 9.07$
NGC 4736 | $12~{}50~{}53.05$ | $+41~{}07~{}12.8$ | 0.158 | $0.90\pm 0.08$ | $54.20\pm 15.74$ | $19.75\pm 5.73$ | $120.39\pm 5.90$ | $27.63\pm 8.45$ | $532.75\pm 49.89$ | $564.44\pm 52.86$
NGC 4826 | $12~{}56~{}43.56$ | $+21~{}41~{}00.6$ | 0.179 | $0.71\pm 0.08$ | $60.55\pm 19.17$ | $22.06\pm 6.98$ | $77.84\pm 3.86$ | $23.13\pm 7.69$ | $421.97\pm 46.64$ | $447.41\pm 49.45$
NGC 5055 | $13~{}15~{}49.31$ | $+42~{}01~{}45.1$ | 0.269 | $0.17\pm 0.02$ | $22.57\pm 8.05$ | $8.22\pm 2.93$ | $16.25\pm 0.79$ | $9.09\pm 3.39$ | $96.54\pm 13.40$ | $103.60\pm 14.38$
NGC 5194 d | $13~{}29~{}52.73$ | $+47~{}11~{}40.6$ | 0.259 | $0.02\pm 0.04$ | $67.40\pm 92.36$ | $23.60\pm 32.34$ | $38.23\pm 1.86$ | $3.43\pm 4.66$ | $11.26\pm 20.83$ | $13.51\pm 24.98$
NGC 5457 | $14~{}03~{}12.53$ | $+54~{}20~{}55.2$ | 0.227 | $0.15\pm 0.02$ | $13.50\pm 4.41$ | $4.92\pm 1.61$ | $23.74\pm 1.13$ | $9.33\pm 3.20$ | $85.55\pm 10.05$ | $91.99\pm 10.80$
NGC 5866 | $15~{}06~{}29.50$ | $+55~{}45~{}47.7$ | 0.519 | $2.49\pm 0.18$ | $56.89\pm 14.34$ | $20.73\pm 5.22$ | $7.88\pm 0.39$ | $25.32\pm 6.77$ | $1473.02\pm 105.16$ | $1557.80\pm 111.21$
NGC 6946 b | $20~{}34~{}52.26$ | $+60~{}09~{}14.3$ | 0.231 | $5.13\pm 0.45$ | $126.78\pm 35.44$ | $46.19\pm 12.91$ | $806.95\pm 40.34$ | $54.55\pm 16.10$ | $3041.22\pm 264.84$ | $3214.55\pm 279.93$
NGC 7331 | $22~{}37~{}04.10$ | $+34~{}24~{}56.0$ | 0.492 | $0.05\pm 0.03$ | $3.02\pm 2.38$ | $1.10\pm 0.87$ | $6.64\pm 0.33$ | $3.66\pm 2.92$ | $27.08\pm 17.21$ | $30.22\pm 19.21$
NGC 7793 a | $23~{}57~{}49.20$ | $-32~{}35~{}24.0$ | 0.133 | $0.05\pm 0.03$ | $3.08\pm 2.10$ | $1.12\pm 0.76$ | $6.34\pm 0.39$ | $7.42\pm 5.14$ | $30.16\pm 14.46$ | $33.47\pm 16.05$
NGC 7793 b | $23~{}57~{}49.58$ | $-32~{}35~{}25.6$ | 0.133 | $0.07\pm 0.02$ | $5.47\pm 3.06$ | $1.99\pm 1.11$ | $9.47\pm 0.60$ | $8.39\pm 4.81$ | $39.24\pm 12.87$ | $43.06\pm 14.13$
SFRS Extranuclear Regions
NGC 0337 a | $00~{}59~{}50.02$ | $-07~{}34~{}33.9$ | 0.655 | $0.34\pm 0.03$ | $14.30\pm 4.16$ | $5.21\pm 1.52$ | $18.62\pm 0.90$ | $8.30\pm 2.55$ | $197.95\pm 18.63$ | $210.73\pm 19.83$
NGC 0337 b | $00~{}59~{}50.68$ | $-07~{}34~{}57.6$ | 0.655 | $0.70\pm 0.06$ | $31.80\pm 8.43$ | $11.58\pm 3.07$ | $32.90\pm 1.60$ | $12.00\pm 3.36$ | $415.80\pm 32.68$ | $440.88\pm 34.65$
NGC 0337 c | $00~{}59~{}51.99$ | $-07~{}34~{}54.9$ | 0.655 | $0.14\pm 0.03$ | $3.12\pm 1.48$ | $1.14\pm 0.54$ | $3.67\pm 0.18$ | $5.35\pm 2.62$ | $81.04\pm 19.44$ | $87.23\pm 20.93$
NGC 0337 d | $00~{}59~{}52.16$ | $-07~{}34~{}38.2$ | 0.655 | $0.28\pm 0.04$ | $3.17\pm 1.08$ | $1.16\pm 0.39$ | $9.17\pm 0.44$ | $7.57\pm 2.70$ | $164.08\pm 20.88$ | $174.95\pm 22.26$
NGC 0628 Enuc. 4 | $01~{}36~{}35.72$ | $+15~{}50~{}07.2$ | 0.244 | $0.15\pm 0.02$ | $2.44\pm 0.88$ | $0.89\pm 0.32$ | $4.58\pm 0.22$ | $9.10\pm 3.44$ | $87.40\pm 12.53$ | $93.94\pm 13.47$
NGC 0628 Enuc. 2 | $01~{}36~{}37.65$ | $+15~{}45~{}07.2$ | 0.244 | $0.21\pm 0.02$ | $7.90\pm 2.48$ | $2.88\pm 0.90$ | $16.47\pm 0.79$ | $10.83\pm 3.57$ | $124.78\pm 13.55$ | $133.43\pm 14.49$
NGC 0628 Enuc. 3 | $01~{}36~{}38.78$ | $+15~{}44~{}23.2$ | 0.244 | $0.27\pm 0.03$ | $6.07\pm 1.89$ | $2.21\pm 0.69$ | $13.82\pm 0.68$ | $12.06\pm 3.94$ | $155.39\pm 16.58$ | $165.77\pm 17.69$
NGC 0628 Enuc. 1 | $01~{}36~{}45.27$ | $+15~{}47~{}48.3$ | 0.244 | $0.45\pm 0.04$ | $11.88\pm 3.24$ | $4.33\pm 1.18$ | $30.31\pm 1.47$ | $15.67\pm 4.51$ | $263.77\pm 21.85$ | $280.27\pm 23.22$
NGC 1097 Enuc. 1 b | $02~{}46~{}22.93$ | $-30~{}17~{}48.1$ | 0.482 | $0.11\pm 0.06$ | $14.31\pm 10.14$ | $5.21\pm 3.69$ | $3.74\pm 0.18$ | $5.52\pm 3.97$ | $62.89\pm 32.53$ | $68.05\pm 35.20$
NGC 1266 | $03~{}16~{}00.76$ | $-02~{}25~{}37.1$ | 1.04 | $1.58\pm 0.31$ | $152.25\pm 65.52$ | $55.46\pm 23.87$ | $142.16\pm 7.11$ | $14.26\pm 6.36$ | $933.46\pm 185.71$ | $987.78\pm 196.51$
NGC 2403 Enuc. 6 | $07~{}36~{}28.69$ | $+65~{}33~{}49.4$ | 0.109 | $0.39\pm 0.03$ | $2.89\pm 0.77$ | $1.05\pm 0.28$ | $7.41\pm 0.61$ | $21.90\pm 6.16$ | $230.09\pm 18.22$ | $244.69\pm 19.37$
NGC 2403 Enuc. 1 a | $07~{}36~{}42.06$ | $+65~{}36~{}51.9$ | 0.109 | $0.21\pm 0.03$ | $8.04\pm 2.68$ | $2.93\pm 0.97$ | $1.97\pm 0.11$ | $16.01\pm 5.58$ | $121.91\pm 14.82$ | $130.40\pm 15.85$
NGC 2403 Enuc. 1 b | $07~{}36~{}45.50$ | $+65~{}37~{}00.9$ | 0.109 | $0.75\pm 0.05$ | $3.75\pm 0.91$ | $1.37\pm 0.33$ | $51.75\pm 3.30$ | $30.34\pm 7.84$ | $443.50\pm 29.62$ | $470.15\pm 31.40$
NGC 2403 Enuc. 2 a | $07~{}36~{}49.12$ | $+65~{}36~{}51.7$ | 0.109 | $0.37\pm 0.03$ | $5.00\pm 1.40$ | $1.82\pm 0.51$ | $14.39\pm 0.69$ | $21.35\pm 6.33$ | $218.47\pm 19.19$ | $232.42\pm 20.41$
NGC 2403 Enuc. 2 b | $07~{}36~{}52.36$ | $+65~{}36~{}46.9$ | 0.109 | $0.71\pm 0.05$ | $5.71\pm 1.39$ | $2.08\pm 0.51$ | $32.25\pm 1.81$ | $29.39\pm 7.62$ | $416.25\pm 27.97$ | $441.37\pm 29.66$
NGC 2403 Enuc. 4 | $07~{}37~{}18.19$ | $+65~{}33~{}48.1$ | 0.109 | $0.38\pm 0.03$ | $2.82\pm 0.74$ | $1.03\pm 0.27$ | $8.72\pm 0.45$ | $21.66\pm 6.04$ | $224.98\pm 17.51$ | $239.29\pm 18.62$
Holmberg II | $08~{}19~{}13.06$ | $+70~{}43~{}08.0$ | 0.104 | $0.22\pm 0.03$ | $2.66\pm 0.87$ | $0.97\pm 0.32$ | $10.55\pm 0.70$ | $16.85\pm 5.79$ | $128.04\pm 15.14$ | $136.88\pm 16.19$
NGC 2976 Enuc. 1 a | $09~{}47~{}05.19$ | $+67~{}55~{}52.0$ | 0.120 | $0.46\pm 0.04$ | $11.25\pm 3.01$ | $4.10\pm 1.10$ | $10.72\pm 0.53$ | $22.61\pm 6.41$ | $270.85\pm 21.76$ | $287.75\pm 23.11$
NGC 2976 Enuc. 1 b | $09~{}47~{}07.64$ | $+67~{}55~{}54.7$ | 0.120 | $1.28\pm 0.08$ | $11.04\pm 2.61$ | $4.02\pm 0.95$ | $51.97\pm 2.52$ | $37.65\pm 9.46$ | $754.83\pm 47.60$ | $799.06\pm 50.39$
NGC 2976 | $09~{}47~{}13.49$ | $+67~{}54~{}54.0$ | 0.120 | $0.12\pm 0.03$ | $2.26\pm 1.01$ | $0.82\pm 0.37$ | $3.56\pm 0.18$ | $11.41\pm 5.29$ | $67.33\pm 14.48$ | $72.74\pm 15.64$
NGC 2976 Enuc. 2 a | $09~{}47~{}23.83$ | $+67~{}53~{}54.9$ | 0.120 | $0.67\pm 0.05$ | $4.05\pm 1.02$ | $1.48\pm 0.37$ | $22.20\pm 1.06$ | $27.35\pm 7.32$ | $397.10\pm 28.47$ | $421.13\pm 30.19$
NGC 2976 Enuc. 2 b | $09~{}47~{}23.94$ | $+67~{}54~{}02.1$ | 0.120 | $0.31\pm 0.03$ | $4.24\pm 1.29$ | $1.54\pm 0.47$ | $9.61\pm 0.46$ | $18.45\pm 5.92$ | $179.65\pm 18.46$ | $191.40\pm 19.67$
IC 2574 a | $10~{}28~{}43.71$ | $+68~{}28~{}26.3$ | 0.129 | $0.33\pm 0.03$ | $2.39\pm 0.73$ | $0.87\pm 0.27$ | $8.65\pm 1.03$ | $18.40\pm 5.93$ | $190.86\pm 19.83$ | $203.25\pm 21.12$
IC 2574 b | $10~{}28~{}48.40$ | $+68~{}28~{}03.5$ | 0.129 | $0.34\pm 0.03$ | $1.48\pm 0.43$ | $0.54\pm 0.16$ | $9.02\pm 0.43$ | $18.83\pm 5.74$ | $199.94\pm 18.58$ | $212.84\pm 19.78$
NGC 3351 b | $10~{}43~{}57.80$ | $+11~{}42~{}18.5$ | 0.317 | $0.82\pm 0.07$ | $49.57\pm 14.16$ | $18.06\pm 5.16$ | $154.04\pm 7.59$ | $18.68\pm 5.62$ | $487.21\pm 44.20$ | $516.33\pm 46.84$
NGC 3521 Enuc. 1 | $11~{}05~{}46.30$ | $-00~{}04~{}09.0$ | 0.380 | $0.05\pm 0.04$ | $3.29\pm 2.92$ | $1.20\pm 1.06$ | $2.57\pm 0.12$ | $4.02\pm 3.59$ | $25.08\pm 19.92$ | $28.11\pm 22.32$
NGC 3521 Enuc. 2 a | $11~{}05~{}49.34$ | $-00~{}03~{}24.2$ | 0.380 | $0.13\pm 0.04$ | $12.12\pm 6.16$ | $4.41\pm 2.24$ | $7.45\pm 0.37$ | $6.87\pm 3.59$ | $77.21\pm 21.12$ | $83.18\pm 22.75$
NGC 3627 Enuc. 1 | $11~{}20~{}16.32$ | $+12~{}57~{}49.2$ | 0.318 | $1.24\pm 0.08$ | $15.40\pm 3.67$ | $5.61\pm 1.34$ | $66.16\pm 3.30$ | $22.85\pm 5.79$ | $734.36\pm 47.10$ | $777.43\pm 49.86$
NGC 3627 Enuc. 2 | $11~{}20~{}16.46$ | $+12~{}58~{}43.4$ | 0.318 | $2.40\pm 0.16$ | $48.89\pm 12.07$ | $17.81\pm 4.40$ | $186.40\pm 9.30$ | $31.79\pm 8.33$ | $1423.14\pm 97.70$ | $1505.10\pm 103.33$
NGC 3938 Enuc. 2 a | $11~{}53~{}00.06$ | $+44~{}08~{}00.0$ | 0.607 | $0.10\pm 0.02$ | $2.58\pm 1.25$ | $0.94\pm 0.46$ | $3.81\pm 0.19$ | $4.60\pm 2.29$ | $54.67\pm 13.63$ | $59.36\pm 14.80$
NGC 3938 Enuc. 2 b | $11~{}53~{}00.20$ | $+44~{}07~{}48.3$ | 0.607 | $0.17\pm 0.03$ | $1.94\pm 0.76$ | $0.71\pm 0.28$ | $9.35\pm 0.44$ | $6.14\pm 2.51$ | $99.30\pm 16.61$ | $106.51\pm 17.82$
NGC 4254 Enuc. 2 a | $12~{}18~{}45.78$ | $+14~{}24~{}10.4$ | 0.489 | $0.04\pm 0.02$ | $2.01\pm 1.34$ | $0.73\pm 0.49$ | $6.09\pm 0.29$ | $3.20\pm 2.16$ | $19.98\pm 9.12$ | $22.72\pm 10.37$
NGC 4254 Enuc. 2 b | $12~{}18~{}46.13$ | $+14~{}24~{}18.8$ | 0.489 | $0.06\pm 0.02$ | $2.11\pm 1.11$ | $0.77\pm 0.40$ | $5.19\pm 0.25$ | $4.17\pm 2.25$ | $35.49\pm 10.31$ | $39.10\pm 11.36$
NGC 4254 a | $12~{}18~{}48.68$ | $+14~{}24~{}42.5$ | 0.489 | $0.05\pm 0.01$ | $10.76\pm 5.46$ | $3.76\pm 1.91$ | $13.10\pm 0.64$ | $3.82\pm 1.97$ | $29.49\pm 7.81$ | $32.76\pm 8.68$
NGC 4254 Enuc. 1 b | $12~{}18~{}50.01$ | $+14~{}24~{}06.9$ | 0.489 | $0.08\pm 0.02$ | $2.14\pm 1.03$ | $0.78\pm 0.38$ | $8.52\pm 0.42$ | $4.57\pm 2.26$ | $42.96\pm 10.57$ | $47.00\pm 11.56$
NGC 4254 c | $12~{}18~{}50.10$ | $+14~{}25~{}11.6$ | 0.489 | $0.04\pm 0.02$ | $2.26\pm 1.55$ | $0.82\pm 0.56$ | $10.51\pm 0.51$ | $3.32\pm 2.31$ | $21.73\pm 10.52$ | $24.57\pm 11.90$
NGC 4254 Enuc. 1 c | $12~{}18~{}50.19$ | $+14~{}24~{}18.6$ | 0.489 | $0.15\pm 0.02$ | $13.76\pm 4.72$ | $5.01\pm 1.72$ | $13.31\pm 0.64$ | $6.47\pm 2.32$ | $88.55\pm 11.42$ | $95.16\pm 12.27$
NGC 4254 d | $12~{}18~{}51.63$ | $+14~{}25~{}08.6$ | 0.489 | $0.12\pm 0.02$ | $17.49\pm 7.42$ | $5.07\pm 2.15$ | $6.08\pm 0.29$ | $5.68\pm 2.26$ | $67.78\pm 10.71$ | $73.22\pm 11.57$
NGC 4254 e | $12~{}18~{}51.90$ | $+14~{}24~{}49.7$ | 0.489 | $0.09\pm 0.02$ | $2.52\pm 1.28$ | $0.92\pm 0.47$ | $10.54\pm 0.50$ | $4.84\pm 2.53$ | $48.45\pm 13.22$ | $52.80\pm 14.41$
NGC 4254 f | $12~{}18~{}51.92$ | $+14~{}24~{}40.1$ | 0.489 | $0.04\pm 0.02$ | $2.15\pm 1.50$ | $0.78\pm 0.55$ | $5.02\pm 0.24$ | $3.18\pm 2.24$ | $19.68\pm 9.80$ | $22.40\pm 11.16$
NGC 4321 Enuc. 2 a | $12~{}22~{}48.84$ | $+15~{}50~{}12.8$ | 0.485 | $0.06\pm 0.02$ | $2.61\pm 1.61$ | $0.95\pm 0.59$ | $3.89\pm 0.19$ | $4.10\pm 2.59$ | $34.04\pm 13.50$ | $37.57\pm 14.90$
NGC 4321 Enuc. 2 b | $12~{}22~{}49.90$ | $+15~{}50~{}27.8$ | 0.485 | $0.07\pm 0.02$ | $1.98\pm 1.11$ | $0.72\pm 0.40$ | $3.32\pm 0.18$ | $4.32\pm 2.49$ | $38.00\pm 12.59$ | $41.76\pm 13.84$
NGC 4321 a | $12~{}22~{}54.65$ | $+15~{}49~{}19.8$ | 0.485 | $0.35\pm 0.04$ | $40.26\pm 13.03$ | $14.67\pm 4.75$ | $39.57\pm 1.89$ | $9.86\pm 3.35$ | $206.83\pm 23.84$ | $220.12\pm 25.37$
NGC 4321 b | $12~{}22~{}55.13$ | $+15~{}49~{}20.4$ | 0.485 | $0.42\pm 0.05$ | $50.26\pm 16.98$ | $18.31\pm 6.18$ | $47.78\pm 2.30$ | $10.77\pm 3.81$ | $247.16\pm 30.92$ | $262.72\pm 32.87$
NGC 4559 a | $12~{}35~{}56.27$ | $+27~{}57~{}40.5$ | 0.237 | $0.08\pm 0.02$ | $2.49\pm 1.27$ | $0.91\pm 0.46$ | $5.20\pm 0.28$ | $6.85\pm 3.60$ | $47.06\pm 13.00$ | $51.32\pm 14.18$
NGC 4559 b | $12~{}35~{}56.45$ | $+27~{}57~{}21.3$ | 0.237 | $0.05\pm 0.02$ | $2.39\pm 1.61$ | $0.87\pm 0.59$ | $4.46\pm 0.35$ | $5.05\pm 3.46$ | $24.57\pm 11.55$ | $27.56\pm 12.96$
NGC 4559 c | $12~{}35~{}58.47$ | $+27~{}57~{}29.7$ | 0.237 | $0.11\pm 0.02$ | $2.43\pm 1.06$ | $0.88\pm 0.38$ | $8.05\pm 0.38$ | $7.90\pm 3.56$ | $63.31\pm 12.88$ | $68.50\pm 13.93$
NGC 4631 Enuc. 1 | $12~{}41~{}40.47$ | $+32~{}31~{}49.1$ | 0.259 | $0.17\pm 0.02$ | $2.43\pm 0.84$ | $0.89\pm 0.30$ | $6.30\pm 0.53$ | $9.46\pm 3.40$ | $100.42\pm 12.95$ | $107.70\pm 13.89$
NGC 4631 a | $12~{}42~{}03.43$ | $+32~{}32~{}17.2$ | 0.259 | $0.38\pm 0.06$ | $55.32\pm 20.73$ | $20.15\pm 7.55$ | $16.83\pm 0.81$ | $13.94\pm 5.45$ | $220.60\pm 33.68$ | $234.66\pm 35.83$
NGC 4631 b | $12~{}42~{}03.98$ | $+32~{}32~{}16.0$ | 0.259 | $0.76\pm 0.06$ | $29.77\pm 7.78$ | $10.85\pm 2.84$ | $34.31\pm 1.67$ | $19.81\pm 5.48$ | $447.39\pm 34.26$ | $474.26\pm 36.32$
NGC 4631 c | $12~{}42~{}04.31$ | $+32~{}32~{}25.3$ | 0.259 | $1.68\pm 0.11$ | $28.95\pm 6.94$ | $10.55\pm 2.53$ | $66.58\pm 3.26$ | $29.47\pm 7.51$ | $993.25\pm 64.42$ | $1050.94\pm 68.16$
NGC 4631 d | $12~{}42~{}05.09$ | $+32~{}32~{}10.6$ | 0.259 | $0.37\pm 0.03$ | $5.90\pm 1.66$ | $2.15\pm 0.60$ | $18.28\pm 0.89$ | $13.90\pm 4.12$ | $219.35\pm 19.30$ | $233.34\pm 20.53$
NGC 4631 e | $12~{}42~{}05.57$ | $+32~{}32~{}29.5$ | 0.259 | $0.78\pm 0.05$ | $14.52\pm 3.54$ | $5.29\pm 1.29$ | $37.27\pm 1.81$ | $20.12\pm 5.21$ | $461.77\pm 31.02$ | $489.45\pm 32.88$
NGC 4631 Enuc. 2 b | $12~{}42~{}21.99$ | $+32~{}32~{}45.0$ | 0.259 | $0.50\pm 0.04$ | $6.07\pm 1.54$ | $2.21\pm 0.56$ | $19.02\pm 0.92$ | $16.11\pm 4.33$ | $295.15\pm 21.33$ | $313.42\pm 22.65$
NGC 4736 Enuc. 1 a | $12~{}50~{}56.41$ | $+41~{}07~{}14.3$ | 0.158 | $0.46\pm 0.04$ | $16.18\pm 4.41$ | $5.90\pm 1.61$ | $36.24\pm 1.72$ | $19.67\pm 5.66$ | $268.99\pm 22.28$ | $285.78\pm 23.67$
NGC 4736 Enuc. 1 b | $12~{}50~{}56.70$ | $+41~{}07~{}05.0$ | 0.158 | $0.47\pm 0.04$ | $24.82\pm 6.94$ | $9.04\pm 2.53$ | $36.65\pm 1.75$ | $19.87\pm 5.86$ | $274.43\pm 23.89$ | $291.53\pm 25.38$
NGC 4736 Enuc. 1 c | $12~{}50~{}56.78$ | $+41~{}06~{}47.6$ | 0.158 | $0.39\pm 0.04$ | $16.38\pm 5.09$ | $5.97\pm 1.85$ | $37.32\pm 1.79$ | $18.20\pm 5.94$ | $229.93\pm 24.51$ | $244.52\pm 26.06$
NGC 5055 Enuc. 1 | $13~{}15~{}58.32$ | $+42~{}00~{}27.4$ | 0.269 | $0.20\pm 0.03$ | $2.69\pm 0.95$ | $0.98\pm 0.35$ | $9.66\pm 0.46$ | $9.88\pm 3.64$ | $114.35\pm 15.49$ | $122.41\pm 16.58$
NGC 5194 Enuc. 6 a | $13~{}29~{}39.32$ | $+47~{}08~{}40.7$ | 0.259 | $0.30\pm 0.03$ | $2.91\pm 0.91$ | $1.06\pm 0.33$ | $10.54\pm 0.51$ | $12.39\pm 4.06$ | $173.62\pm 18.62$ | $185.03\pm 19.85$
NGC 5194 Enuc. 3 | $13~{}29~{}45.13$ | $+47~{}09~{}57.4$ | 0.259 | $0.38\pm 0.04$ | $1.82\pm 0.53$ | $0.66\pm 0.19$ | $27.13\pm 1.29$ | $14.03\pm 4.29$ | $223.28\pm 20.86$ | $237.50\pm 22.19$
NGC 5194 Enuc. 11 a | $13~{}29~{}47.14$ | $+47~{}13~{}41.3$ | 0.259 | $0.11\pm 0.02$ | $12.25\pm 4.84$ | $4.18\pm 1.65$ | $7.00\pm 0.37$ | $7.68\pm 3.07$ | $65.32\pm 10.44$ | $70.61\pm 11.28$
NGC 5194 Enuc. 11 c | $13~{}29~{}49.67$ | $+47~{}14~{}00.2$ | 0.259 | $0.02\pm 0.01$ | $18.85\pm 14.70$ | $6.87\pm 5.36$ | $1.22\pm 0.06$ | $3.43\pm 2.70$ | $11.30\pm 7.02$ | $13.55\pm 8.42$
NGC 5194 c | $13~{}29~{}50.02$ | $+47~{}11~{}31.9$ | 0.259 | $0.24\pm 0.03$ | $22.10\pm 7.38$ | $8.05\pm 2.69$ | $19.75\pm 0.98$ | $11.26\pm 3.94$ | $143.13\pm 17.54$ | $152.82\pm 18.73$
NGC 5194 Enuc. 11 e | $13~{}29~{}50.64$ | $+47~{}13~{}44.9$ | 0.259 | $0.19\pm 0.03$ | $5.76\pm 2.08$ | $2.10\pm 0.76$ | $9.18\pm 0.45$ | $9.85\pm 3.71$ | $108.91\pm 15.46$ | $116.67\pm 16.56$
NGC 5194 b | $13~{}29~{}51.64$ | $+47~{}12~{}06.7$ | 0.259 | $0.15\pm 0.02$ | $29.61\pm 11.45$ | $10.79\pm 4.17$ | $15.28\pm 0.73$ | $8.92\pm 3.59$ | $89.09\pm 14.45$ | $95.73\pm 15.53$
NGC 5194 e | $13~{}29~{}52.55$ | $+47~{}11~{}52.6$ | 0.259 | $0.13\pm 0.04$ | $54.21\pm 30.15$ | $19.75\pm 10.98$ | $13.34\pm 0.63$ | $8.06\pm 4.59$ | $72.17\pm 23.46$ | $77.86\pm 25.31$
NGC 5194 Enuc. 1 c | $13~{}29~{}53.13$ | $+47~{}12~{}39.4$ | 0.259 | $0.20\pm 0.03$ | $3.48\pm 1.32$ | $1.27\pm 0.48$ | $12.97\pm 0.63$ | $10.06\pm 3.98$ | $113.76\pm 17.80$ | $121.80\pm 19.06$
NGC 5194 Enuc. 10 a | $13~{}29~{}55.35$ | $+47~{}10~{}47.2$ | 0.259 | $0.12\pm 0.03$ | $3.96\pm 1.89$ | $1.44\pm 0.69$ | $7.87\pm 0.39$ | $7.91\pm 3.91$ | $69.57\pm 16.98$ | $75.11\pm 18.33$
NGC 5194 Enuc. 4 c | $13~{}29~{}55.61$ | $+47~{}13~{}50.2$ | 0.259 | $0.10\pm 0.02$ | $7.58\pm 3.15$ | $2.76\pm 1.15$ | $5.57\pm 0.26$ | $7.37\pm 3.18$ | $60.09\pm 11.20$ | $65.09\pm 12.13$
NGC 5194 a | $13~{}29~{}55.79$ | $+47~{}11~{}45.1$ | 0.259 | $0.44\pm 0.04$ | $11.91\pm 3.38$ | $4.34\pm 1.23$ | $43.45\pm 2.08$ | $15.18\pm 4.54$ | $261.81\pm 23.44$ | $278.20\pm 24.91$
NGC 5194 Enuc. 10 b | $13~{}29~{}56.52$ | $+47~{}10~{}46.9$ | 0.259 | $0.19\pm 0.03$ | $8.47\pm 3.07$ | $3.09\pm 1.12$ | $15.37\pm 0.75$ | $9.82\pm 3.71$ | $108.43\pm 15.47$ | $116.16\pm 16.58$
NGC 5194 Enuc. 4 d | $13~{}29~{}58.73$ | $+47~{}14~{}09.4$ | 0.259 | $0.09\pm 0.02$ | $6.80\pm 3.13$ | $2.48\pm 1.14$ | $10.18\pm 0.48$ | $6.86\pm 3.26$ | $51.68\pm 11.71$ | $56.21\pm 12.74$
NGC 5194 Enuc. 5 a | $13~{}29~{}59.60$ | $+47~{}13~{}59.8$ | 0.259 | $0.23\pm 0.03$ | $7.33\pm 2.55$ | $2.67\pm 0.93$ | $15.28\pm 0.75$ | $10.80\pm 3.93$ | $131.39\pm 17.40$ | $140.41\pm 18.59$
NGC 5194 Enuc. 9 | $13~{}29~{}59.78$ | $+47~{}11~{}12.3$ | 0.259 | $0.21\pm 0.03$ | $2.98\pm 1.14$ | $1.08\pm 0.41$ | $20.77\pm 1.01$ | $10.44\pm 4.15$ | $122.66\pm 19.43$ | $131.20\pm 20.78$
NGC 5194 Enuc. 7 a | $13~{}30~{}01.03$ | $+47~{}09~{}28.6$ | 0.259 | $0.07\pm 0.03$ | $2.70\pm 1.70$ | $0.98\pm 0.62$ | $4.22\pm 0.20$ | $5.94\pm 3.82$ | $38.24\pm 15.78$ | $42.01\pm 17.34$
NGC 5194 Enuc. 8 | $13~{}30~{}01.48$ | $+47~{}12~{}51.7$ | 0.259 | $0.44\pm 0.04$ | $8.36\pm 2.38$ | $3.05\pm 0.87$ | $44.47\pm 2.15$ | $15.01\pm 4.50$ | $256.11\pm 23.03$ | $272.18\pm 24.48$
NGC 5194 Enuc. 7 b | $13~{}30~{}02.38$ | $+47~{}09~{}48.7$ | 0.259 | $0.22\pm 0.03$ | $15.14\pm 4.92$ | $5.52\pm 1.79$ | $21.18\pm 1.01$ | $10.68\pm 3.64$ | $128.46\pm 14.94$ | $137.33\pm 15.97$
NGC 5194 Enuc. 7 c | $13~{}30~{}03.47$ | $+47~{}09~{}40.3$ | 0.259 | $0.04\pm 0.02$ | $10.81\pm 6.58$ | $3.94\pm 2.40$ | $6.50\pm 0.31$ | $4.66\pm 2.89$ | $22.65\pm 8.74$ | $25.53\pm 9.86$
NGC 5398 | $14~{}01~{}20.11$ | $-33~{}04~{}09.2$ | 0.260 | $0.91\pm 0.07$ | $8.06\pm 2.12$ | $2.94\pm 0.77$ | $40.97\pm 1.94$ | $21.66\pm 6.02$ | $538.22\pm 41.56$ | $570.22\pm 44.03$
NGC 5457 Enuc. 6 a | $14~{}02~{}28.20$ | $+54~{}16~{}27.2$ | 0.227 | $0.52\pm 0.04$ | $14.42\pm 3.67$ | $5.25\pm 1.34$ | $32.88\pm 1.79$ | $17.44\pm 4.70$ | $304.26\pm 22.09$ | $323.05\pm 23.45$
NGC 5457 Enuc. 6 b | $14~{}02~{}29.61$ | $+54~{}16~{}15.8$ | 0.227 | $0.56\pm 0.04$ | $21.98\pm 5.68$ | $8.01\pm 2.07$ | $26.46\pm 1.29$ | $18.08\pm 4.95$ | $327.33\pm 24.53$ | $347.42\pm 26.03$
NGC 5457 Enuc. 6 c | $14~{}02~{}30.57$ | $+54~{}16~{}9.8$ | 0.227 | $0.43\pm 0.04$ | $13.48\pm 3.65$ | $4.91\pm 1.33$ | $27.35\pm 1.30$ | $15.97\pm 4.57$ | $254.78\pm 20.88$ | $270.78\pm 22.20$
NGC 5457 Enuc. 5 a | $14~{}03~{}01.20$ | $+54~{}14~{}28.4$ | 0.227 | $1.01\pm 0.06$ | $8.92\pm 2.11$ | $3.25\pm 0.77$ | $40.27\pm 2.22$ | $24.41\pm 6.15$ | $598.38\pm 37.93$ | $633.77\pm 40.18$
NGC 5457 Enuc. 1 | $14~{}03~{}10.20$ | $+54~{}20~{}57.8$ | 0.227 | $0.08\pm 0.02$ | $2.73\pm 1.31$ | $0.99\pm 0.48$ | $5.67\pm 0.27$ | $6.81\pm 3.38$ | $44.53\pm 10.98$ | $48.65\pm 12.00$
NGC 5457 Enuc. 3 a | $14~{}03~{}38.32$ | $+54~{}18~{}51.4$ | 0.227 | $0.16\pm 0.06$ | $11.96\pm 6.98$ | $4.36\pm 2.54$ | $6.89\pm 0.33$ | $9.62\pm 5.74$ | $91.03\pm 32.46$ | $97.78\pm 34.87$
NGC 5457 Enuc. 3 b | $14~{}03~{}39.89$ | $+54~{}18~{}56.8$ | 0.227 | $0.55\pm 0.06$ | $2.25\pm 0.71$ | $0.82\pm 0.26$ | $22.31\pm 1.06$ | $18.01\pm 6.00$ | $324.84\pm 36.07$ | $344.79\pm 38.29$
NGC 5457 Enuc. 3 c | $14~{}03~{}41.44$ | $+54~{}19~{}04.9$ | 0.227 | $4.20\pm 0.26$ | $24.79\pm 5.76$ | $9.03\pm 2.10$ | $240.39\pm 11.67$ | $49.76\pm 12.29$ | $2493.04\pm 152.17$ | $2635.42\pm 160.86$
NGC 5457 Enuc. 3 d | $14~{}03~{}42.91$ | $+54~{}19~{}24.7$ | 0.227 | $0.26\pm 0.05$ | $13.02\pm 5.29$ | $4.74\pm 1.93$ | $15.56\pm 0.76$ | $12.29\pm 5.19$ | $149.91\pm 26.74$ | $159.98\pm 28.53$
NGC 5457 Enuc. 4 a | $14~{}03~{}52.04$ | $+54~{}21~{}52.5$ | 0.227 | $0.14\pm 0.02$ | $7.26\pm 2.63$ | $2.64\pm 0.96$ | $6.67\pm 0.61$ | $9.07\pm 3.43$ | $80.65\pm 11.53$ | $86.81\pm 12.41$
NGC 5457 Enuc. 4 b | $14~{}03~{}53.00$ | $+54~{}21~{}57.3$ | 0.227 | $0.19\pm 0.02$ | $4.94\pm 1.63$ | $1.80\pm 0.59$ | $9.11\pm 0.74$ | $10.71\pm 3.71$ | $113.40\pm 13.58$ | $121.41\pm 14.54$
NGC 5457 Enuc. 4 c | $14~{}03~{}53.20$ | $+54~{}22~{}06.3$ | 0.227 | $0.37\pm 0.03$ | $10.91\pm 2.89$ | $3.97\pm 1.05$ | $20.07\pm 1.04$ | $14.72\pm 4.12$ | $216.24\pm 16.96$ | $230.06\pm 18.05$
NGC 5457 Enuc. 4 d | $14~{}03~{}53.99$ | $+54~{}22~{}10.8$ | 0.227 | $0.30\pm 0.03$ | $14.42\pm 3.96$ | $5.25\pm 1.44$ | $14.87\pm 0.84$ | $13.27\pm 3.85$ | $175.25\pm 14.73$ | $186.76\pm 15.70$
NGC 5713 Enuc. 2 a | $14~{}40~{}10.80$ | $-00~{}17~{}35.5$ | 0.726 | $0.32\pm 0.03$ | $22.21\pm 6.47$ | $8.09\pm 2.36$ | $24.11\pm 1.15$ | $7.65\pm 2.35$ | $185.98\pm 17.52$ | $198.09\pm 18.66$
NGC 5713 Enuc. 2 b | $14~{}40~{}10.86$ | $-00~{}17~{}50.2$ | 0.726 | $0.16\pm 0.02$ | $15.12\pm 5.52$ | $5.51\pm 2.01$ | $12.88\pm 0.63$ | $5.51\pm 2.10$ | $95.48\pm 13.86$ | $102.49\pm 14.88$
NGC 5713 | $14~{}40~{}11.30$ | $-00~{}17~{}27.0$ | 0.726 | $0.38\pm 0.04$ | $31.66\pm 9.35$ | $11.53\pm 3.41$ | $22.29\pm 1.07$ | $8.40\pm 2.61$ | $224.80\pm 21.77$ | $239.10\pm 23.16$
NGC 5713 Enuc. 1 | $14~{}40~{}11.36$ | $-00~{}17~{}18.2$ | 0.726 | $1.24\pm 0.12$ | $66.19\pm 19.22$ | $24.11\pm 7.00$ | $142.99\pm 7.08$ | $15.11\pm 4.62$ | $732.27\pm 68.58$ | $775.23\pm 72.60$
NGC 6946 Enuc. 4 c | $20~{}34~{}22.74$ | $+60~{}10~{}34.2$ | 0.231 | $0.77\pm 0.05$ | $3.79\pm 0.92$ | $1.38\pm 0.33$ | $47.20\pm 2.36$ | $21.16\pm 5.45$ | $455.66\pm 30.23$ | $483.00\pm 32.04$
NGC 6946 Enuc. 8 | $20~{}34~{}32.28$ | $+60~{}10~{}19.3$ | 0.231 | $0.69\pm 0.05$ | $22.58\pm 5.70$ | $8.23\pm 2.08$ | $30.70\pm 1.53$ | $19.98\pm 5.34$ | $406.02\pm 29.05$ | $430.55\pm 30.81$
NGC 6946 Enuc. 5 b | $20~{}34~{}39.36$ | $+60~{}04~{}52.4$ | 0.231 | $0.16\pm 0.02$ | $2.35\pm 0.83$ | $0.85\pm 0.30$ | $4.37\pm 0.21$ | $9.58\pm 3.55$ | $91.68\pm 12.57$ | $98.47\pm 13.50$
NGC 6946 Enuc. 3 a | $20~{}34~{}49.87$ | $+60~{}12~{}40.7$ | 0.231 | $0.04\pm 0.02$ | $1.88\pm 1.10$ | $0.68\pm 0.40$ | $2.11\pm 0.10$ | $5.05\pm 3.04$ | $23.94\pm 8.65$ | $26.90\pm 9.72$
NGC 6946 Enuc. 3 b | $20~{}34~{}52.24$ | $+60~{}12~{}43.7$ | 0.231 | $0.22\pm 0.02$ | $2.54\pm 0.82$ | $0.93\pm 0.30$ | $8.93\pm 0.43$ | $11.17\pm 3.77$ | $125.46\pm 14.31$ | $134.15\pm 15.31$
NGC 6946 Enuc. 6 b | $20~{}35~{}06.96$ | $+60~{}10~{}46.5$ | 0.231 | $0.22\pm 0.03$ | $4.22\pm 1.47$ | $1.54\pm 0.54$ | $16.59\pm 0.82$ | $11.22\pm 4.09$ | $126.52\pm 16.81$ | $135.27\pm 17.97$
NGC 6946 Enuc. 9 | $20~{}35~{}11.09$ | $+60~{}08~{}57.5$ | 0.231 | $0.95\pm 0.06$ | $4.22\pm 1.02$ | $1.54\pm 0.37$ | $45.00\pm 2.24$ | $23.49\pm 6.03$ | $562.30\pm 37.05$ | $595.66\pm 39.25$
NGC 6946 Enuc. 7 | $20~{}35~{}12.97$ | $+60~{}08~{}50.5$ | 0.231 | $0.29\pm 0.03$ | $2.72\pm 0.81$ | $0.99\pm 0.29$ | $27.76\pm 1.38$ | $12.95\pm 4.05$ | $169.41\pm 16.58$ | $180.58\pm 17.67$
NGC 6946 Enuc. 1 | $20~{}35~{}16.80$ | $+60~{}11~{}00.0$ | 0.231 | $0.38\pm 0.03$ | $2.94\pm 0.80$ | $1.07\pm 0.29$ | $18.42\pm 0.91$ | $14.87\pm 4.30$ | $223.95\pm 18.69$ | $238.20\pm 19.88$
NGC 7793 Enuc. 1 | $23~{}57~{}48.80$ | $-32~{}36~{}58.0$ | 0.133 | $0.02\pm 0.02$ | $4.80\pm 4.76$ | $1.75\pm 1.73$ | $6.21\pm 0.64$ | $4.34\pm 4.31$ | $8.90\pm 8.77$ | $11.02\pm 10.85$
NGC 7793 Enuc. 3 | $23~{}57~{}48.80$ | $-32~{}34~{}52.0$ | 0.133 | $0.10\pm 0.03$ | $3.16\pm 1.76$ | $1.15\pm 0.64$ | $13.31\pm 0.70$ | $10.21\pm 5.83$ | $59.08\pm 19.29$ | $64.02\pm 20.90$
GOALS Nuclear Regions
MCG-02-01-051 | $00~{}18~{}50.87$ | $-10~{}22~{}36.5$ | 2.05 | $4.81\pm 0.77$ | $62.09\pm 23.06$ | $62.09\pm 23.06$ | $237.73\pm 237.73$ | $20.47\pm 20.47$ | $2854.21\pm 459.43$ | $3016.98\pm 485.63$
IC 1623 | $01~{}07~{}47.51$ | $-17~{}30~{}25.2$ | 1.52 | $11.45\pm 0.47$ | $226.37\pm 66.49$ | $226.37\pm 66.49$ | $759.33\pm 759.33$ | $36.63\pm 36.63$ | $6792.80\pm 280.13$ | $7177.97\pm 296.01$
MCG-03-04-014 | $01~{}10~{}08.96$ | $-16~{}51~{}9.8$ | 2.64 | $3.05\pm 0.08$ | $131.44\pm 55.47$ | $131.44\pm 55.47$ | $293.36\pm 293.36$ | $14.36\pm 14.36$ | $1807.11\pm 47.65$ | $1910.76\pm 50.38$
NGC 0838 | $02~{}09~{}38.63$ | $-10~{}08~{}47.7$ | 0.972 | $4.21\pm 0.22$ | $107.09\pm 24.51$ | $107.09\pm 24.51$ | $331.20\pm 331.20$ | $27.82\pm 27.82$ | $2498.49\pm 132.46$ | $2641.18\pm 140.02$
IC 0214 | $02~{}14~{}05.52$ | $+05~{}10~{}23.8$ | 2.27 | $2.22\pm 0.16$ | $83.20\pm 25.85$ | $83.20\pm 25.85$ | $94.01\pm 94.01$ | $13.19\pm 13.19$ | $1312.98\pm 95.27$ | $1388.72\pm 100.77$
NGC 0877 | $02~{}18~{}00.05$ | $+14~{}32~{}38.9$ | 0.981 | $0.09\pm 0.18$ | $8.37\pm 7.26$ | $8.37\pm 7.26$ | $31.65\pm 31.65$ | $4.07\pm 4.07$ | $51.74\pm 102.07$ | $56.27\pm 111.01$
UGC 02369 | $02~{}54~{}01.81$ | $+14~{}58~{}14.9$ | 2.36 | $4.04\pm 0.08$ | $169.70\pm 61.47$ | $169.70\pm 61.47$ | $189.46\pm 189.46$ | $17.46\pm 17.46$ | $2394.12\pm 47.18$ | $2530.91\pm 49.88$
CGCG 465-012 | $03~{}54~{}15.96$ | $+15~{}55~{}43.3$ | 1.68 | $3.39\pm 0.07$ | $86.76\pm 22.08$ | $86.76\pm 22.08$ | $133.68\pm 133.68$ | $18.97\pm 18.97$ | $2011.22\pm 42.43$ | $2126.39\pm 44.86$
UGC 02982 | $04~{}12~{}22.57$ | $+05~{}32~{}50.3$ | 1.31 | $1.54\pm 0.10$ | $80.61\pm 29.00$ | $80.61\pm 29.00$ | $167.73\pm 167.73$ | $14.48\pm 14.48$ | $911.12\pm 56.86$ | $964.17\pm 60.18$
UGC 03094 | $04~{}35~{}33.81$ | $+19~{}10~{}17.7$ | 1.99 | $2.84\pm 0.14$ | $95.80\pm 23.55$ | $95.80\pm 23.55$ | $178.15\pm 178.15$ | $15.95\pm 15.95$ | $1680.50\pm 82.89$ | $1777.00\pm 87.65$
IRAS 05442+173 | $05~{}47~{}11.16$ | $+17~{}33~{}46.7$ | 1.45 | $7.93\pm 0.11$ | $195.59\pm 31.46$ | $195.59\pm 31.46$ | $368.30\pm 368.30$ | $31.20\pm 31.20$ | $4701.50\pm 63.36$ | $4968.58\pm 66.96$
IC 0563 | $09~{}46~{}20.30$ | $+03~{}02~{}44.9$ | 1.69 | $0.39\pm 0.15$ | $43.75\pm 19.56$ | $43.75\pm 19.56$ | $43.91\pm 43.91$ | $6.45\pm 6.45$ | $231.41\pm 87.35$ | $246.09\pm 92.89$
NGC 3110 | $10~{}04~{}02.10$ | $-06~{}28~{}28.9$ | 1.42 | $2.27\pm 0.11$ | $93.69\pm 26.42$ | $93.69\pm 26.42$ | $183.44\pm 183.44$ | $16.87\pm 16.87$ | $1345.16\pm 63.12$ | $1422.72\pm 66.76$
IC 2810 | $11~{}25~{}45.07$ | $+14~{}40~{}35.7$ | 2.77 | $2.96\pm 0.05$ | $116.65\pm 22.19$ | $116.65\pm 22.19$ | $92.66\pm 92.66$ | $13.81\pm 13.81$ | $1753.96\pm 30.78$ | $1854.60\pm 32.54$
NGC 5257 | $13~{}39~{}52.92$ | $+00~{}50~{}24.5$ | 1.91 | $0.79\pm 0.07$ | $46.09\pm 11.83$ | $46.09\pm 11.83$ | $39.98\pm 39.98$ | $8.58\pm 8.58$ | $465.30\pm 44.15$ | $493.18\pm 46.80$
NGC 5258 | $13~{}39~{}57.85$ | $+00~{}49~{}54.1$ | 1.91 | $0.33\pm 0.10$ | $18.44\pm 9.12$ | $18.44\pm 9.12$ | $28.92\pm 28.92$ | $5.52\pm 5.52$ | $191.27\pm 57.54$ | $203.68\pm 61.28$
NGC 5331 | $13~{}52~{}16.20$ | $+02~{}06~{}05.5$ | 2.70 | $3.30\pm 0.11$ | $125.39\pm 33.38$ | $125.39\pm 33.38$ | $175.42\pm 175.42$ | $14.78\pm 14.78$ | $1957.60\pm 65.34$ | $2069.74\pm 69.08$
NGC 5990 | $15~{}46~{}16.37$ | $+02~{}24~{}55.7$ | 1.08 | $2.99\pm 0.06$ | $114.29\pm 24.02$ | $114.29\pm 24.02$ | $304.09\pm 304.09$ | $22.25\pm 22.25$ | $1772.87\pm 37.26$ | $1874.58\pm 39.40$
IRASF 16516-09 | $16~{}54~{}23.74$ | $-09~{}53~{}20.7$ | 1.88 | $2.41\pm 0.22$ | $71.94\pm 20.04$ | $71.94\pm 20.04$ | $121.18\pm 121.18$ | $15.12\pm 15.12$ | $1426.18\pm 130.23$ | $1508.32\pm 137.73$
IRASF 17138-10 | $17~{}16~{}35.81$ | $-10~{}20~{}39.5$ | 1.47 | $6.43\pm 0.15$ | $152.54\pm 26.33$ | $152.54\pm 26.33$ | $240.59\pm 240.59$ | $27.94\pm 27.94$ | $3814.59\pm 88.51$ | $4031.59\pm 93.55$
NGC 7592 | $23~{}18~{}21.81$ | $-04~{}24~{}56.9$ | 1.84 | $5.14\pm 0.10$ | $141.63\pm 19.77$ | $141.63\pm 19.77$ | $138.60\pm 138.60$ | $22.31\pm 22.31$ | $3050.39\pm 62.00$ | $3224.23\pm 65.53$
NGC 7679 | $23~{}28~{}46.67$ | $+03~{}30~{}40.9$ | 1.31 | $1.97\pm 0.30$ | $60.82\pm 20.95$ | $60.82\pm 20.95$ | $205.84\pm 205.84$ | $16.38\pm 16.38$ | $1169.23\pm 179.43$ | $1236.86\pm 189.81$
GOALS Extranuclear Regions
IC 1623 Enuc. 1 | $01~{}07~{}46.70$ | $-17~{}30~{}20.6$ | 1.52 | $1.48\pm 0.50$ | $58.15\pm 31.84$ | $21.18\pm 11.60$ | $55.71\pm 55.71$ | $13.16\pm 13.16$ | $874.87\pm 298.08$ | $925.87\pm 315.46$
IC 1623 Enuc. 2 | $01~{}07~{}46.74$ | $-17~{}30~{}27.2$ | 1.52 | $3.06\pm 0.34$ | $107.28\pm 47.03$ | $39.08\pm 17.13$ | $116.52\pm 116.52$ | $18.95\pm 18.95$ | $1815.96\pm 201.97$ | $1920.10\pm 213.55$
IC 1623 Enuc. 3 | $01~{}07~{}47.07$ | $-17~{}30~{}26.6$ | 1.52 | $2.69\pm 0.30$ | $101.53\pm 44.33$ | $35.54\pm 15.52$ | $102.87\pm 102.87$ | $17.76\pm 17.76$ | $1594.69\pm 178.67$ | $1686.34\pm 188.94$
MCG-03-04-014 Enuc. 1 | $01~{}10~{}09.06$ | $-16~{}51~{}12.7$ | 2.64 | $1.52\pm 0.07$ | $79.90\pm 31.19$ | $27.01\pm 10.54$ | $126.94\pm 126.94$ | $10.14\pm 10.14$ | $900.69\pm 44.31$ | $953.15\pm 46.89$
NGC 0838 Enuc. 1 | $02~{}09~{}38.31$ | $-10~{}08~{}49.4$ | 0.972 | $1.85\pm 0.36$ | $44.27\pm 15.95$ | $16.13\pm 5.81$ | $158.33\pm 158.33$ | $18.46\pm 18.46$ | $1098.44\pm 211.23$ | $1162.07\pm 223.46$
IC 0214 Enuc. 3 | $02~{}14~{}05.02$ | $+05~{}10~{}29.9$ | 2.27 | $1.25\pm 0.19$ | $53.20\pm 22.28$ | $19.38\pm 8.12$ | $28.69\pm 28.69$ | $9.89\pm 9.89$ | $737.10\pm 110.69$ | $780.33\pm 117.18$
IC 0214 Enuc. 2 | $02~{}14~{}05.31$ | $+05~{}10~{}28.7$ | 2.27 | $1.29\pm 0.17$ | $63.44\pm 24.76$ | $23.11\pm 9.02$ | $60.11\pm 60.11$ | $10.06\pm 10.06$ | $762.68\pm 98.95$ | $807.35\pm 104.75$
NGC 0877 Enuc. 2 | $02~{}18~{}00.27$ | $+14~{}32~{}27.1$ | 0.981 | $0.08\pm 0.12$ | $18.14\pm 11.17$ | $6.61\pm 4.07$ | $37.16\pm 37.16$ | $3.74\pm 3.74$ | $43.54\pm 69.23$ | $47.60\pm 75.70$
NGC 0877 Enuc. 1 | $02~{}18~{}00.42$ | $+14~{}32~{}44.7$ | 0.981 | $0.49\pm 0.14$ | $16.75\pm 7.98$ | $6.10\pm 2.91$ | $26.32\pm 26.32$ | $9.44\pm 9.44$ | $288.51\pm 81.38$ | $306.41\pm 86.43$
UGC 02369 Enuc. 1 | $02~{}54~{}01.90$ | $+14~{}58~{}11.5$ | 2.36 | $1.67\pm 0.10$ | $82.00\pm 29.44$ | $29.87\pm 10.72$ | $65.07\pm 65.07$ | $11.23\pm 11.23$ | $989.91\pm 57.10$ | $1047.42\pm 60.42$
CGCG 465-012 Enuc. 2 | $03~{}54~{}16.15$ | $+15~{}55~{}35.1$ | 1.68 | $0.67\pm 0.08$ | $30.62\pm 18.07$ | $11.16\pm 6.58$ | $28.86\pm 28.86$ | $8.44\pm 8.44$ | $396.07\pm 48.75$ | $420.04\pm 51.70$
CGCG 465-012 Enuc. 1 | $03~{}54~{}16.19$ | $+15~{}55~{}47.5$ | 1.68 | $2.11\pm 0.07$ | $59.94\pm 20.63$ | $21.84\pm 7.51$ | $74.02\pm 74.02$ | $14.94\pm 14.94$ | $1247.63\pm 43.77$ | $1319.68\pm 46.30$
CGCG 465-012 Enuc. 3 | $03~{}54~{}16.57$ | $+15~{}55~{}37.3$ | 1.68 | $0.44\pm 0.10$ | $21.34\pm 15.52$ | $7.78\pm 5.66$ | $17.20\pm 17.20$ | $6.83\pm 6.83$ | $258.88\pm 57.31$ | $275.11\pm 60.90$
UGC 02982 Enuc. 5 | $04~{}12~{}22.09$ | $+05~{}32~{}59.1$ | 1.31 | $0.41\pm 0.09$ | $38.91\pm 33.05$ | $14.17\pm 12.04$ | $22.80\pm 22.80$ | $7.45\pm 7.45$ | $239.30\pm 52.56$ | $254.42\pm 55.89$
UGC 02982 Enuc. 2 | $04~{}12~{}22.15$ | $+05~{}32~{}48.3$ | 1.31 | $0.44\pm 0.14$ | $35.48\pm 18.92$ | $12.93\pm 6.89$ | $42.12\pm 42.12$ | $7.76\pm 7.76$ | $260.02\pm 83.70$ | $276.31\pm 88.94$
UGC 02982 Enuc. 3 | $04~{}12~{}22.41$ | $+05~{}32~{}57.9$ | 1.31 | $0.50\pm 0.08$ | $47.48\pm 56.66$ | $17.30\pm 20.64$ | $35.26\pm 35.26$ | $8.26\pm 8.26$ | $294.92\pm 46.31$ | $313.18\pm 49.17$
UGC 02982 Enuc. 4 | $04~{}12~{}23.74$ | $+05~{}32~{}47.1$ | 1.31 | $0.33\pm 0.13$ | $29.55\pm 22.71$ | $10.77\pm 8.27$ | $26.88\pm 26.88$ | $6.72\pm 6.72$ | $194.38\pm 74.81$ | $206.96\pm 79.65$
UGC 03094 Enuc. 2 | $04~{}35~{}33.91$ | $+19~{}10~{}10.1$ | 1.99 | $0.79\pm 0.39$ | $7.96\pm 4.64$ | $2.90\pm 1.69$ | $65.18\pm 65.18$ | $8.42\pm 8.42$ | $466.36\pm 232.62$ | $494.30\pm 246.56$
IRAS05442+173 Enuc. 1 | $05~{}47~{}10.83$ | $+17~{}33~{}46.3$ | 1.45 | $0.77\pm 0.12$ | $38.72\pm 18.26$ | $12.28\pm 5.79$ | $26.46\pm 26.46$ | $9.70\pm 9.70$ | $452.48\pm 71.82$ | $479.64\pm 76.13$
IC 0563 Enuc. 2 | $09~{}46~{}20.07$ | $+03~{}02~{}43.7$ | 1.69 | $0.23\pm 0.17$ | $22.44\pm 16.84$ | $8.18\pm 6.14$ | $24.12\pm 24.12$ | $4.93\pm 4.93$ | $134.30\pm 97.03$ | $143.49\pm 103.67$
NGC 3110 Enuc. 3 | $10~{}04~{}01.53$ | $-06~{}28~{}26.0$ | 1.43 | $0.78\pm 0.20$ | $32.87\pm 15.33$ | $11.98\pm 5.58$ | $55.79\pm 55.79$ | $9.87\pm 9.87$ | $458.87\pm 115.56$ | $486.38\pm 122.49$
NGC 3110 Enuc. 2 | $10~{}04~{}02.03$ | $-06~{}28~{}33.1$ | 1.42 | $1.08\pm 0.12$ | $67.12\pm 23.60$ | $24.45\pm 8.60$ | $124.85\pm 124.85$ | $11.62\pm 11.62$ | $637.58\pm 69.95$ | $675.19\pm 74.08$
NGC 3110 Enuc. 1 | $10~{}04~{}02.58$ | $-06~{}28~{}46.4$ | 1.42 | $0.70\pm 0.13$ | $40.48\pm 20.96$ | $14.75\pm 7.63$ | $41.60\pm 41.60$ | $9.39\pm 9.39$ | $415.31\pm 77.59$ | $440.36\pm 82.27$
NGC 3110 Enuc. 4 | $10~{}04~{}02.69$ | $-06~{}28~{}35.1$ | 1.42 | $0.46\pm 0.14$ | $28.62\pm 19.91$ | $10.43\pm 7.25$ | $37.83\pm 37.83$ | $7.56\pm 7.56$ | $268.30\pm 81.52$ | $285.05\pm 86.61$
IC 2810 Enuc. 1 | $11~{}25~{}44.92$ | $+14~{}40~{}30.9$ | 2.77 | $1.22\pm 0.16$ | $6.94\pm 2.25$ | $2.53\pm 0.82$ | $20.67\pm 20.67$ | $8.88\pm 8.88$ | $723.57\pm 97.31$ | $766.04\pm 103.02$
NGC 5257 Enuc. 3 | $13~{}39~{}52.19$ | $+00~{}50~{}22.3$ | 1.91 | $0.38\pm 0.06$ | $29.69\pm 13.19$ | $10.82\pm 4.80$ | $14.96\pm 14.96$ | $5.98\pm 5.98$ | $225.47\pm 37.31$ | $239.81\pm 39.69$
NGC 5257 Enuc. 2 | $13~{}39~{}52.59$ | $+00~{}50~{}15.5$ | 1.91 | $0.58\pm 0.06$ | $38.95\pm 13.39$ | $14.19\pm 4.88$ | $15.82\pm 15.82$ | $7.35\pm 7.35$ | $341.30\pm 37.77$ | $362.19\pm 40.08$
NGC 5257 Enuc. 1 | $13~{}39~{}52.94$ | $+00~{}50~{}12.6$ | 1.91 | $1.22\pm 0.12$ | $24.83\pm 6.29$ | $9.04\pm 2.29$ | $21.24\pm 21.24$ | $10.68\pm 10.68$ | $723.00\pm 72.90$ | $765.43\pm 77.18$
NGC 5257 Enuc. 4 | $13~{}39~{}53.54$ | $+00~{}50~{}28.9$ | 1.91 | $0.31\pm 0.07$ | $27.40\pm 13.50$ | $9.98\pm 4.92$ | $24.25\pm 24.25$ | $5.39\pm 5.39$ | $182.68\pm 38.57$ | $194.60\pm 41.09$
NGC 5258 Enuc. 2 | $13~{}39~{}57.10$ | $+00~{}49~{}40.5$ | 1.91 | $0.69\pm 0.08$ | $36.97\pm 11.77$ | $13.47\pm 4.29$ | $23.31\pm 23.31$ | $8.03\pm 8.03$ | $407.72\pm 44.55$ | $432.36\pm 47.24$
NGC 5258 Enuc. 3 | $13~{}39~{}57.12$ | $+00~{}49~{}44.0$ | 1.91 | $1.49\pm 0.07$ | $58.72\pm 12.33$ | $21.39\pm 4.49$ | $60.52\pm 60.52$ | $11.79\pm 11.79$ | $882.28\pm 43.15$ | $933.71\pm 45.66$
NGC 5258 Enuc. 1 | $13~{}39~{}57.23$ | $+00~{}49~{}47.4$ | 1.91 | $1.16\pm 0.05$ | $66.80\pm 18.41$ | $24.34\pm 6.71$ | $84.49\pm 84.49$ | $10.41\pm 10.41$ | $686.89\pm 30.29$ | $727.28\pm 32.07$
NGC 5331 Enuc. 2 | $13~{}52~{}16.03$ | $+02~{}06~{}05.3$ | 2.70 | $1.33\pm 0.08$ | $70.57\pm 21.63$ | $25.71\pm 7.88$ | $66.29\pm 66.29$ | $9.38\pm 9.38$ | $787.34\pm 46.61$ | $833.41\pm 49.34$
NGC 5331 Enuc. 4 | $13~{}52~{}16.21$ | $+02~{}06~{}08.5$ | 2.70 | $0.99\pm 0.09$ | $55.45\pm 18.22$ | $20.20\pm 6.64$ | $52.36\pm 52.36$ | $8.09\pm 8.09$ | $584.83\pm 50.31$ | $619.46\pm 53.29$
NGC 5331 Enuc. 1 | $13~{}52~{}16.30$ | $+02~{}05~{}59.1$ | 2.70 | $0.28\pm 0.07$ | $35.85\pm 29.95$ | $13.06\pm 10.91$ | $28.54\pm 28.54$ | $4.32\pm 4.32$ | $165.52\pm 38.77$ | $176.47\pm 41.34$
NGC 5331 Enuc. 3 | $13~{}52~{}16.32$ | $+02~{}06~{}05.7$ | 2.70 | $1.46\pm 0.07$ | $86.55\pm 30.82$ | $31.53\pm 11.23$ | $98.98\pm 98.98$ | $9.84\pm 9.84$ | $866.63\pm 39.86$ | $917.17\pm 42.18$
NGC 5990 Enuc. 2 | $15~{}46~{}16.08$ | $+02~{}25~{}02.1$ | 1.08 | $0.55\pm 0.09$ | $35.12\pm 14.89$ | $12.80\pm 5.42$ | $60.68\pm 60.68$ | $9.58\pm 9.58$ | $326.45\pm 51.37$ | $346.49\pm 54.52$
NGC 5990 Enuc. 1 | $15~{}46~{}16.53$ | $+02~{}24~{}47.7$ | 1.08 | $1.13\pm 0.09$ | $45.96\pm 14.02$ | $16.74\pm 5.11$ | $79.55\pm 79.55$ | $13.65\pm 13.65$ | $665.93\pm 54.48$ | $705.14\pm 57.69$
IRASF16516-09 Enuc. 4 | $16~{}54~{}23.40$ | $-09~{}53~{}22.2$ | 1.88 | $0.64\pm 0.20$ | $24.92\pm 11.68$ | $9.08\pm 4.25$ | $31.25\pm 31.25$ | $7.77\pm 7.77$ | $374.98\pm 117.33$ | $397.76\pm 124.46$
IRASF16516-09 Enuc. 1 | $16~{}54~{}23.81$ | $-09~{}53~{}30.6$ | 1.88 | $0.55\pm 0.22$ | $19.02\pm 10.11$ | $6.93\pm 3.68$ | $22.80\pm 22.80$ | $7.25\pm 7.25$ | $326.02\pm 130.43$ | $346.03\pm 138.44$
IRASF17138-10 Enuc. 1 | $17~{}16~{}35.80$ | $-10~{}20~{}41.9$ | 1.47 | $6.24\pm 0.16$ | $139.47\pm 24.25$ | $50.81\pm 8.83$ | $240.76\pm 240.76$ | $27.52\pm 27.52$ | $3700.53\pm 95.29$ | $3911.09\pm 100.71$
NGC 7592 Enuc. 1 | $23~{}18~{}22.18$ | $-04~{}25~{}08.0$ | 1.84 | $0.13\pm 0.19$ | $12.13\pm 10.42$ | $4.42\pm 3.79$ | $9.80\pm 9.80$ | $3.56\pm 3.56$ | $75.50\pm 110.40$ | $81.37\pm 118.98$
NGC 7679 Enuc. 1 | $23~{}28~{}46.50$ | $+03~{}30~{}43.7$ | 1.31 | $2.04\pm 0.10$ | $63.99\pm 13.10$ | $23.31\pm 4.77$ | $149.96\pm 149.96$ | $16.65\pm 16.65$ | $1207.80\pm 61.22$ | $1277.61\pm 64.76$
NGC 7679 Enuc. 2 | $23~{}28~{}46.90$ | $+03~{}30~{}41.3$ | 1.31 | $1.27\pm 0.12$ | $42.57\pm 11.92$ | $15.51\pm 4.34$ | $103.72\pm 103.72$ | $13.16\pm 13.16$ | $753.55\pm 68.15$ | $797.71\pm 72.15$

Note. — † Pressure estimates derived using empirical relations between $P_{\rm
DE}$, $P_{\rm turb}$ and $\Sigma_{\rm SFR}$ given in Sun et al. (2020).
Associated uncertainties provided here do not include the intrinsic scatter of
these empirical relations.