# 183 GHz water megamasers in active galactic nuclei: a new accretion disk
tracer

Dominic W. Pesce Center for Astrophysics $|$ Harvard & Smithsonian, 60 Garden
Street, Cambridge, MA 02138, USA  Black Hole Initiative at Harvard University,
20 Garden Street, Cambridge, MA 02138, USA  James A. Braatz National Radio
Astronomy Observatory, 520 Edgemont Road, Charlottesville, VA 22903, USA
Christian Henkel Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69,
D-53121 Bonn, Germany  Astronomy Department, Faculty of Science, King
Abdulaziz University, P.O. Box 80203, Jeddah 21589, Saudi Arabia  Elizabeth M.
L. Humphreys Joint ALMA Observatory, Alonso de Cordova 3107, Vitacura,
Santiago, Chile  European Southern Observatory (ESO) Vitacura, Alonso de
Cordova 3107, Vitacura, Santiago, Chile  C. M. Violette Impellizzeri Leiden
Observatory, Leiden University, PO Box 9513, 2300 RA Leiden, The Netherlands
National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, VA
22903, USA  Cheng-Yu Kuo Institute of Astronomy and Astrophysics, Academia
Sinica, 11F of Astronomy-Mathematics Building, AS/NTU No. 1, Sec. 4, Roosevelt
Rd, Taipei 10617, Taiwan, R.O.C.  Physics Department, National Sun Yat-Sen
University, No. 70, Lien-Hai Road, Kaosiung City 80424, Taiwan, R.O.C.

###### Abstract

We present the results of an ALMA survey to identify 183 GHz H2O maser
emission from AGN already known to host 22 GHz megamaser systems. Out of 20
sources observed, we detect significant 183 GHz maser emission from 13; this
survey thus increases the number of AGN known to host (sub)millimeter
megamasers by a factor of 5. We find that the 183 GHz emission is
systematically fainter than the 22 GHz emission from the same targets, with
typical flux densities being roughly an order of magnitude lower at 183 GHz
than at 22 GHz. However, the isotropic luminosities of the detected 183 GHz
sources are comparable to their 22 GHz values. For two of our sources – ESO
269-G012 and the Circinus galaxy – we detect rich 183 GHz spectral structure
containing multiple line complexes. The 183 GHz spectrum of ESO 269-G012
exhibits the triple-peaked structure characteristic of an edge-on AGN disk
system. The Circinus galaxy contains the strongest 183 GHz emission detected
in our sample, peaking at a flux density of nearly 5 Jy. The high signal-to-
noise ratios achieved by these strong lines enable a coarse mapping of the 183
GHz maser system, in which the masers appear to be distributed similarly to
those seen in VLBI maps of the 22 GHz system in the same galaxy and may be
tracing the circumnuclear accretion disk at larger orbital radii than are
occupied by the 22 GHz masers. This newly identified population of AGN disk
megamasers presents a motivation for developing VLBI capabilities at 183 GHz.

††facilities: ALMA††software: astropy (Astropy Collaboration et al., 2013,
2018, 2022), CASA (McMullin et al., 2007), dynesty (Speagle, 2020), matplotlib
(Hunter, 2007), numpy (Harris et al., 2020), scipy (Virtanen et al., 2020)

##  1 Introduction

H2O megamasers residing in the accretion disks of active galactic nuclei (AGN)
have proven to be unique tools for mapping circumnuclear molecular gas on sub-
pc scales (e.g., Greenhill et al., 1995; Argon et al., 2007; Reid et al.,
2009), providing precise constraints on supermassive black hole (SMBH) masses
(e.g., Kuo et al., 2011; Gao et al., 2017; Zhao et al., 2018), and enabling
geometric distance measurements to their host galaxies (e.g., Herrnstein et
al., 1999; Braatz et al., 2010; Humphreys et al., 2013; Kuo et al., 2013; Gao
et al., 2016; Reid et al., 2019; Pesce et al., 2020a). For galaxies situated
in the Hubble flow, these distance measurements can be used to provide direct
constraints on the local cosmological expansion rate (Reid et al., 2013; Kuo
et al., 2015; Pesce et al., 2020b).

To date, most of the observational work on AGN disk megamasers has focused on
the $6_{1,6}$-$5_{2,3}$ rotational transition in the ground vibrational state
of the ortho-H2O molecule (i.e., the version of the molecule in which the
nuclear spins of the two hydrogen atoms are parallel). This transition emits
at a rest-frame frequency of $\sim$22.23508 GHz111In astrophysical systems,
the observed $6_{1,6}$-$5_{2,3}$ line is typically composed of a blend of six
hyperfine transitions spanning $\sim$0.5 MHz; the 22.23508 GHz nominal rest-
frame frequency represents an intensity-weighted (assuming thermal
equilibrium) mean of these six transitions. (Kukolich, 1969), though we adopt
the community standard and refer to it in this paper as simply “the 22 GHz
line.” Nearly 200 galaxies have been detected in this line so far – the result
of more than 4800 galaxies surveyed (Kuo et al., 2018, 2020) – and $\sim$150
of these detections are associated with
AGN.222https://safe.nrao.edu/wiki/bin/view/Main/PublicWaterMaserList

A $\sim$20% subset of these AGN megamasers originate from circumnuclear
accretion disks and are identifiable by their characteristic spectral profiles
containing three distinct groups of maser features distributed approximately
symmetrically around the recession velocity of the host galaxy (see, e.g.,
Pesce et al., 2015). The central group of features (called the “systemic
masers”) coincides roughly with the recession velocity of the galaxy; these
features originate from masing gas that resides along our line of sight
through the disk to the central AGN. The two other groups of features – which
are typically offset by several hundred km s-1 redward and blueward of the
galaxy recession velocity – are collectively known as the “high-velocity”
features and originate from masing gas on lines of sight through the disk that
are tangent to its orbital motion.

Though the 22 GHz megamaser transition is by far the best studied, the H2O
molecule has a rich energy spectrum containing many other transitions capable
of sustaining maser activity under physical conditions comparable to those
that support 22 GHz emission (Neufeld & Melnick, 1991; Yates et al., 1997;
Gray et al., 2016). To date, there have been only three AGN observed to host
H2O megamaser emission at frequencies other than 22 GHz:

  * •

The first detection of an extragalactic H2O maser source in a transition other
than the 22 GHz line was made by Humphreys et al. (2005), who used the
Submillimeter Array (SMA) to observe 183 GHz megamaser emission from NGC 3079.
The 183 GHz emission from NGC 3079 is consistent with an origin near the AGN
and has a flux density about $\sim$3 times fainter than the corresponding 22
GHz emission; a tentative detection of 439 GHz maser emission with the James
Clerk Maxwell Telescope (JCMT) towards the same galaxy was also identified by
the authors.

  * •

Hagiwara et al. (2013) used the Atacama Large Millimeter/submillimeter Array
(ALMA) to detect 321 GHz H2O megamaser emission towards the Circinus galaxy,
where it is observed to be roughly two orders of magnitude fainter than the
corresponding 22 GHz emission. The spectral structure at 321 GHz is similar to
that at 22 GHz, and repeated observations have demonstrated that it persists
over years but exhibits large ($\sim$order of magnitude) variability in the
maser line strength (Pesce et al., 2016; Hagiwara et al., 2021).

  * •

Humphreys et al. (2016) used the Atacama Pathfinder Experiment (APEX)
telescope to detect 183 GHz megamaser emission from NGC 4945, where it is
observed to have a flux density that is a factor of $\sim$1.5 times greater
than that of the corresponding 22 GHz emission. Pesce et al. (2016) and
Hagiwara et al. (2016) additionally identified 321 GHz megamaser emission from
NGC 4945 using ALMA observations, at a level more than two orders of magnitude
fainter than the 22 GHz or 183 GHz emission but with a similar spectral
structure. The 321 GHz maser system appears to be time-variable, as followup
ALMA observations failed to detect any emission from NGC 4945 (Hagiwara et
al., 2021).

The 183 GHz transition has also been detected towards $\sim$10 ultraluminous
infrared galaxies (ULIRGs), where it is typically associated with spatially
distributed star formation rather than centrally concentrated AGN activity
(Cernicharo et al., 2006; König et al., 2017; Imanishi et al., 2022). One
possible exception is the so-called “Superantennae” pair of interacting
galaxies, which are ULIRGs but for which the 183 GHz emission appears to be at
least partially circumnuclear and has been associated with an AGN (Imanishi et
al., 2021).

In this paper we present the results of an ALMA survey that increases by a
factor of 5 the number of AGN known to host H2O megamaser emission at a
frequency other than 22 GHz. Our survey targets the $3_{1,3}$-$2_{2,0}$
rotational transition in the ground vibrational state of the para-H2O molecule
(where the nuclear spins of the two hydrogen atoms are anti-parallel), which
has a rest-frame emission frequency of 183.310087
GHz333https://splatalogue.online/ (Pickett et al., 1998). We refer to this
transition as “the 183 GHz line.”

This paper is organized as follows. In Section 2 we describe the observations,
data reduction, and imaging procedure. Section 3 presents the results of our
survey and compares the 183 GHz spectra to prior 22 GHz observations. We
summarize and conclude in Section 4. Throughout this paper, all velocities are
quoted using the optical convention in the heliocentric reference frame unless
otherwise specified, and we adopt a Hubble constant value of $H_{0}=70$ km s-1
Mpc-1.

##  2 Data reduction and imaging

The data presented in this paper were taken as part of ALMA projects
2017.1.00909.S and 2018.1.00321.S. We used ALMA’s Band 5 to observe 20 AGN
known to host 22 GHz water megamaser systems and spanning a range of redshifts
out to $z\approx 0.05$; the targets are listed in LABEL:tab:Observations. Each
of the targeted galaxies was selected either because it was identified as a
disk maser by Pesce et al. (2015), and/or – as in the case of NGC 1068, NGC
1386, Circinus, and NGC 5643 – because it is a strong 22 GHz megamaser source
associated with an AGN. Observations of the science targets were interwoven
with observations of calibrator targets, which are used to determine bandpass,
flux density, and station gain calibrations. All observations recorded only
dual linear polarization correlation products (XX, YY), and cross-hand data
are not available.

The spectral setup for each observation consisted of four spectral windows.
Three of the spectral windows were configured to have fine frequency
resolution (122 kHz, corresponding to $\sim$0.2 km s-1) and to span a total of
$\pm$ $\sim$1200 km s-1 around the expected locations of the maser lines
(i.e., centered on the recession velocity of each galaxy) with a small
($\sim$50 km s-1) overlap to ensure continuous coverage. The fourth spectral
window was wider (2 GHz, corresponding to $\sim$3200 km s-1) with coarser
channels (15.625 MHz, corresponding to $\sim$25 km s-1), and was offset toward
lower frequency to capture continuum emission.

Though most of the observations taken as part of these ALMA projects have
corresponding quality assurance 2 (QA2) data processing, we also include a
handful of observations that did not achieve sufficient sensitivity to pass
the QA2 level. Additionally, there are some desirable data reduction
procedures (e.g., self-calibration) that are not included in the standard QA2
pipeline. To maintain uniformity, we thus carry out our own data reduction on
all datasets. Our data reduction procedure for each observation largely
follows standard practice for ALMA datasets. We have carried out all data
reduction for this paper using CASA444https://casa.nrao.edu/ version 6.2.1.7.
Where the calibration procedure for individual science targets deviates from
the general procedure described in this section, or where there is relevant
information specific to an individual source, we have provided additional
details in Appendix A.

J0109-0332 | 01:09:45.101 |  $-03$:32:32.97 | 16369 | 2018 Nov 06
---|---|---|---|---
J0126-0417 | 01:26:01.640 |  $-04$:17:56.42 | 5639 | 2018 Nov 16
NGC 1068 | 02:42:40.771 |  $-00$:00:47.84 | 1137 | 2018 Nov 9, Nov 16
NGC 1194 | 03:03:49.100 |  $-01$:06:12.99 | 4086 | 2018 Nov 16
NGC 1386 | 03:36:46.238 |  $-35$:59:57.39 | 868 | 2018 Nov 16
ESO 558-G009 | 07:04:21.002 |  $-21$:35:19.36 | 7674 | 2018 Oct 31
IC 485 | 08:00:19.772 |  $+26$:42:05.35 | 8342 | 2018 Oct 31, Dec 06
J0847-0022 | 08:47:47.692 |  $-00$:22:51.47 | 15243 | 2018 Nov 08
Mrk 1419 | 09:40:36.376 |  $+03$:34:36.96 | 4947 | 2018 Nov 14
IC 2560 | 10:16:18.666 |  $-33$:33:49.85 | 2925 | 2018 Oct 02, Oct 30, Nov 08
NGC 3393 | 10:48:23.467 |  $-25$:09.43.30 | 3750 | 2018 Oct 30, Nov 14
UGC 6093 | 11:00:47.955 |  $+10$.43.41.76 | 10803 | 2018 Oct 31
ESO 269-G012 | 12:56:40.521 |  $-46$:55:33.79 | 5014 | 2018 Dec 01
CGCG 074-064 | 14:03:04.459 |  $+08$:56:51.03 | 6915 | 2018 Dec 06
NGC 5495 | 14:12:23.353 |  $-27$:06:28.72 | 6737 | 2018 Dec 06
Circinus | 14:13:09.950 |  $-65$:20:21.20 | 434 | 2018 Dec 01
NGC 5643 | 14:32:40.778 |  $-44$:10:28.60 | 1199 | 2018 Aug 27, Sep 20
NGC 5765b | 14:50:51.500 |  $+05$:06:52.00 | 8257 | 2018 Dec 06
CGCG 165-035 | 15:14:39.793 |  $+26$:35:39.11 | 9622 | 2018 Dec 06
NGC 6264 | 16:57:16.110 |  $+27$:50:58.71 | 10145 | 2018 Nov 18

Note. — Information about the targets surveyed in this paper for 183 GHz
emission. Coordinates are specified in J2000 and indicate the pointing center
for the science targets; observing dates are specified in UT. Recession
velocities $v_{\text{rec}}$ have been obtained from the NASA/IPAC
Extragalactic Database (NED).

###  2.1 Data reduction procedure

We start by importing the raw ALMA Science Data Model files and converting
them to CASA Measurement Sets (MSs) using the importasdm task. Initial
flagging is carried out using the flagdata task; we remove pointing scans,
autocorrelation products, and any data associated with shadowed antennas. We
use the gencal task to create system temperature and water vapor radiometer
calibration tables, which we then apply to the science and calibrator targets
using the applycal task. We then perform a round of data inspection and manual
flagging, the latter typically restricted to edge channels of the continuum
spectral windows and to the initial integrations of some scans.

Prior to determining the bandpass solution, we first track the phase
variations of the bandpass calibrator on every integration time by averaging
over a small number of frequency channels ($\sim$5% of each spectral window)
using the gaincal task. We select as our reference antenna the one that is
most centrally located within the array configuration while also being present
for all scans, and phase solutions are referenced to the phase at this
antenna. These phase solutions are then applied to the bandpass calibrator
prior to deriving the phase and amplitude bandpass solutions on a per-channel
basis using the bandpass task.

We determine the absolute flux density scale using observations of targets
from the ALMA calibrator source catalog555https://almascience.nrao.edu/sc/.
For each flux density calibrator, we assume a power-law spectral energy
distribution to extrapolate from the ALMA calibrator source catalog
measurements to the $\sim$183 GHz observing frequencies relevant for the
observations in this paper. Specifically, we use

$S_{\nu}=S_{\nu,0}\left(\frac{\nu}{\nu_{0}}\right)^{\alpha},$ (1)

where $S_{\nu}$ is the flux density at frequency $\nu$, $S_{\nu,0}$ is the
flux density at a reference frequency $\nu_{0}$, and $\alpha$ is the spectral
index. We specify the absolute flux density calibration using CASA’s setjy
task. The flux densities and spectral indices assumed for each of our
calibrators are detailed in Appendix A.

We carry out gain calibration in a few steps. Applying bandpass calibration
on-the-fly, we first use the gaincal task to derive gain phase solutions on
the calibrator targets per integration time. We separately solve for the gain
phases on the calibrator targets per scan time, for later application to the
science target. We apply the integration-time gain phase solutions to the
calibrator targets during another run of gaincal in which we solve for both
the gain amplitudes and phases per scan time. Following this gain calibration,
we bootstrap the absolute flux density calibration from the flux density
calibrator to the gain calibrator using the fluxscale task. The gain phase and
amplitude solutions – the latter of which now carry over the absolute flux
density calibration – are then applied to the science target using the
applycal task.

###  2.2 Imaging procedure

We carry out all imaging for this paper using the Högbom CLEAN method (Högbom,
1974) as implemented in CASA’s tclean task, combining the calibrated XX and YY
visibilities to form Stokes I images. We use the default “standard”
$(u,v)$-gridding setting in tclean, and we specify natural visibility
weighting to maximize the point source sensitivity of our images. Each image
is reconstructed on a $300\times 300$ pixel field of view, with pixel sizes
chosen to be equal to ${\sim}1/8$ of the nominal beam FWHM. We do not place
any CLEAN boxes, and instead we permit the algorithm to operate on the entire
field of view. We use a loop gain of 0.1, and as our stopping criterion we
require that the peak residual be $\sim$5 times the expected image RMS level.

We carry out imaging separately for every channel in each spectral window to
produce image cubes, and we also carry out continuum imaging by combining
multiple channels to form a single image. For sources with detected 183 GHz
spectral line emission (see Section 3), we carry out continuum imaging using
only the continuum spectral window. For sources with no detected spectral
lines, we combine all spectral windows during continuum imaging. The spectra
and continuum images for each of our targets are shown in Appendix B. For the
two sources with other molecular line emission, we instead perform continuum
imaging using only those channels in the continuum spectral window that lack
spectral line emission; the images of these other molecular lines are shown in
Appendix C.

###  2.3 Spectral line identification

We manually inspect spectra extracted from the central region of each of the
image cubes for signs of spectral line emission, both before and after
carrying out boxcar smoothing in frequency by 10 channels (corresponding to a
$\sim$2 km s-1 post-averaging spectral resolution). For some targets, the
spectral line emission is sufficiently strong for the detection to be
unambiguous. Other targets have weaker emission, and so we establish some
detection criteria below.

After initially estimating the spectral sensitivity by computing the RMS
across the full spectral window666We note that for our lowest-redshift targets
(e.g., Circinus, NGC 1386), the steep transmission gradient caused by the
atmospheric 183 GHz line causes the the RMS noise level to change
substantially across the spectral windows. We do not attempt to compensate for
this effect when computing the RMS, so our estimate is an average value across
the entire spectral range., we identify any peaks that reach the
${\gtrsim}5\sigma$ level. If the source shows comparably elevated emission
across multiple channels, then we consider it to be a detection.

##  3 Results and discussion

In total, we detect 183 GHz spectral line emission in 13 out of 20 targeted
sources; LABEL:tab:Detections provides a summary of the targets detected in
183 GHz line emission and in continuum, and at what sensitivities. The
detected 183 GHz spectra are shown in Figures 1, 2, and 3, where they are also
compared against 22 GHz spectra. The full set of 183 GHz spectra and continuum
images for each of our targets is shown in Appendix B; for sources detected in
both spectral line and continuum emission, we typically find that the line and
continuum emission are spatially coincident.

Because maser sources are expected to be heavily anisotropic emitters, the
intrinsic luminosity of a maser is difficult to estimate without knowledge of
the solid angle into which the emission is beamed. Instead, we characterize
the strength of a maser system using the “isotropic luminosity,”
$L_{\text{iso}}$, which corresponds to the luminosity that an isotropic
emitter would need to have to produced the observed flux level. We compute
isotropic luminosities using

$L_{\text{iso}}=\frac{4\pi D^{2}\nu_{0}}{c}\int S_{v}dv,$ (2)

where $D$ is the distance to the system, $\nu_{0}$ is the rest-frame frequency
of the emission line, and $S_{v}$ is the flux density as a function of
velocity $v$. Scaled to convenient units for 183 GHz maser emission, this
equation becomes

$L_{\text{iso}}=\left(1.9\times 10^{-4}\text{
L}_{\odot}\right)\left(\frac{D}{1\text{ Mpc}}\right)^{2}\left(\frac{\int
S_{v}dv}{1\text{ mJy km\,s${}^{-1}$}}\right).$ (3)

The true intrinsic luminosity is smaller than $L_{\text{iso}}$ by a factor
$\Omega/4\pi$, where $\Omega$ is the solid angle of the maser beam.

In this section, we describe the properties of each detection, and we compare
the 183 GHz maser line characteristics to those of the 22 GHz masers in the
same system. We note that because these megamaser sources are known to be
intrinsically variable on timescales of $\sim$weeks at 22 GHz (Lo, 2005), the
typically large temporal separation between the observations of the 183 GHz
masers presented in this paper and their 22 GHz counterparts from the
literature renders a line-by-line comparison of the two spectra potentially
fraught. Instead, we opt here for a somewhat more qualitative comparison,
aiming to determine whether the two sets of spectra exhibit comparable gross
structures.

J0109-0332 | 1.6 | 0.26 | no detection | 4.80 | …
---|---|---|---|---|---
J0126-0417 | 1.9 | 0.28 | no detection | 1.17 | …
NGC 1068 | 7.4 | 1.13 | no detection | 128 | …
NGC 1194 | 2.2 | 0.34 | 35.2† | 11.1 | $64\pm 12$
NGC 1386 | 9.3 | 0.67 | no detection | 22.7 | …
ESO 558-G009 | 2.2 | 0.35 | 11.8 | 1.67 | $402\pm 41$
IC 485 | 1.6 | 0.15 | 8.4 | 1.15 | $527\pm 57$
J0847-0022 | 1.2 | 0.24 | 6.3 | 1.46 | $1180\pm 112$
Mrk 1419 | 2.2 | 0.29 | 23.6 | 1.07 | $542\pm 71$
IC 2560 | 2.8 | 0.25 | 29.6 | 5.42 | $282\pm 59$
NGC 3393 | 1.8 | 0.21 | no detection | 4.23 | …
UGC 6093 | 1.9 | 0.32 | 16.0 | no detection | $978\pm 86$
ESO 269-G012 | 2.0 | 0.09 | 30.8 | no detection | $1337\pm 165$
CGCG 074-064 | 1.9 | 0.08 | 17.9 | 0.28 | $484\pm 88$
NGC 5495 | 2.1 | 0.11 | 9.9 | no detection | $163\pm 27$
Circinus‡ | 21.8 | 0.42 | 4810 | 90.2 | $165\pm 63$
NGC 5643 | 4.9 | 1.08 | 32.9 | 14.2 | $27\pm 3$
NGC 5765b | 1.3 | 0.05 | 6.5 | 0.73 | $379\pm 75$
CGCG 165-035 | 1.9 | 0.33 | no detection | no detection | …
NGC 6264 | 2.0 | 0.23 | no detection | 0.78 | …

Note. — The RMS spectral sensitivity and peak 183 GHz line flux density is
determined after boxcar smoothing by 10 channels to achieve a $\sim$2 km s-1
velocity resolution. Spectral line detections are determined by the criteria
described in Section 2.3; continuum nondetections are only reported if the
emission level at the expected location of the AGN does not rise above the
$3\sigma$ level.

∗When computing the uncertainty in isotropic luminosity, we incorporate an
assumed 300 km s-1 peculiar-velocity-induced uncertainty in the recession
velocity for galaxies whose distances are estimated using the Hubble law.

†The peak flux density of the NGC 1194 spectrum occurs in a single narrow
(FWHM $<$1 km s-1) line that is offset from the systemic velocity by $\sim$500
km s-1 (see Section 3.1), so we report the flux density here prior to
smoothing. The peak flux density after smoothing over 10 channels occurs in
the systemic set of features and reaches a value of $\sim$16 mJy.

‡Because the Circinus spectrum contains many narrow ($\lesssim$1 km s-1)
lines, we report the RMS spectral sensitivity and peak 183 GHz line flux
density at the native channel resolution of $\sim$0.2 km s-1 (i.e., with no
smoothing applied).

Figure 1: Comparison of 183 GHz spectra (in black) with their 22 GHz
counterparts (in red, vertically offset). The 183 GHz spectra have been boxcar
smoothed across 10 channels, corresponding to a $\sim$2 km s-1 post-averaging
spectral resolution. Note that the pair of spectra in each panel were not
observed concurrently; in most cases, the 22 GHz spectrum predates the 183 GHz
spectrum by multiple years.

###  3.1 NGC 1194

The 22 GHz maser system in NGC 1194 was discovered by the Megamaser Cosmology
Project (MCP), and VLBI observations presented in Kuo et al. (2011) confirm
that it is a disk maser system. Relative to other 22 GHz disk maser systems,
NGC 1194 is notable for its large apparent disk size ($\sim$0.54-1.33 pc
across; see Kuo et al. 2011) and for having several strong (up to $\sim$1 Jy)
and narrow (linewidths $<$1 km s-1) individual maser lines exhibiting a large
degree of variability (Pesce et al., 2015). Our 183 GHz spectrum of NGC 1194
is shown in the upper lefthand panel of Figure 1, along with a 22 GHz spectrum
reproduced from Pesce et al. (2015) (see Fig. 1 in that paper). The 22 GHz
spectrum is an average of multiple epochs of spectral monitoring observations
taken with the Green Bank Telescope (GBT).

The most prominent component in the 183 GHz spectrum of NGC 1194 is the
central set of systemic maser features, which span a velocity range of
$\sim$10 km s-1 around a central velocity of ${\sim}4104$ km s-1. The peak
flux density of the systemic features is $\sim$16 mJy, which is a factor of
$\sim$2 fainter than the typical peak flux density of the 22 GHz systemic
masers.

A second component in the 183 GHz spectrum of NGC 1194 is a single narrow
($\text{FWHM}\approx 0.69$ km s-1) line at a velocity of $\sim$4593.5 km s-1,
which peaks at a flux density of $\sim$35.2 mJy. This spectral line coincides
with the redshifted complex of high-velocity maser features seen at 22 GHz.

Adopting a Hubble law distance of 58 Mpc, the isotropic luminosity of the 183
GHz maser emission from NGC 1194 is $L_{\text{iso}}\approx 64$ L⊙. This value
is a factor of $\sim$2 lower than the $\sim$130 L⊙ isotropic luminosity
observed at 22 GHz (Pesce et al., 2015).

###  3.2 ESO 558-G009

The 22 GHz maser system in ESO 558-G009 was discovered by the MCP, and VLBI
observations presented in Gao et al. (2017) confirm that it is a disk maser
system. Our 183 GHz spectrum of ESO 558-G009 is shown in the upper righthand
panel of Figure 1, along with a 22 GHz spectrum reproduced from Pesce et al.
(2015) (see Fig. 1 in that paper). The 22 GHz spectrum is an average of
multiple epochs of spectral monitoring observations taken with the GBT.

The only significant component seen in the 183 GHz spectrum of ESO 558-G009 is
the central set of systemic maser features, which span a velocity range of
$\sim$40 km s-1 around a central velocity of ${\sim}7585$ km s-1. The peak
flux density of the systemic features is $\sim$11.8 mJy, which is a factor of
$\sim$6 fainter than the typical peak flux density of the 22 GHz systemic
masers.

With the overlaid 22 GHz spectrum in Figure 1, apparent features in the 183
GHz spectrum near 7075 km s-1, 7150 km s-1, and 8120 km s-1 are suggestive
because they coincide with high-velocity features at 22 GHz. However, none of
these features rise above the $3\sigma$ significance level, so we cannot
confirm high-velocity masers in the 183 GHz spectrum of ESO 558-G009.

Adopting a Hubble law distance of 110 Mpc, the isotropic luminosity of the 183
GHz maser emission from ESO 558-G009 is $L_{\text{iso}}\approx 402$ L⊙. This
value is a factor of $\sim$2 lower than the $\sim$709 L⊙ isotropic luminosity
observed at 22 GHz (Pesce et al., 2015).

###  3.3 IC 485

The 22 GHz maser system in IC 485 was discovered by the MCP and identified as
a disk maser system because of its spectral structure (Pesce et al., 2015),
though the faintness of the high-velocity features makes confirmation of this
system’s disk nature with VLBI observations difficult (Ladu et al. in prep.).
Our 183 GHz spectrum of IC 485 is shown in the middle lefthand panel of Figure
1, along with a 22 GHz spectrum reproduced from Pesce et al. (2015) (see Fig.
1 in that paper). The 22 GHz spectrum is an average of multiple epochs of
spectral monitoring observations taken with the GBT.

The only component seen in the 183 GHz spectrum of IC 485 is the central set
of systemic maser features, which span a velocity range of $\sim$60 km s-1
around a central velocity of ${\sim}8355$ km s-1. The peak flux density of the
systemic features is $\sim$8.4 mJy, which is a factor of $\sim$8 fainter than
the typical peak flux density of the 22 GHz systemic masers. No high-velocity
maser features are detected at 183 GHz, consistent with their relative
faintness at 22 GHz (where they are more than an order of magnitude fainter
than the systemic features).

Adopting a Hubble law distance of 119 Mpc, the isotropic luminosity of the 183
GHz maser emission from IC 485 is $L_{\text{iso}}\approx 527$ L⊙. This value
is a factor of $\sim$2 lower than the $\sim$1061 L⊙ isotropic luminosity
observed at 22 GHz (Pesce et al., 2015).

###  3.4 J0847-0022

The 22 GHz maser system in J0847-0022 was discovered by the MCP and identified
as a disk maser system because of its spectral structure (Pesce et al., 2015),
though to date there have been no published VLBI observations to confirm the
disk maser nature of this system. Our 183 GHz spectrum of J0847-0022 is shown
in the middle righthand panel of Figure 1, along with a 22 GHz spectrum
reproduced from Pesce et al. (2015) (see Fig. 1 in that paper). The 22 GHz
spectrum is an average of multiple epochs of spectral monitoring observations
taken with the GBT.

The only component seen in the 183 GHz spectrum of J0847-0022 is the central
set of systemic maser features, which span a velocity range of $\sim$70 km s-1
around a central velocity of ${\sim}15295$ km s-1. The peak flux density of
the systemic features is $\sim$6.3 mJy, which is a factor of $\sim$5 fainter
than the typical peak flux density of the 22 GHz systemic masers. No high-
velocity maser features are detected at 183 GHz despite them being more
prominent than the systemic features at 22 GHz; the high-velocity features
must thus be at least an order of magnitude fainter at 183 GHz than at 22 GHz.

Adopting a Hubble law distance of 218 Mpc, the isotropic luminosity of the 183
GHz maser emission from J0847-0022 is $L_{\text{iso}}\approx 1180$ L⊙. This
value is a factor of $\sim$2.5 lower than the $\sim$2945 L⊙ isotropic
luminosity observed at 22 GHz (Pesce et al., 2015).

###  3.5 Mrk 1419

The 22 GHz maser system in Mrk 1419 was discovered by Henkel et al. (2002) and
VLBI observations presented in Kuo et al. (2011) confirm that it is a disk
maser system. Our 183 GHz spectrum of Mrk 1419 is shown in the lower lefthand
panel of Figure 1, along with a 22 GHz spectrum reproduced from Pesce et al.
(2015) (see Fig. 1 in that paper). The 22 GHz spectrum is an average of
multiple epochs of spectral monitoring observations taken with the GBT.

The only component seen in the 183 GHz spectrum of Mrk 1419 is the central set
of systemic maser features, which span a velocity range of $\sim$100 km s-1
around a central velocity of ${\sim}4885$ km s-1. The peak flux density of the
systemic features is $\sim$23.6 mJy, which is comparable to the typical peak
flux density of the 22 GHz systemic masers. However, no high-velocity maser
features are detected at 183 GHz despite them being comparably prominent to
the systemic features at 22 GHz; the high-velocity features must thus be at
least $\sim$4 times fainter at 183 GHz than at 22 GHz.

Adopting a Hubble law distance of 70 Mpc, the isotropic luminosity of the 183
GHz maser emission from Mrk 1419 is $L_{\text{iso}}\approx 542$ L⊙. This value
is comparable to the $\sim$565 L⊙ isotropic luminosity observed at 22 GHz
(Pesce et al., 2015).

###  3.6 IC 2560

The 22 GHz maser system in IC 2560 was discovered by Braatz et al. (1996), and
VLBI observations presented in Nakai et al. (1998) and Ishihara et al. (2001)
confirm that it is a disk maser system. Our 183 GHz spectrum of IC 2560 is
shown in the lower righthand panel of Figure 1, along with a 22 GHz spectrum
reproduced from Pesce et al. (2015) (see Fig. 1 in that paper). The 22 GHz
spectrum is an average of multiple epochs of spectral monitoring observations
taken with the GBT.

The only component seen in the 183 GHz spectrum of IC 2560 is the central set
of systemic maser features, which span a velocity range of $\sim$60 km s-1
around a central velocity of ${\sim}2925$ km s-1. The peak flux density of the
systemic features is $\sim$29.6 mJy, which is a factor of $\sim$5 fainter than
the typical peak flux density of the 22 GHz systemic masers. The heavily
blended structure of the systemic maser complex at 183 GHz is similar to that
at 22 GHz, with no prominent narrow features present.

Adopting a Hubble law distance of 42 Mpc, the isotropic luminosity of the 183
GHz maser emission from IC 2560 is $L_{\text{iso}}\approx 282$ L⊙. This value
is comparable to the $\sim$210 L⊙ isotropic luminosity observed at 22 GHz
(Pesce et al., 2015).

###  3.7 UGC 6093

Figure 2: Continuation of Figure 1.

The 22 GHz maser system in UGC 6093 was discovered by the MCP, and VLBI
observations presented in Zhao et al. (2018) confirm that it is a disk maser
system. Our 183 GHz spectrum of UGC 6093 is shown in the upper lefthand panel
of Figure 2, along with a 22 GHz spectrum reproduced from Pesce et al. (2015)
(see Fig. 1 in that paper). The 22 GHz spectrum is an average of multiple
epochs of spectral monitoring observations taken with the GBT.

The only significant component seen in the 183 GHz spectrum of UGC 6093 is the
central set of systemic maser features, which span a velocity range of
$\sim$30 km s-1 around a central velocity of ${\sim}10848$ km s-1. The peak
flux density of the systemic features is $\sim$16.0 mJy, which is a factor of
$\sim$12 fainter than the typical peak flux density of the 22 GHz systemic
masers.

Similar to ESO 558-G009, the overlaid 22 GHz spectrum in Figure 2 draws the
eye to some features in the 183 GHz spectrum of UGC 6093 near 10100 km s-1 and
11600 km s-1 that are suggestive because they coincide with high-velocity
features at 22 GHz. However, we again find that none of these features rise
above the $3\sigma$ significance level, so we cannot confirm the presence of
high-velocity masers in the 183 GHz spectrum of UGC 6093.

Adopting a Hubble law distance of 154 Mpc, the isotropic luminosity of the 183
GHz maser emission from UGC 6093 is $L_{\text{iso}}\approx 978$ L⊙. This value
is comparable to the $\sim$1048 L⊙ isotropic luminosity observed at 22 GHz
(Pesce et al., 2015).

###  3.8 ESO 269-G012

The 22 GHz maser system in ESO 269-G012 was discovered by Greenhill et al.
(2003a) and identified as a disk maser system because of its spectral
structure, though VLBI observations have not been carried out to confirm the
disk maser nature of this system. Our 183 GHz spectrum of ESO 269-G012 is
shown in the upper righthand panel of Figure 2, along with a 22 GHz spectrum
reproduced from Greenhill et al. (2003a) (see Fig. 1 in that paper). The 22
GHz spectrum was observed with the 70 m Deep Space Network dish at
Tidbinbilla.

The 183 GHz spectrum of ESO 269-G012 shows three prominent sets of maser
features, coinciding in velocity with the three groups of maser features seen
in the 22 GHz spectrum and matching the expected spectral structure for an
edge-on disk system. The line complexes in the 183 GHz spectrum of ESO
269-G012 are relatively rich compared to many of the other spectra in our
sample, with each of the three sets of lines containing multiple identifiable
maser features. The peak flux density of the systemic features in the 183 GHz
spectrum is $\sim$17.2 mJy, which is comparable to the peak flux density of
the 22 GHz systemic masers. The peak flux density of the high-velocity
features is $\sim$30.8 mJy, which is a factor of $\sim$3 fainter than the peak
flux density seen in the 22 GHz high-velocity masers.

Adopting a Hubble law distance of 72 Mpc, the isotropic luminosity of the 183
GHz maser emission from ESO 269-G012 is $L_{\text{iso}}\approx 1337$ L⊙. This
value is roughly 2.5 times larger than the $\sim$496 L⊙ isotropic luminosity
observed at 22 GHz (Greenhill et al., 2003a; Pesce et al., 2015).

###  3.9 CGCG 074-064

The 22 GHz maser system in CGCG 074-064 was discovered by the MCP, and VLBI
observations presented in Pesce et al. (2020a) confirm that it is a disk maser
system. Our 183 GHz spectrum of CGCG 074-064 is shown in the middle lefthand
panel of Figure 2, along with a 22 GHz spectrum reproduced from Pesce et al.
(2015) (see Fig. 1 in that paper). The 22 GHz spectrum is an average of
multiple epochs of spectral monitoring observations taken with the GBT.

The most prominent component in the 183 GHz spectrum of CGCG 074-064 is the
central set of systemic maser features, which span a velocity range of
$\sim$60 km s-1 around a central velocity of ${\sim}6910$ km s-1. The peak
flux density of the systemic features is $\sim$17.9 mJy, which is a factor of
$\sim$11 fainter than the typical peak flux density of the 22 GHz systemic
masers.

A second component in the 183 GHz spectrum of CGCG 074-064 is a weaker (peak
flux density of $\sim$6.7 mJy) feature spanning a velocity range of $\sim$20
km s-1 around a central velocity of ${\sim}6387$ km s-1. This spectral feature
coincides with the most redshifted end of the blueshifted complex of high-
velocity maser features seen at 22 GHz.

Adopting the symmetrized $87.6\pm 7.5$ Mpc distance measurement from Pesce et
al. (2020a), the isotropic luminosity of the 183 GHz maser emission from CGCG
074-064 is $L_{\text{iso}}\approx 484$ L⊙. This value is a factor of $\sim$2
lower than the $\sim$851 L⊙ isotropic luminosity observed at 22 GHz (Pesce et
al., 2015).

###  3.10 NGC 5495

The 22 GHz maser system in NGC 5495 was discovered by Kondratko et al. (2006),
and VLBI observations presented in Gao et al. (2017) confirm that it is a disk
maser system. Our 183 GHz spectrum of NGC 5495 is shown in the middle
righthand panel of Figure 2, along with a 22 GHz spectrum reproduced from
Pesce et al. (2015) (see Fig. 1 in that paper). The 22 GHz spectrum is an
average of multiple epochs of spectral monitoring observations taken with the
GBT.

The only component seen in the 183 GHz spectrum of NGC 5495 is the central set
of systemic maser features, which span a velocity range of $\sim$15 km s-1
around a central velocity of ${\sim}6787$ km s-1. The peak flux density of the
systemic features is $\sim$9.9 mJy, which is a factor of $\sim$12 fainter than
the typical peak flux density of the 22 GHz systemic masers.

Adopting a Hubble law distance of 96 Mpc, the isotropic luminosity of the 183
GHz maser emission from NGC 5495 is $L_{\text{iso}}\approx 163$ L⊙. This value
is a factor of $\sim$4 lower than the $\sim$625 L⊙ isotropic luminosity
observed at 22 GHz (Pesce et al., 2015).

###  3.11 NGC 5643

The 22 GHz maser system in NGC 5643 was discovered by Greenhill et al.
(2003a). The 22 GHz spectrum contains only systemic features, and thus may not
be associated with an edge-on disk; VLBI observations have also not yet been
carried out on this system. Our 183 GHz spectrum of NGC 5643 is shown in the
lower lefthand panel of Figure 2, along with a 22 GHz spectrum reproduced from
Greenhill et al. (2003a) (see Fig. 1 in that paper). The 22 GHz spectrum was
observed with the 70 m Deep Space Network dish at Tidbinbilla.

The only component seen in the 183 GHz spectrum of NGC 5643 is the central set
of systemic maser features, which span a velocity range of $\sim$90 km s-1
around a central velocity of ${\sim}1275$ km s-1. The peak flux density of the
systemic features is $\sim$32.9 mJy, which is a factor of $\sim$6 fainter than
the typical peak flux density of the 22 GHz systemic masers.

Because NGC 5643 is nearby enough that the Hubble law distance is expected to
be very uncertain, we instead adopt a Type Ia Supernova distance of $12.4\pm
0.5$ Mpc to this system from Burns et al. (2020). The isotropic luminosity of
the 183 GHz maser emission from NGC 5643 is $L_{\text{iso}}\approx 27$ L⊙.
This value is a factor of $\sim$2 higher than the $\sim$12.6 L⊙ isotropic
luminosity observed at 22 GHz (Greenhill et al., 2003a).

###  3.12 NGC 5765b

The 22 GHz maser system in NGC 5765b was discovered by the MCP, and VLBI
observations presented in Gao et al. (2016) confirm that it is a disk maser
system. Our 183 GHz spectrum of NGC 5765b is shown in the bottom righthand
panel of Figure 2, along with a 22 GHz spectrum reproduced from Pesce et al.
(2015) (see Fig. 1 in that paper). The 22 GHz spectrum is an average of
multiple epochs of spectral monitoring observations taken with the GBT.

The most prominent component seen in the 183 GHz spectrum of NGC 5765b is the
central set of systemic maser features, which span a velocity range of
$\sim$20 km s-1 around a central velocity of ${\sim}8260$ km s-1. The peak
flux density of the systemic features is $\sim$6.5 mJy, which is a factor of
$\sim$30 fainter than the typical peak flux density of the 22 GHz systemic
masers.

A potential second component near $\sim$8740 km s-1 is most apparent in the
zoomed-in spectrum shown in Figure 4, and is coincident with the brightest
redshifted maser features seen in the 22 GHz spectrum (see Figure 2). However,
this feature does not rise above the $3\sigma$ significance level, so we
cannot confirm the presence of high-velocity masers in the 183 GHz spectrum of
NGC 5765b.

Adopting the $126.3\pm 11.6$ Mpc distance from Gao et al. (2016), the
isotropic luminosity of the 183 GHz maser emission from NGC 5765b is
$L_{\text{iso}}\approx 379$ L⊙. This value is a factor of $\sim$7 lower than
the $\sim$2553 L⊙ isotropic luminosity observed at 22 GHz (Pesce et al.,
2015).

###  3.13 Circinus

Figure 3: Spectra (left) and map (right) for the Circinus galaxy. The left
panel shows a comparison of the 183 GHz spectrum (in black) and the 22 GHz
spectrum (in red), as in Figure 1 and Figure 2; the 22 GHz spectrum was taken
with the Parkes Telescope in 1997 August and is reproduced from Braatz et al.
(2003) (see Fig. 1 in that paper). The systemic velocity of the galaxy is
marked by the blue line at the top. The right panel shows the best-fit
positions of the detected 183 GHz maser spots as cross markers along with
their corresponding $1\sigma$ error ellipses; we show only those masers with
mean position uncertainties less than 20 mas. Each maser spot in the right
panel is colored blue if it has a velocity that is blueshifted with respect to
the recession velocity of the galaxy and red if it has a velocity that is
redshifted with respect to the recession velocity of the galaxy. The opacity
of the error ellipses scales the signal-to-noise ratio of the fit, with more
tightly-constrained positions having a more opaque ellipse. The overplotted
blue, green, and red lines in the right panel mark the warped edge-on disk
from Greenhill et al. (2003b) (see Fig. 3 in that paper).

The 22 GHz maser system in Circinus was only the second ever discovered in an
AGN (Gardner & Whiteoak, 1982), but because of its very southern location
(declination of $-65$ degrees) it was not mapped with VLBI observations until
the work of Greenhill et al. (2003b). The VLBI observations reveal that the
masers in this system originate from both a warped, edge-on accretion disk
(within the innermost $\sim$0.4 pc) as well as a wide-angle outflow extending
out to $\sim$1 pc from the AGN, with the entire complex covering an angular
region ${\sim}90\text{\,mas}\times 50$ mas on the sky. Greenhill et al.
(2003b) showed that the masers associated with the disk and those associated
with the outflow largely span different velocity ranges, with the latter
primarily emitting in the range $\sim$300–570 km s-1 and the former primarily
emitting outside of that range. Our 183 GHz spectrum of Circinus is shown in
the lefthand panel of Figure 3, along with a 22 GHz spectrum reproduced from
Braatz et al. (2003). The 22 GHz spectrum was observed with the 64 m Parkes
Telescope in 1997 August.

The 183 GHz spectrum of Circinus is by far the brightest in our sample; its
peak flux density of nearly 5 Jy is more than two orders of magnitude brighter
than that of the next-brightest source, and it is comparable to the peak flux
density of the 22 GHz spectrum from Braatz et al. (2003). However, we note
that the 22 GHz spectrum is known to be highly time-variable (Greenhill et
al., 1997; McCallum et al., 2007), and it has been observed with peak flux
densities as high as $\sim$40 Jy (e.g., Greenhill et al., 2003b). The 183 GHz
maser complex in Circinus is characterized by a broad pedestal of emission at
a level of $\sim$50–100 mJy and spanning a velocity range from $\sim$250–650
km s-1, interspersed with a number of bright (flux density $\gtrsim$1 Jy),
narrow ($\text{FWHM}\lesssim 1$ km s-1) individual maser lines. The velocity
range covered by the 183 GHz maser lines overlaps primarily with the velocity
range associated with the outflow masers at 22 GHz (Greenhill et al., 2003b).

Because Circinus is nearby enough that the Hubble law distance is very
uncertain, we instead adopt the Tully-Fisher distance measurement of $4.2\pm
0.8$ Mpc to this system from Karachentsev et al. (2013). The isotropic
luminosity of the 183 GHz maser emission from Circinus is
$L_{\text{iso}}\approx 171$ L⊙. This value is a factor of $\sim$2.5 higher
than the $\sim$73 L⊙ isotropic luminosity observed at 22 GHz (e.g., Braatz et
al., 2003).

Although ALMA was in a relatively compact (C-4) configuration for the 2018
December observations of Circinus (with a corresponding resolution of
$\sim$0.6 arcseconds at 183 GHz), the high signal-to-noise ratio provided by
the bright maser lines enables coarse mapping of the 183 GHz system. When
fitting an emission structure observed with some width $\Delta$ and signal-to-
noise ratio $\rho$, the expected uncertainty in the recovered position of the
structure is proportional to $\Delta/\rho$ (see, e.g., Kaper et al., 1966;
Condon, 1997). With a $\sim$0.6 arcsecond beam and signal-to-noise ratios
exceeding 200 for some lines, we expect to achieve relative positional
uncertainties of several mas for a number of the maser features in the 183 GHz
spectrum. To map the 183 GHz maser system in Circinus, we fit a point source
model directly to the visibilities in each spectral channel individually. Our
fitting procedure is detailed in Appendix D, and the resulting map is shown in
the righthand panel of Figure 3. We plot only the maser features that have
geometric mean fitted positional uncertainties less than 20 mas; i.e., we
require that $\sqrt{\sigma_{x}\sigma_{y}}<20$ mas, given an uncertainty in the
$x$-position of $\sigma_{x}$ and an uncertainty in the $y$-position of
$\sigma_{y}$.

Our map shows that the 183 GHz maser system in Circinus is spatially resolved
by these ALMA observations, covering a region that appears to be well-matched
in size and shape to the ${\sim}90\text{\,mas}\times 50$ mas distribution of
22 GHz masers. We have colored each plotted maser spot in Figure 3 by whether
it is redshifted or blueshifted with respect to the 434 km s-1 systemic
velocity of the galaxy. We find that the orientation and sense of rotation for
the 183 GHz system also qualitatively matches that seen at 22 GHz, with
redshifted material located towards the southwest and blueshifted material
located towards the northeast.

To compare more directly with the 22 GHz maser distribution, we overplot in
Figure 3 the midline of the edge-on disk inferred in Greenhill et al. (2003b).
We note that Greenhill et al. (2003b) referenced their VLBI map to a maser
spot with a velocity777Greenhill et al. (2003b) quote the velocity of their
reference spot as being 565.2 km s-1 using the radio convention in the
heliocentric reference frame. We have converted the velocity here to the
optical convention, as is used throughout the rest of this paper. of 566.3 km
s-1, while we have referenced our map to a maser spot with a velocity of 561.6
km s-1; in both cases, the reference maser is the brightest one in the
spectrum. It is plausible that both of these lines originate from the same
parcel of masing gas (with the $\sim$5 km s-1 difference in velocity then
attributable to dynamical evolution during the $\sim$20 years separating the
two observations), but the lack of absolute astrometric referencing between
the two observations means that we cannot be sure; the relative locations of
the 183 GHz maser system and the overplotted 22 GHz disk midline in Figure 3
should thus be understood to be an assumption. Proceeding with this
assumption, we find that the 183 GHz masers live primarily outside of the disk
region, with a $\sim$40 mas gap separating the blueshifted and redshifted
complexes.

The 183 GHz masers are consistent with being located in the disk identified at
22 GHz, but at larger orbital radii than are occupied by the 22 GHz disk
masers. Having 183 GHz masers probe out to larger orbital radii than their 22
GHz counterparts would also be consistent with the results of radiative
transfer modeling, in which the 183 GHz transition is optimally pumped at
lower gas densities and temperatures than the 22 GHz transition (Yates et al.,
1997; Gray et al., 2016). However, the current large uncertainties in our
maser position measurements (relative to the size of the system) ultimately
make it difficult to unambiguously distinguish between a disk and outflow
origin for the 183 GHz masers in Circinus.

###  3.14 Nondetections

The remaining 7 of the 20 systems targeted in our survey are not detected in
183 GHz maser emission. Of these nondetections, 5 of them are disk systems
known to host relatively weak (${\lesssim}$60 mJy peak flux density) 22 GHz
maser systems (Pesce et al., 2015). If the 183 GHz emission in these systems
is an order of magnitude fainter than the corresponding 22 GHz emission – as
is the case for a number of our detections – then our 183 GHz observations
would simply not be sensitive enough to detect it.

The other two nondetections – NGC 1068 and NGC 1386 – host strong ($\sim$1 Jy
peak flux density) 22 GHz maser systems (Claussen et al., 1984; Braatz et al.,
1996), and we would have strongly detected 183 GHz emission from these systems
even if it were an order of magnitude fainter than the corresponding 22 GHz
emission. It is thus perhaps surprising that these systems are not detected by
our 183 GHz observations, because the nondetection implies that any 183 GHz
emission that might be present is substantially more than an order of
magnitude fainter than the corresponding 22 GHz emission. It is also
interesting to note that the 22 GHz maser systems in both NGC 1068 and NGC
1386 exhibit more complicated dynamics and geometries than many of the more
ordered edge-on disk systems in our survey. The 22 GHz maser system in NGC
1068 is distributed in an approximately linear structure on the sky (Greenhill
et al., 1996), but it exhibits sub-Keplerian rotation velocities and may trace
the upper layers of a geometrically thick torus or disk structure whose
surface is irradiated by the jet (e.g., Gallimore et al., 2001); the 22 GHz
maser system in NGC 1386 is not consistent with a rotating structure at all
(Braatz et al., 1997). The lack of observed 183 GHz maser emission from these
systems may thus be associated with their distinct underlying physical
conditions compared with most of the other systems in our survey.

##  4 Summary and conclusions

In this paper we present the results of a survey to identify 183 GHz water
maser emission from AGN targets that are known to host 22 GHz megamaser
emission. We used ALMA to observe 20 targets with a typical spectral
sensitivity of $\sim$2 mJy per $\sim$2 km s-1 of bandwidth, and we detect 183
GHz maser emission in 13 of these targets. In those targets for which both 183
GHz continuum and spectral line emission is detected, we find that the
spectral line emission is spatially coincident with the peak of the continuum
emission.

For the majority (9 out of 13) of our 183 GHz maser detections, the spectrum
exhibits only a single, largely featureless line complex located near the
systemic velocity of the galaxy. Of the remaining 4 sources, two (NGC 1194 and
CGCG 074-064) have a single additional spectral feature detected at velocities
that are offset by several hundred km s-1 from the systemic velocity of the
galaxy, while the other two (Circinus and ESO 269-G012) exhibit considerably
more complex spectral structure. Across all targets, we find that the 183 GHz
maser emission is systematically fainter than the corresponding 22 GHz
emission from the same system (typically by nearly an order of magnitude in
flux density). However, because the emission frequency of the 183 GHz
transition is also approximately an order of magnitude higher than that of the
22 GHz transition, the isotropic luminosities at 183 GHz and 22 GHz are
comparable (typically within a factor of $\sim$2 of one another). I.e.,
assuming similar maser beaming angles, the energy emitted from the 183 GHz
transition is similar to that from the 22 GHz transition.

The 183 GHz spectra of Circinus and ESO 269-G012 contain the richest line
complexes in our sample, with the latter clearly exhibiting the triple-peaked
spectral structure characteristic of edge-on disk systems seen at 22 GHz
(Pesce et al., 2015). The qualitative difference in spectral complexity
between these two sources and the others in our sample likely arises from
their higher signal-to-noise ratio relative to most of the other spectra in
our sample; the situation is reminiscent of the increased spectral richness
observed in 22 GHz maser systems with the GBT compared to prior observations
with less sensitive facilities (Braatz et al., 2003).

Among the galaxies targeted in our sample, Circinus exhibits the strongest 183
GHz emission by more than two orders of magnitude in flux density, with the
brightest line peaking at nearly 5 Jy. Even though the nominal beam size of
these ALMA observations is arcsecond-scale, the high signal-to-noise ratios
achieved by the strong lines in the Circinus spectrum enable us to carry out
coarse mapping of the mas-scale system. We find that the spatial distribution
of the 183 GHz masers is qualitatively similar to that of the 22 GHz masers
(which were mapped with VLBI by Greenhill et al. 2003b), with the blueshifted
masers located to the northeast of the redshifted masers and the entire system
covering ${\sim}90\text{\,mas}\times 50$ mas on the sky. The current
observations are most consistent with the 183 GHz material originating from
the disk identified in the 22 GHz system, but at larger orbital radii than are
occupied by the 22 GHz disk masers.

The survey presented in this paper demonstrates the existence of a population
of 183 GHz AGN megamasers that shows every indication of being associated with
the same edge-on accretion disk material as the known 22 GHz population. We
thus expect that VLBI observations of these systems at 183 GHz – which we note
are not currently possible with any existing facility – would reveal simple
orbital dynamics that present the same utility for measurements of distances
and SMBH masses as the 22 GHz systems currently do, but with roughly an order
of magnitude finer angular resolution. In the meantime, we have demonstrated
that even connected-element observations with ALMA are capable of coarse
mapping of the strongest 183 GHz systems. Extended-configuration ALMA
observations of the Circinus galaxy should permit mapping of the 183 GHz maser
system with an effective resolution that is comparable to that of the best
current 22 GHz maps.

Support for this work was provided by the NSF through grant AST-1935980, and
by the Gordon and Betty Moore Foundation through grant GBMF-5278. This work
has been supported in part by the Black Hole Initiative at Harvard University,
which is funded by grants from the John Templeton Foundation and the Gordon
and Betty Moore Foundation to Harvard University. This paper makes use of the
following ALMA data: ADS/JAO.ALMA#2018.1.00321.S, ADS/JAO.ALMA#2017.1.00909.S.
ALMA is a partnership of ESO (representing its member states), NSF (USA) and
NINS (Japan), together with NRC (Canada), MOST and ASIAA (Taiwan), and KASI
(Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA
Observatory is operated by ESO, AUI/NRAO and NAOJ. The National Radio
Astronomy Observatory is a facility of the National Science Foundation
operated under cooperative agreement by Associated Universities, Inc. 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.

## References

  * Argon et al. (2007) Argon, A. L., Greenhill, L. J., Reid, M. J., Moran, J. M., & Humphreys, E. M. L. 2007, ApJ, 659, 1040, doi: 10.1086/512718
  * Astropy Collaboration et al. (2013) Astropy Collaboration, Robitaille, T. P., Tollerud, E. J., et al. 2013, A&A, 558, A33, doi: 10.1051/0004-6361/201322068
  * Astropy Collaboration et al. (2018) Astropy Collaboration, Price-Whelan, A. M., Sipőcz, B. M., et al. 2018, AJ, 156, 123, doi: 10.3847/1538-3881/aabc4f
  * Astropy Collaboration et al. (2022) Astropy Collaboration, Price-Whelan, A. M., Lim, P. L., et al. 2022, ApJ, 935, 167, doi: 10.3847/1538-4357/ac7c74
  * Braatz et al. (1997) Braatz, J., Greenhill, L., Moran, J., Wilson, A., & Herrnstein, J. 1997, in American Astronomical Society Meeting Abstracts, Vol. 191, American Astronomical Society Meeting Abstracts, 104.02
  * Braatz et al. (2010) Braatz, J. A., Reid, M. J., Humphreys, E. M. L., et al. 2010, ApJ, 718, 657, doi: 10.1088/0004-637X/718/2/657
  * Braatz et al. (1996) Braatz, J. A., Wilson, A. S., & Henkel, C. 1996, ApJS, 106, 51, doi: 10.1086/192328
  * Braatz et al. (2003) Braatz, J. A., Wilson, A. S., Henkel, C., Gough, R., & Sinclair, M. 2003, ApJS, 146, 249, doi: 10.1086/374417
  * Burns et al. (2020) Burns, C. R., Ashall, C., Contreras, C., et al. 2020, ApJ, 895, 118, doi: 10.3847/1538-4357/ab8e3e
  * Cernicharo et al. (2006) Cernicharo, J., Pardo, J. R., & Weiss, A. 2006, ApJ, 646, L49, doi: 10.1086/506473
  * Claussen et al. (1984) Claussen, M. J., Heiligman, G. M., & Lo, K. Y. 1984, Nature, 310, 298, doi: 10.1038/310298a0
  * Condon (1997) Condon, J. J. 1997, PASP, 109, 166, doi: 10.1086/133871
  * Curran et al. (2001) Curran, S. J., Johansson, L. E. B., Bergman, P., Heikkilä, A., & Aalto, S. 2001, A&A, 367, 457, doi: 10.1051/0004-6361:20000462
  * Gallimore et al. (2001) Gallimore, J. F., Henkel, C., Baum, S. A., et al. 2001, ApJ, 556, 694, doi: 10.1086/321616
  * Gao et al. (2016) Gao, F., Braatz, J. A., Reid, M. J., et al. 2016, ApJ, 817, 128, doi: 10.3847/0004-637X/817/2/128
  * Gao et al. (2017) —. 2017, ApJ, 834, 52, doi: 10.3847/1538-4357/834/1/52
  * García-Burillo et al. (2019) García-Burillo, S., Combes, F., Ramos Almeida, C., et al. 2019, A&A, 632, A61, doi: 10.1051/0004-6361/201936606
  * Gardner & Whiteoak (1982) Gardner, F. F., & Whiteoak, J. B. 1982, MNRAS, 201, 13P, doi: 10.1093/mnras/201.1.13P
  * Gray et al. (2016) Gray, M. D., Baudry, A., Richards, A. M. S., et al. 2016, MNRAS, 456, 374, doi: 10.1093/mnras/stv2437
  * Greenhill et al. (1997) Greenhill, L. J., Ellingsen, S. P., Norris, R. P., et al. 1997, ApJ, 474, L103, doi: 10.1086/310434
  * Greenhill et al. (1996) Greenhill, L. J., Gwinn, C. R., Antonucci, R., & Barvainis, R. 1996, ApJ, 472, L21, doi: 10.1086/310346
  * Greenhill et al. (1995) Greenhill, L. J., Jiang, D. R., Moran, J. M., et al. 1995, ApJ, 440, 619, doi: 10.1086/175301
  * Greenhill et al. (2003a) Greenhill, L. J., Kondratko, P. T., Lovell, J. E. J., et al. 2003a, ApJ, 582, L11, doi: 10.1086/367602
  * Greenhill et al. (2003b) Greenhill, L. J., Booth, R. S., Ellingsen, S. P., et al. 2003b, ApJ, 590, 162, doi: 10.1086/374862
  * Hagiwara et al. (2016) Hagiwara, Y., Horiuchi, S., Doi, A., Miyoshi, M., & Edwards, P. G. 2016, ApJ, 827, 69, doi: 10.3847/0004-637X/827/1/69
  * Hagiwara et al. (2021) Hagiwara, Y., Horiuchi, S., Imanishi, M., & Edwards, P. G. 2021, ApJ, 923, 251, doi: 10.3847/1538-4357/ac3089
  * Hagiwara et al. (2013) Hagiwara, Y., Miyoshi, M., Doi, A., & Horiuchi, S. 2013, ApJ, 768, L38, doi: 10.1088/2041-8205/768/2/L38
  * Harris et al. (2020) Harris, C. R., Millman, K. J., van der Walt, S. J., et al. 2020, Nature, 585, 357, doi: 10.1038/s41586-020-2649-2
  * Henkel et al. (2002) Henkel, C., Braatz, J. A., Greenhill, L. J., & Wilson, A. S. 2002, A&A, 394, L23, doi: 10.1051/0004-6361:20021297
  * Herrnstein et al. (1999) Herrnstein, J. R., Moran, J. M., Greenhill, L. J., et al. 1999, Nature, 400, 539, doi: 10.1038/22972
  * Högbom (1974) Högbom, J. A. 1974, A&AS, 15, 417
  * Humphreys et al. (2005) Humphreys, E. M. L., Greenhill, L. J., Reid, M. J., et al. 2005, ApJ, 634, L133, doi: 10.1086/498890
  * Humphreys et al. (2013) Humphreys, E. M. L., Reid, M. J., Moran, J. M., Greenhill, L. J., & Argon, A. L. 2013, ApJ, 775, 13, doi: 10.1088/0004-637X/775/1/13
  * Humphreys et al. (2016) Humphreys, E. M. L., Vlemmings, W. H. T., Impellizzeri, C. M. V., et al. 2016, A&A, 592, L13, doi: 10.1051/0004-6361/201629168
  * Hunter (2007) Hunter, J. D. 2007, Computing in Science & Engineering, 9, 90, doi: 10.1109/MCSE.2007.55
  * Imanishi et al. (2021) Imanishi, M., Hagiwara, Y., Horiuchi, S., Izumi, T., & Nakanishi, K. 2021, MNRAS, 502, L79, doi: 10.1093/mnrasl/slab006
  * Imanishi et al. (2022) Imanishi, M., Nakanishi, K., Izumi, T., & Baba, S. 2022, ApJ, 926, 159, doi: 10.3847/1538-4357/ac3a68
  * Ishihara et al. (2001) Ishihara, Y., Nakai, N., Iyomoto, N., et al. 2001, PASJ, 53, 215, doi: 10.1093/pasj/53.2.215
  * Israel (1992) Israel, F. P. 1992, A&A, 265, 487
  * Jones et al. (1999) Jones, K. L., Koribalski, B. S., Elmouttie, M., & Haynes, R. F. 1999, MNRAS, 302, 649, doi: 10.1046/j.1365-8711.1999.02057.x
  * Kaper et al. (1966) Kaper, H. G., Smits, D. W., Schwarz, U., Takakubo, K., & van Woerden, H. 1966, Bull. Astron. Inst. Netherlands, 18, 465
  * Karachentsev et al. (2013) Karachentsev, I. D., Makarov, D. I., & Kaisina, E. I. 2013, AJ, 145, 101, doi: 10.1088/0004-6256/145/4/101
  * Kondratko et al. (2006) Kondratko, P. T., Greenhill, L. J., Moran, J. M., et al. 2006, ApJ, 638, 100, doi: 10.1086/498641
  * König et al. (2017) König, S., Martín, S., Muller, S., et al. 2017, A&A, 602, A42, doi: 10.1051/0004-6361/201630331
  * Kukolich (1969) Kukolich, S. G. 1969, J. Chem. Phys., 50, 3751, doi: 10.1063/1.1671623
  * Kuo et al. (2013) Kuo, C. Y., Braatz, J. A., Reid, M. J., et al. 2013, ApJ, 767, 155, doi: 10.1088/0004-637X/767/2/155
  * Kuo et al. (2011) Kuo, C. Y., Braatz, J. A., Condon, J. J., et al. 2011, ApJ, 727, 20, doi: 10.1088/0004-637X/727/1/20
  * Kuo et al. (2015) Kuo, C. Y., Braatz, J. A., Lo, K. Y., et al. 2015, ApJ, 800, 26, doi: 10.1088/0004-637X/800/1/26
  * Kuo et al. (2018) Kuo, C. Y., Constantin, A., Braatz, J. A., et al. 2018, ApJ, 860, 169, doi: 10.3847/1538-4357/aac498
  * Kuo et al. (2020) Kuo, C. Y., Hsiang, J. Y., Chung, H. H., et al. 2020, ApJ, 892, 18, doi: 10.3847/1538-4357/ab781d
  * Lo (2005) Lo, K. Y. 2005, ARA&A, 43, 625, doi: 10.1146/annurev.astro.41.011802.094927
  * McCallum et al. (2007) McCallum, J. N., Ellingsen, S. P., & Lovell, J. E. J. 2007, MNRAS, 376, 549, doi: 10.1111/j.1365-2966.2007.11492.x
  * McMullin et al. (2007) McMullin, J. P., Waters, B., Schiebel, D., Young, W., & Golap, K. 2007, in Astronomical Society of the Pacific Conference Series, Vol. 376, Astronomical Data Analysis Software and Systems XVI, ed. R. A. Shaw, F. Hill, & D. J. Bell, 127
  * Müller et al. (2005) Müller, H. S. P., Schlöder, F., Stutzki, J., & Winnewisser, G. 2005, Journal of Molecular Structure, 742, 215, doi: 10.1016/j.molstruc.2005.01.027
  * Nakai et al. (1998) Nakai, N., Inoue, M., Hagiwara, Y., Miyoshi, M., & Diamond, P. J. 1998, in Astronomical Society of the Pacific Conference Series, Vol. 144, IAU Colloq. 164: Radio Emission from Galactic and Extragalactic Compact Sources, ed. J. A. Zensus, G. B. Taylor, & J. M. Wrobel, 237
  * Nakajima et al. (2018) Nakajima, T., Takano, S., Kohno, K., Harada, N., & Herbst, E. 2018, PASJ, 70, 7, doi: 10.1093/pasj/psx153
  * Neufeld & Melnick (1991) Neufeld, D. A., & Melnick, G. J. 1991, ApJ, 368, 215, doi: 10.1086/169685
  * Pesce et al. (2015) Pesce, D. W., Braatz, J. A., Condon, J. J., et al. 2015, ApJ, 810, 65, doi: 10.1088/0004-637X/810/1/65
  * Pesce et al. (2016) Pesce, D. W., Braatz, J. A., & Impellizzeri, C. M. V. 2016, ApJ, 827, 68, doi: 10.3847/0004-637X/827/1/68
  * Pesce et al. (2020a) Pesce, D. W., Braatz, J. A., Reid, M. J., et al. 2020a, ApJ, 890, 118, doi: 10.3847/1538-4357/ab6bcd
  * Pesce et al. (2020b) —. 2020b, ApJ, 891, L1, doi: 10.3847/2041-8213/ab75f0
  * Pickett et al. (1998) Pickett, H. M., Poynter, R. L., Cohen, E. A., et al. 1998, J. Quant. Spec. Radiat. Transf., 60, 883, doi: 10.1016/S0022-4073(98)00091-0
  * Qiu et al. (2020) Qiu, J., Zhang, J., Zhang, Y., Jia, L., & Tang, X. 2020, A&A, 634, A125, doi: 10.1051/0004-6361/201935800
  * Reid et al. (2009) Reid, M. J., Braatz, J. A., Condon, J. J., et al. 2009, ApJ, 695, 287, doi: 10.1088/0004-637X/695/1/287
  * Reid et al. (2013) —. 2013, ApJ, 767, 154, doi: 10.1088/0004-637X/767/2/154
  * Reid et al. (2019) Reid, M. J., Pesce, D. W., & Riess, A. G. 2019, ApJ, 886, L27, doi: 10.3847/2041-8213/ab552d
  * Schinnerer et al. (2000) Schinnerer, E., Eckart, A., Tacconi, L. J., Genzel, R., & Downes, D. 2000, ApJ, 533, 850, doi: 10.1086/308702
  * Speagle (2020) Speagle, J. S. 2020, MNRAS, 493, 3132, doi: 10.1093/mnras/staa278
  * Thompson et al. (2017) Thompson, A. R., Moran, J. M., & Swenson, George W., J. 2017, Interferometry and Synthesis in Radio Astronomy, 3rd Edition (Springer Nature), doi: 10.1007/978-3-319-44431-4
  * Virtanen et al. (2020) Virtanen, P., Gommers, R., Oliphant, T. E., et al. 2020, Nature Methods, 17, 261, doi: 10.1038/s41592-019-0686-2
  * Viti et al. (2014) Viti, S., García-Burillo, S., Fuente, A., et al. 2014, A&A, 570, A28, doi: 10.1051/0004-6361/201424116
  * Yates et al. (1997) Yates, J. A., Field, D., & Gray, M. D. 1997, MNRAS, 285, 303, doi: 10.1093/mnras/285.2.303
  * Zhao et al. (2018) Zhao, W., Braatz, J. A., Condon, J. J., et al. 2018, ApJ, 854, 124, doi: 10.3847/1538-4357/aaa95c

##  Appendix A Additional calibration details

In this section we provide specific calibration details relevant for
individual science targets. We focus primarily on the flux density calibration
fidelity, which is unique to each target, though we also highlight any
elements of the calibration that deviate from the standard procedure described
in Section 2.1. The calibrator details for each science target are listed in
LABEL:tab:Calibrators.

The flux density calibrators contained in the ALMA catalog typically exhibit
substantial time variability (i.e., varying by factors of several in total
flux density over the multi-year monitoring period), so when applying the flux
density calibration we pick the dates that fall closest in time to the
observations for each source. The ALMA monitoring observations typically do
not have observations at 183 GHz directly, so we instead interpolate the
observed flux densities – one higher frequency $\nu_{\text{high}}$ and one
lower frequency $\nu_{\text{low}}$ – to determine the appropriate 183 GHz
calibration.

We also use the historical ALMA monitoring observations to empirically
estimate the reliability of the flux density calibration for each target. For
all monitoring epochs in which a third frequency $\nu_{\text{mid}}$ that falls
between $\nu_{\text{low}}$ and $\nu_{\text{high}}$ was also observed, we take
a difference between the observed flux density at $\nu_{\text{mid}}$ and that
predicted by the spectral index derived from the observations at
$\nu_{\text{low}}$ and $\nu_{\text{high}}$. We take the standard deviation of
these differences over all historical ALMA monitoring observations as our
measure of flux density calibration fidelity, which we add in quadrature to
the statistical uncertainty and then cast as a fraction $f_{\text{cal}}$ of
the flux density itself; this value is reported in LABEL:tab:Calibrators for
each target. We find a typical flux density precision between 5% and 10%.

J0109-0332 | J0006-0623 | J0108+0135 | J0006-0623 | 1.98 | 103.5 | $-0.687$ | 8.0
---|---|---|---|---|---|---|---
J0126-0417 | J0006-0623 | J0141-0928 | J0006-0623 | 1.98 | 103.5 | $-0.687$ | 8.0
NGC 1068 (2018 Nov 09) | J0423-0120 | J0217+0144 | J0423-0120 | 3.59 | 103.5 | $-0.496$ | 9.0
NGC 1068 (2018 Nov 16) | J0423-0120 | J0217+0144 | J0423-0120 | 3.59 | 103.5 | $-0.496$ | 9.0
NGC 1194 | J0238+1636 | J0241-0815 | J0238+1636 | 0.59 | 103.5 | $-0.653$ | 5.7
NGC 1386 | J0522-3627 | J0403-3605 | J0522-3627 | 5.46 | 103.5 | $-0.132$ | 4.1
ESO 558-G009 | J0522-3627 | J0702-1951 | J0522-3627 | 5.35 | 103.5 | $-0.230$ | 4.1
IC 485 (2018 Oct 31) | J0854+2006 | J0748+2400 | J0854+2006 | 4.51 | 103.5 | $-0.337$ | 4.2
IC 485 (2018 Dec 06) | J0854+2006 | J0748+2400 | J0854+2006 | 4.52 | 103.5 | $-0.487$ | 4.2
J0847-0022 | J0750+1231 | J0839+0104 | J0750+1231 | 0.97 | 103.5 | $-0.643$ | 8.7
Mrk 1419 | J0854+2006 | J0930+0034 | J0854+2006 | 4.08 | 103.5 | $-0.496$ | 4.3
IC 2560 (2018 Oct 02) | J1058+0133 | J1037-2934 | J1058+0133 | 4.10 | 91.5 | $-0.488$ | 5.2
IC 2560 (2018 Oct 30) | J1058+0133 | J1037-2934 | J1058+0133 | 4.06 | 103.5 | $-0.542$ | 4.0
IC 2560 (2018 Nov 08) | J1107-4449 | J1037-2934 | J1107-4449 | 1.37 | 103.5 | $-0.677$ | 6.6
NGC 3393 (2018 Oct 30) | J1107-4449 | J1037-2934 | J1107-4449 | 1.36 | 103.5 | $-0.712$ | 6.1
NGC 3393 (2018 Nov 14) | J1107-4449 | J1037-2934 | J1107-4449 | 1.41 | 103.5 | $-0.701$ | 6.6
UGC 6093 | J1058+0133 | J1041+0610 | J1058+0133 | 4.06 | 103.5 | $-0.542$ | 4.0
ESO 269-G012 | J1427-4206 | J1326-5256 | J1427-4206 | 3.77 | 91.5 | $-0.585$ | 6.2
CGCG 074-064 | J1256-0547 | J1415+1320 | J1256-0547 | 14.93 | 91.5 | $-0.537$ | 6.5
NGC 5495 | J1517-2422 | J1408-2900 | J1517-2422 | 3.15 | 91.5 | $-0.321$ | 5.7
Circinus | J1427-4206 | J1326-5256 | J1427-4206 | 3.77 | 91.5 | $-0.585$ | 6.2
NGC 5643 (2018 Aug 27) | J1256-0547 | J1427-4206 | J1256-0547 | 11.42 | 103.5 | $-0.548$ | 6.4
NGC 5643 (2018 Sep 20) | J1517-2422 | J1427-4206 | J1517-2422 | 2.84 | 103.5 | $-0.256$ | 5.4
NGC 5765b | J1504+1029 | J1504+1029 | J1504+1029 | 1.96 | 103.5 | $-0.626$ | 4∗
CGCG 165-035 | J1550+0527 | J1516+1932 | J1550+0527 | 0.90 | 91.5 | $-0.809$ | 5.7
NGC 6264 | J1751+0939 | J1653+3107 | J1751+0939 | 3.57 | 103.5 | $-0.451$ | 4.4

Note. — Information about the calibrator targets used for each of the
observations reported in this paper.

∗Because we use the targeted gain calibrator rather than the targeted flux
calibrator for NGC 5765b flux density calibration, there are insufficient
historical ALMA observations to estimate the flux density calibration
precision. The $f_{\text{cal}}$ value reported here thus reflects just the
statistical uncertainty in the observations at $\nu_{\text{low}}$ and
$\nu_{\text{high}}$.

###  A.1 J0109-0332

The flux density calibrator for J0109-0332 is J0006-0623. Its flux density as
reported by ALMA on 2018 Nov 08 was $0.88\pm 0.04$ Jy at 337.5 GHz, and on
2018 Nov 09 it was $1.98\pm 0.04$ Jy at 103.5 GHz. We use these values to
derive a spectral index of $\alpha=-0.687\pm 0.042$.

###  A.2 J0126-0417

The flux density calibrator for J0126-0417 is J0006-0623, which is the same as
for J0109-0332.

###  A.3 NGC 1068

NGC 1068 was observed twice, once on 2018 Nov 09 and again on 2018 Nov 16. The
flux density calibrator for NGC 1068 was J0423-0120 during both observations.
Its flux density as reported by ALMA on 2018 Nov 10 was $1.98\pm 0.05$ Jy at
343.5 GHz, and on 2018 Nov 12 it was $3.59\pm 0.07$ Jy at 103.5 GHz. We use
these values to derive a spectral index of $\alpha=-0.496\pm 0.027$.

Because NGC 1068 was observed on two separate days, we combine the MSs using
CASA’s concat task after calibration and prior to imaging.

###  A.4 NGC 1194

The flux density calibrator for NGC 1194 is J0238+1636. Its flux density as
reported by ALMA on 2018 Nov 19 was $0.27\pm 0.02$ Jy at 343.5 GHz and
$0.59\pm 0.03$ Jy at 103.5 GHz, which we use to derive a spectral index of
$\alpha=-0.653\pm 0.076$.

###  A.5 NGC 1386

The flux density calibrator for NGC 1386 is J0522-3627. Its flux density as
reported by ALMA on 2018 Nov 19 was $4.65\pm 0.10$ Jy at 349.5 GHz and
$5.46\pm 0.14$ Jy at 103.5 GHz, which we use to derive a spectral index of
$\alpha=-0.132\pm 0.028$.

###  A.6 ESO 558-G009

The flux density calibrator for ESO 558-G009 is J0522-3627. Its flux density
as reported by ALMA on 2018 Oct 30 was $4.06\pm 0.09$ Jy at 343.5 GHz and
$5.35\pm 0.14$ Jy at 103.5 GHz, which we use to derive a spectral index of
$\alpha=-0.230\pm 0.029$.

###  A.7 IC 485

IC 485 was observed twice, once on 2018 Oct 31 and again on 2018 Dec 06. The
flux density calibrator for IC 485 was J0854+2006 during both observations.
Its flux density as reported by ALMA on 2018 Oct 30 was $3.01\pm 0.06$ Jy at
343.5 GHz and $4.51\pm 0.07$ Jy at 103.5 GHz. We use these values to derive a
spectral index of $\alpha=-0.337\pm 0.021$.

The flux density of J0854+2006 as reported by ALMA on 2018 Dec 01 was $2.52\pm
0.05$ Jy at 343.5 GHz, and on 2018 Dec 02 it was $4.52\pm 0.07$ Jy at 103.5
GHz. We use these values to derive a spectral index of $\alpha=-0.487\pm
0.021$.

Because IC 485 was observed on two separate days, we combine the MSs using
CASA’s concat task after calibration and prior to imaging.

###  A.8 J0847-0022

The flux density calibrator for J0847-0022 is J0750+1231. Its flux density as
reported by ALMA on 2018 Nov 10 was $0.97\pm 0.03$ Jy at 103.5 GHz and
$0.45\pm 0.04$ Jy at 343.5 GHz, which we use to derive a spectral index of
$\alpha=-0.643\pm 0.079$.

###  A.9 Mrk 1419

The flux density calibrator for Mrk 1419 is J0854+2006. Its flux density as
reported by ALMA on 2018 Nov 9 was $4.08\pm 0.09$ Jy at 103.5 GHz and $2.25\pm
0.06$ Jy at 343.5 GHz, which we use to derive a spectral index of
$\alpha=-0.496\pm 0.029$.

###  A.10 IC 2560

IC 2560 was observed three times, on 2018 Oct 02, 2018 Oct 30, and 2018 Nov
08. The flux density calibrator for IC 2560 was J1058+0133 during the first
two observations and J1107-4449 during the third observation.

The flux density of J1058+0133 as reported by ALMA on 2018 Sep 27 was $4.10\pm
0.10$ Jy at 91.5 GHz, and on 2018 Oct 02 it was $2.15\pm 0.05$ Jy at 343.5
GHz. We use these values to derive a spectral index of $\alpha=-0.488\pm
0.025$ for the 2018 Oct 02 observation.

The flux density of J1058+0133 as reported by ALMA on 2018 Oct 30 was $4.06\pm
0.08$ Jy at 103.5 GHz and $2.12\pm 0.06$ Jy at 343.5 GHz. We use these values
to derive a spectral index of $\alpha=-0.542\pm 0.029$ for the 2018 Oct 30
observation.

The flux density of J1107-4449 as reported by ALMA on 2018 Nov 10 was $1.37\pm
0.04$ Jy at 103.5 GHz and $0.61\pm 0.05$ Jy at 343.5 GHz. We use these values
to derive a spectral index of $\alpha=-0.677\pm 0.073$ for the 2018 Nov 08
observation.

Because IC 2560 was observed on three separate days, we combine the MSs using
CASA’s concat task after calibration and prior to imaging.

###  A.11 NGC 3393

NGC 3393 was observed twice, once on 2018 Oct 30 and again on 2018 Nov 14. The
flux density calibrator for NGC 3393 was J1107-4449 during both observations.
Its flux density as reported by ALMA on 2018 Oct 30 was $1.36\pm 0.03$ Jy at
103.5 GHz and $0.58\pm 0.04$ Jy at 343.5 GHz, which we use to derive a
spectral index of $\alpha=-0.712\pm 0.061$ for the 2018 Oct 30 observation.
Its flux density as reported by ALMA on 2018 Nov 10 was $1.41\pm 0.04$ Jy at
103.5 GHz and $0.61\pm 0.05$ Jy at 343.5 GHz, which we use to derive a
spectral index of $\alpha=-0.701\pm 0.073$ for the 2018 Nov 14 observation.

Because NGC 3393 was observed on two separate days, we combine the MSs using
CASA’s concat task after calibration and prior to imaging.

###  A.12 UGC 6093

The flux density calibrator for UGC 6093 is J1058+0133. Its flux density as
reported by ALMA on 2018 Oct 30 was $4.06\pm 0.08$ Jy at 103.5 GHz and
$2.12\pm 0.06$ Jy at 343.5 GHz, which we use to derive a spectral index of
$\alpha=-0.542\pm 0.029$.

###  A.13 ESO 269-G012

The flux density calibrator for ESO 269-G012 is J1427-4206. Its flux density
as reported by ALMA on 2018 Dec 01 was $1.74\pm 0.05$ Jy at 343.5 GHz, and on
2018 Dec 02 it was $3.77\pm 0.08$ Jy at 91.5 GHz. We use these values to
derive a spectral index of $\alpha=-0.585\pm 0.027$.

###  A.14 CGCG 074-064

The flux density calibrator for CGCG 074-064 is J1256-0547. Its flux density
as reported by ALMA on 2018 Dec 01 was $7.34\pm 0.16$ Jy at 343.5 GHz, and on
2018 Dec 02 it was $14.93\pm 0.29$ Jy at 91.5 GHz. We use these values to
derive a spectral index of $\alpha=-0.537\pm 0.022$.

###  A.15 NGC 5495

The flux density calibrator for NGC 5495 is J1517-2422. Its flux density as
reported by ALMA on 2018 Dec 06 was $2.06\pm 0.07$ Jy at 343.5 GHz, and on
2018 Dec 08 it was $3.15\pm 0.07$ Jy at 91.5 GHz. We use these values to
derive a spectral index of $\alpha=-0.321\pm 0.031$.

###  A.16 Circinus

The flux density calibrator for Circinus is J1427-4206. Its flux density as
reported by ALMA on 2018 Dec 01 was $1.74\pm 0.05$ Jy at 343.5 GHz, and on
2018 Dec 02 it was $3.77\pm 0.08$ Jy at 91.5 GHz. We use these values to
derive a spectral index of $\alpha=-0.585\pm 0.027$.

After initial gain calibration, we further self-calibrate the phases to the
strongest spectral line, averaging over the three channels ($\sim$0.6 km s-1)
centered at 561.6 km s-1 and using a solution interval of 60 seconds. These
phase solutions are applied to all four spectral windows prior to imaging.

###  A.17 NGC 5643

Note that the observing setup for NGC 5643 is different than for the other
targets. It was observed over two days – 2018 Aug 27 and 2018 Sep 20 – and its
spectral window configuration did not contain a dedicated continuum spectral
window. The flux density calibrator for NGC 5643 was J1256-0547 during the
first observation and J1517-2422 during the second observation.

The flux density of J1256-0547 as reported by ALMA on 2018 Aug 28 was
$11.42\pm 0.19$ Jy at 103.5 GHz and $5.92\pm 0.14$ Jy at 343.5 GHz, which we
use to derive a spectral index of $\alpha=-0.548\pm 0.024$ for the 2018 Aug 27
observation.

The flux density of J1517-2422 as reported by ALMA was $2.09\pm 0.06$ Jy at
343.5 GHz on 2018 Sep 18, and $2.84\pm 0.06$ Jy at 103.5 GHz on 2018 Sep 20,
which we use to derive a spectral index of $\alpha=-0.256\pm 0.030$ for the
2018 Sep 20 observation.

Because NGC 5643 was observed on two separate days, we combine the MSs using
CASA’s concat task after calibration and prior to imaging.

###  A.18 NGC 5765b

The targeted flux density calibrator for NGC 5765b, J1550+0527, was quite weak
(flux density less than $\sim$0.5 Jy) at the time of the observation. We thus
use the gain calibrator, J1504+1029, to also calibrate bandpass and flux
density for NGC 5765b. The flux density of J1504+1029 as reported by ALMA on
2018 Nov 11 was $1.96\pm 0.1$ Jy at 103.5 GHz and $0.92\pm 0.05$ Jy at 343.5
GHz, which we use to derive a spectral index of $\alpha=-0.626\pm 0.060$.

###  A.19 CGCG 165-035

The flux density calibrator for CGCG 165-035 is J1550+0527. Its flux density
as reported by ALMA on 2018 Dec 06 was $0.31\pm 0.03$ Jy at 343.5 GHz, and on
2018 Dec 08 it was $0.90\pm 0.03$ Jy at 91.5 GHz. We use these values to
derive a spectral index of $\alpha=-0.809\pm 0.078$.

###  A.20 NGC 6264

The flux density calibrator for NGC 6264 is J1751+0939. Its flux density as
reported by ALMA on 2018 Nov 18 was $3.57\pm 0.06$ Jy at 103.5 GHz and
$2.08\pm 0.08$ Jy at 343.5 GHz, which we use to derive a spectral index of
$\alpha=-0.451\pm 0.035$.

##  Appendix B Spectral line and continuum images at 183 GHz

The 183 GHz spectra for each of our targets are shown in Figure 4, including
those for which no 183 GHz spectral line emission is detected. Figure 5 shows
the corresponding continuum images, again with nondetections included.

Figure 4: 183 GHz spectra for each of the survey targets, after carrying out
boxcar smoothing in frequency by 10 channels (corresponding to a $\sim$2 km
s-1 post-averaging spectral resolution). The velocity axis for spectrum is
referenced to the recession velocities listed in LABEL:tab:Observations. All
panels share the same velocity range, which is explicitly labeled in the lower
left panel. We note that the vertical axis range for the Circinus spectrum has
been restricted (the strongest emission lines are nearly 5 Jy) to more clearly
show the broad pedestal of emission around the recession velocity. Figure 5:
183 GHz continuum images for each of the survey targets. The outermost
contours mark a brightness level of three times the image RMS (see
LABEL:tab:Detections), and each subsequent contour marks a brightness level
that is larger by a factor of two than that of the previous contour. In each
panel, the image is centered on the location of the maser emission (if it is
present) or on the peak of the continuum emission (if maser emission is not
present); for the one galaxy without either continuum or maser emission
detections (CGCG 165-035), we center the image on the expected location of the
galaxy (see LABEL:tab:Observations). The FWHM restoring beam size is shown in
the lower lefthand corner of each panel.

##  Appendix C Other molecular species in Circinus and NGC 1068

In the continuum spectral windows of both Circinus and NGC 1068, we detect
spectral lines other than the 183 GHz H2O maser line. We use CASA’s tclean
task to make image cubes from the continuum spectral windows for both
galaxies, and we use the curve_fit function from the scipy Python library to
to fit a Gaussian function to each pixel in the image cube. We fit for four
free parameters at every pixel location, with the parameterization given by

$S(v)=S_{0}+A\exp\left(-\frac{\left(v-v_{0}\right)^{2}}{2\sigma^{2}}\right).$
(C.1)

Here, $S_{0}$ is a constant-valued background continuum brightness level, $A$
is the peak brightness level of the Gaussian, $v_{0}$ is the center velocity
of the Gaussian, and $\sigma$ is its standard deviation. We use a least
squares optimization scheme to carry out the fits, taking the uncertainty in
each pixel’s value to be given by the continuum RMS level (see
LABEL:tab:Detections). The corresponding uncertainty estimates for each
modeled parameter are then given by the diagonal elements of the inverse
Hessian matrix.

###  C.1 Circinus

There is one additional spectral line detected towards Circinus, which we
associate with the HNC ($\text{v}=0$, $\text{J}=2\text{ - }1$) transition with
a rest-frame frequency of 181.324758 GHz (Müller et al., 2005). Figure 6 shows
the spatial distributions of the $A$, $v_{0}$, and $\sigma$ parameters
associated with the Gaussian fits to the image cube, which effectively provide
an image of the line brightness and its velocity structure. The emission shows
clear and spatially-resolved rotation signatures, with blueshifted material
residing in the northeastern region of the galaxy and the redshifted material
residing in the southwestern region. This sense and orientation of rotation
matches that seen in large-scale neutral hydrogen maps of Circinus (Jones et
al., 1999), and it also matches that seen in the 183 GHz maser system
presented in this paper (see Section 3.13). Other transitions of HNC have
previously been observed towards Circinus by Israel (1992) and Curran et al.
(2001).

Figure 6: HNC emission in Circinus. The left panel shows the spatial
distribution of the fitted line brightness, the middle panel shows the spatial
distribution of the fitted line velocity, and the right panel shows the
spatial distribution of the fitted line dispersion; these quantities
correspond to the $A$, $v_{0}$, and $\sigma$ parameters from Equation C.1,
respectively. Each panel’s image is centered as in Figure 5. We mask out all
regions of the image that have a fractional uncertainty in $A$ larger than 20%
($5\sigma$).

###  C.2 NGC 1068

We detect two different lines in the continuum window of the NGC 1068
spectrum, both of which appear to originate in the circumnuclear disk
(Schinnerer et al., 2000). We associated one of the lines with the same HNC
($\text{v}=0$, $\text{J}=2\text{ - }1$) transition seen towards Circinus; the
distribution of the HNC emission is shown in Figure 7. The other line is
weaker and has a more ambiguous identification, but we associate it with the
HC3N ($\text{v}=0$, $\text{J}=20\text{ - }19$) transition with a rest-frame
frequency of 181.944923 GHz (Müller et al., 2005); the distribution of the
HC3N emission is shown in Figure 8. Both HNC and HC3N have been previously
detected towards NGC 1068 (e.g., Viti et al., 2014; Nakajima et al., 2018; Qiu
et al., 2020).

The emission from both spectral lines shows spatially-resolved rotation
signatures, with blueshifted material residing in the southeastern region of
the circumnuclear disk and the redshifted material residing in the
northwestern region. The emission is strongest in an eastern “knot” within the
circimnuclear disk, as seen in previous observations (e.g. Viti et al., 2014;
García-Burillo et al., 2019; Qiu et al., 2020).

Figure 7: Same as Figure 6, but showing HNC emission in NGC 1068. Figure 8:
Same as Figure 6 and Figure 7, but showing HC3N emission in NGC 1068.

##  Appendix D Maser position modeling

We assume that each maser spot is a point source, for which we fit the
position $(x,y)$ and flux density $S$ in every spectral channel. Our model for
the point source visibility $V(u,v)$ is given by

$V(u,v)=Se^{2\pi i(ux+vy)}.$ (D.1)

We fit directly to the Stokes I complex visibilities in each spectral channel,
after self-calibration and averaging in time across the entire observing
track. Because we have self-calibrated using the brightest spectral line at
561.6 km s-1 (see Section A.16), all positions are referenced to the location
of the associated maser spot.

Within a single spectral channel, our likelihood function for each data point
$i$ is a complex Gaussian (Thompson et al., 2017),

$\ell_{i}=\frac{1}{2\pi\sigma^{2}}\exp\left(-\frac{1}{2\sigma^{2}}\left|V_{i}-\hat{V}_{i}\right|^{2}\right),$
(D.2)

where $\hat{V}_{i}$ represents an observed complex visibility and $\sigma$ is
the uncertainty (which we assume to be the same for all data points within a
single channel). Because we do not have a good measure of $\sigma$ from CASA,
we treat it as a model parameter and fit for it alongside the position and
flux density in each spectral channel. The total likelihood for a single
channel is then the product over all individual data point likelihoods,

$\mathcal{L}=\prod_{i}\ell_{i}.$ (D.3)

We specify uniform priors of $p_{S}\sim\mathcal{U}(0,10)$ (in units of Jy) for
$S$, $p_{x}\sim\mathcal{U}(-0.5,0.5)$ (in units of arcseconds) for $x$ and
$y$, and $p_{\sigma}\sim\mathcal{U}(0,5)$ (in units of Jy) for $\sigma$. We
use the nested sampling package dynesty (Speagle, 2020) to carry out the
fitting separately for each spectral channel.