# The H i in Ring Galaxies Survey (H i-RINGS) - Effects of the bar on the H i
gas in ring galaxies

C. Murugeshan ATNF, CSIRO, Space and Astronomy, PO Box 1130, Bentley, WA 6102,
Australia [ R. Džudžar Centre for Astrophysics and Supercomputing, Swinburne
University of Technology, Hawthorn, Victoria 3122, Australia R. Bagge School
of Physics, The University of New South Wales, Kensington Campus, Old Main
Building, Kensington, Sydney NSW Australia T. O’Beirne International Centre
for Radio Astronomy Research, The University of Western Australia, 35 Stirling
Highway, Crawley, WA 6009, Australia O. I. Wong ATNF, CSIRO, Space and
Astronomy, PO Box 1130, Bentley, WA 6102, Australia V. A. Kilborn Centre for
Astrophysics and Supercomputing, Swinburne University of Technology, Hawthorn,
Victoria 3122, Australia M. E. Cluver Centre for Astrophysics and
Supercomputing, Swinburne University of Technology, Hawthorn, Victoria 3122,
Australia K. A. Lutz Observatoire Astronomique de Strasbourg, Université de
Strasbourg, CNRS, UMR 7550, 67000 Strasbourg, France A. Elagali Minderoo
Foundation, 171 - 173 Mounts Bay Road, Perth, WA 6000, Australia

(dd Mmm YYYY; dd Mmm YYYY; dd Mmm YYYY)

###### Abstract

We present a new high-resolution neutral atomic hydrogen (H i) survey of ring
galaxies using the Australia Telescope Compact Array (ATCA). We target a
sample of 24 ring galaxies from the Buta (1995) Southern Ring Galaxy Survey
Catalogue in order to study the origin of resonance-, collisional- and
interaction-driven ring galaxies. In this work, we present an overview of the
sample and study their global and resolved H i properties. In addition, we
also probe their star formation properties by measuring their star formation
rates (SFR) and their resolved SFR surface density profiles. We find that a
majority of the barred galaxies in our sample are H i deficient, alluding to
the effects of the bar in driving their H i deficiency. Furthermore, for the
secularly evolving barred ring galaxies in our sample, we apply Lindblad’s
resonance theory to predict the location of the resonance rings and find very
good agreement between predictions and observations. We identify rings of H i
gas and/or star formation co-located at one or the other major resonances.
Lastly, we measure the bar pattern speed ($\Omega_{\textrm{bar}}$) for a sub-
sample of our galaxies and find that the values range from 10 – 90 km
s-1kpc-1, in good agreement with previous studies.

###### doi:

10.1017/pasa.2020.32

###### keywords:

galaxies: evolution– galaxies: fundamental parameters– galaxies: ISM–
galaxies: kinematics and dynamics

C<EMAIL_ADDRESS>Tribalism PTY LTD, Lvl 1, 25
King Street Melbourne, 3000 Victoria, Australia International Centre for Radio
Astronomy Research, The University of Western Australia, 35 Stirling Highway,
Crawley, WA 6009, Australia ATNF, CSIRO, Space and Astronomy, PO Box 1130,
Bentley, WA 6102, Australia International Centre for Radio Astronomy Research,
The University of Western Australia, 35 Stirling Highway, Crawley, WA 6009,
Australia ARC Centre of Excellence for All Sky Astrophysics in 3 Dimensions
(ASTRO 3D), Australia Department of Physics and Astronomy, University of the
Western Cape, Robert Sobukwe Road, Bellville, 7535, South Africa School of
Biological Sciences, The University of Western Australia, 35 Stirling Highway,
Crawley, WA 6009, Australia 22 September 2020

##  1 Introduction

Galaxies come in a variety of shapes and sizes and have been traditionally
classified on the basis of their optical morphology (e.g., Hubble 1926; de
Vaucouleurs 1959). Hubble developed the famous “tunning-fork” diagram which
sorted galaxies broadly into spheroidal/elliptical early-type galaxies towards
one end of the fork, gradually transitioning into more and more late-type
disky galaxies, which are observed to possess features such as a nucleus, bar
and spiral arms to name a few. Galaxies were further classified on the basis
of how tightly wound their spiral arms were and on the presence or absence of
a bar, forming the two prongs in the tuning fork diagram. Apart from the
early- and late-type galaxies, Hubble also identified galaxies that did not
specifically fit into one or the other group and classified them as peculiar
galaxies. In the following decades a number of atlases of galaxies were
curated in an attempt to uniformly classify galaxies, one of the most
important ones being the atlas compiled by Sandage (1961), who expanded on the
existing Hubble classification scheme and introduced a special class of
galaxies called ring galaxies. Such galaxies were observed to posses ring-like
structures based on their optical images. This classification scheme was
further revised by de Vaucouleurs et al. (1991), who introduced a number of
sub-classes of ring galaxies, including those galaxies having nuclear, inner
and outer (pseudo-) rings and lenses.

Depending on the morphology of the ring, its position, alignment and size,
ring galaxies are typically classified into the following major types:
Resonance rings, collisional, accretion and polar rings (for a full review see
Buta & Combes 1996). While the exact mechanisms driving the formation of rings
in galaxies are still debated, the following are believed to be some of the
major scenarios:

1\. Resonance rings driven by the bar: Stars and gas orbiting the inner parts
of a galaxy possessing a bar respond to the torques exerted by the non-
axisymmetric potential of the bar and settle into orbits of resonance
(Lindblad 1964; Schwarz 1981). In addition to the accumulation of stars and
gas from the disk, due to the increased density and pressure at these
resonance locations, the gas collapses to form new stars, resulting in the
formation of stellar ring-like structures that is visible in the UV, optical
and infrared (IR) bands (Buta & Combes 1996; Buta et al. 2004; Grouchy et al.
2010 and references therein). Numerous sub-classes exist within resonance
rings, such as nuclear rings, inner (pseudo) rings and outer (pseudo) rings
(e.g. Comerón et al. 2014). Outer pseudo rings are observed to be formed by
two tightly wound spiral arms in late-type galaxies that almost form a
complete ring in the outskirts of galaxies (Buta, 1993). Pseudo rings are
purported to be reflective of an earlier stage in ring formation in galaxies,
as they are more common in late-type galaxies, while complete rings are more
common among early-type disk galaxies (Elmegreen et al. 1992; Buta 1993; Buta
& Combes 1996). The theoretical models of Schwarz (1981) however predict both
rings and pseudo rings to occur close to the Lindblad resonance radii in
galaxies. We present a more detailed discussion of resonance rings in Section
4.

2\. Driven by galaxy-galaxy collisions: Simulations have shown that galaxy
collisions could lead to the formation of ring galaxies (Appleton et al. 1996;
Wong et al. 2006; Elagali et al. 2018 and references therein). In a drop-
through collision, typically a smaller “bullet” galaxy passes through the
(off-)centre of a larger disk galaxy, sending out shock waves that push both
the stellar and gas material outward to form a ring-like structure. Many such
systems have thus far been observed including the famous Cartwheel galaxy AM
2026-424 (Zwicky, 1941), and the recently discovered high-redshift collisional
ring galaxy R5519 (Yuan et al., 2020).

3\. Driven by accretion of gas from the IGM: In few ring systems such as
Hoag’s galaxy (Hoag, 1950), the Leo ring (Schneider et al., 1983), AGC 203001
(Bait et al., 2020), the formation of the ring structure is attributed to the
accretion of fresh gas from the surrounding IGM onto a ring, under favourable
conditions (Lynds & Toomre, 1976; Theys & Spiegel, 1977; Schweizer et al.,
1987). Star formation starts as soon as the accreted H i gas column density
exceeds a threshold. In some cases, for example the Leo ring and AGC 203001
(Bait et al., 2020), the accreted H i gas column density is low enough that
significant star formation does not occur, leaving only an H i ring behind
with no stellar counterpart. Polar rings are another class of ring galaxies
which are believed to have accreted H i gas, but that gas has settled into
polar orbits around the galaxy. Typically, polar ring galaxies are observed to
be S0-type galaxies which have recently accreted fresh H i gas from either a
smaller companion dwarf galaxy or directly from the IGM. To date over 400
potential polar ring galaxy candidates have been identified, but only a dozen
or more galaxies have been confirmed to be polar rings (Moiseev et al. 2011;
Egorov & Moiseev 2019; Nishimura et al. 2022). As such, polar ring galaxies
are a rare class of ring galaxies.

Ring galaxies have been studied extensively in the optical, UV, NIR and FIR
bands. Early photometric studies revealed that the location of the rings were
hot-spots for active star formation, which was determined from their distinct
patchy blue distribution (Sandage 1961; Athanassoula & Bosma 1985). Buta &
Combes (1996) find that among the catalogue of bright nearby galaxies, outer
(pseudo)rings are found only for about 10% of the sample. Buta (1995) compiled
a comprehensive list of 3692 ring galaxies, with the objective of studying if
the rings in galaxies were linked to orbital resonances. They speculate that
the observed outer (pseudo)rings may be associated with the outer Lindblad
resonance, while the inner rings are associated with the Ultraharmonic
resonances. They suggest that measuring the location of the inner and outer
rings may be useful to indirectly measure the bar pattern speed. For a full
review of the various studies undertaken in the different wavelength regimes
we refer the reader to Buta & Combes (1996).

In terms of the H i gas in ring galaxies, most studies have focused on the
unresolved, global H i gas properties with a handful of high-resolution H i
studies limited to small samples. One of the first high-resolution study of
the H i gas distribution in a ring galaxy was undertaken by Bosma et al.
(1977) for the spiral galaxy NGC 4736, via synthesis imaging using the
Westerbork Radio Synthesis Telescope (WSRT). They find that the H i gas in
this galaxy is observed to trace not only the optically bright central
regions, but also the faint outer stellar ring. Following this, a few targeted
H i studies focusing on a few galaxies were conducted, most of which revealed
H i rings to be roughly co-located with the outer stellar rings (see for
example Knapp et al. 1984; van Driel et al. 1988; van Driel & Buta 1991). Some
early H i studies have also identified galaxies for which rings of H i gas are
found with no obvious optical ring-like structures (see for example Krumm et
al. 1985; Tilanus & Allen 1993). Overall, such studies have hinted at the
importance of the H i gas in forming the ring-like stellar structures observed
in galaxies.

Apart from a few targeted H i studies, to date, there is no large
comprehensive study of the resolved H i gas distribution and kinematics of
ring galaxies. H i gas is an excellent tracer of both internal and external
processes affecting galactic evolution. In addition, H i gas is the
fundamental source of molecular hydrogen (H2), which is eventually converted
to stars in low-z galaxies (Meurer et al., 2006). This link proves even more
useful in the context of tracing star formation in the different types of ring
galaxies. Therefore, such a study will be an important step towards
understanding how the H i gas in these systems is distributed and also enable
us to study the effects of the environment on such systems. This may be
particularly useful in the context of understanding ring formation in galaxies
linked to tidal interactions and external perturbations (e.g., Combes et al.
1990). In addition, from the gas kinematics we can measure the angular
momentum (AM) in the ring and examine the link between gas stability and the
role of AM in regulating star formation in the ring. H i studies of ring
galaxies are therefore very important to understand both the morphology and
kinematics of the rings, allowing us to understand the various processes (both
internal and environmental) driving their formation.

It is important to distinguish mechanisms that lead to the formation of ring
galaxies (visible in their optical and H i morphology) as discussed above
compared to H i “holes” and H i “depressions” at the centers of galaxies that
may be linked to the natural process of star formation where the H i gas is
converted to H2 and eventually stars (see Walter et al. 2008; Leroy et al.
2008; Murugeshan et al. 2019) and do not necessarily coincide with the optical
rings. In this paper, we introduce a sample of ring galaxies which have been
identified through their optical morphology to possess ring-like structures
(see Section 2.1 for more details on the selection criteria). We probe the
distribution of the H i gas in such systems and study their kinematics. In the
next sub-section we introduce this new survey of ring galaxies.

###  1.1 The HI in Ring Galaxies Survey (HI-RINGS)

In order to address the diversity of ring galaxy evolution in terms of the gas
kinematics, we have put together a new high-resolution H i survey of ring
galaxies – the H i in Ring Galaxies Survey (H i-RINGS), targeting different
types of ring galaxies, to study their H i and star formation properties using
a uniform and large sample of 24 galaxies. We have procured new Australia
Telescope Compact Array (ATCA) observations for 15 ring galaxies in addition
to compiling ATCA data from the archive for 9 galaxies. Our sample consists of
various types of ring galaxies including resonance, collisional and
interaction rings, enabling us to probe their H i gas distribution, resolved
star formation and star formation rate (SFR) surface densities, as well as
their kinematic properties, using a multiwavelength approach including optical
and infrared data.

Through this paper and the following series of papers, we hope to address
important questions pertaining to the formation and evolution of ring galaxies
in the local Universe. Below we list some of the main goals of the H i-RINGS
survey.

  * •

Test Lindblad’s resonance theory using robust rotation curves for a large
sample of ring galaxies. We will utilise the rotation curves to derive the bar
pattern speeds and predict the location of the resonance rings. In addition,
we aim to study the H i gas and SFR surface densities to probe the nature of
the distribution of both gas and star formation in our sample galaxies.

  * •

We aim explore the H i/H2 ratio in ring galaxies. It has been observed in
several collisional ring galaxies that the H i/H2 ratio is large, this could
mean that in such galaxies the H2 is being efficiently converted into stars
(Higdon et al. 2011; Wong et al. 2017). We intend to compare the H i/H2 ratio
among the different types of ring galaxies in our sample to understand the H i
to H2 conversion phase in such systems. More than half the galaxies in our
sample have high-resolution ALMA observations of CO. By complementing this
with high-resolution H i data, we will be able to study the H i/H2 ratios for
a large sample of spatially resolved ring galaxies.

  * •

Obreschkow et al. (2016) introduced the $f_{\mathrm{atm}}-q$ relation, where
$f_{\mathrm{atm}}$ is the neutral atomic gas fraction and $q\propto
j_{\mathrm{b}}/M_{\mathrm{b}}$ is the integrated disk stability parameter,
where $M_{\mathrm{b}}$ is the total baryonic mass and $j_{\mathrm{b}}$ is the
total specific baryonic angular momentum (sAM) of the galaxy. Both H i-excess
and -deficient late-type galaxies have been found to follow this relation
consistently (Lutz et al., 2017, 2018; Murugeshan et al., 2019), indicating
that sAM plays a fundamental role in regulating the H i gas content in
galaxies. Murugeshan et al. (2019) find two ring galaxies in their sample that
follow the $f_{\mathrm{atm}}-q$ relation. However, do all galaxies possessing
rings that are in gravitational equilibrium follow the $f_{\mathrm{atm}}-q$
relation? We will use the high-resolution H i data to accurately model the
kinematics of ring galaxies, measure their sAM and study their behaviour on
the $f_{\mathrm{atm}}-q$ plane. Such a study will give us important clues
about the role of angular momentum in regulating the stability of the H i gas
in the ring.

  * •

What is the effect of the local and global environment on the formation of
ring galaxies? Elagali et al. (2018) found that 83% of ring galaxies in the
EAGLE simulations were formed via a ‘drop-through’ collision. Our sample of
ring galaxies spans across a range of local density environments, being
associated with the field, group or cluster. With our sample of ring galaxies,
we aim to study how the environment affects the properties of ring galaxies.

In this work, we introduce the H i-RINGS survey and give an overview of the
sample. We model the H i kinematics for the sample galaxies using 3D tilted-
ring fitting and derive robust rotation curves. Using the rotation curves we
accurately derive the bar pattern speed for the barred galaxies in our sample
and test Lindblad’s theory of ring formation. In addition, we also explore the
effects of the bar on the resolved and global H i and star formation
properties of galaxies. We structure the paper as follows – in Section 2 we
give details on the sample selection, observations, data reduction, source
finding and 3D modelling. Following this, in Section 3 we present the main
results, including an overview of the sample and their measured properties. In
Section 4, we test resonance theory of ring formation based on the analysis of
a sub-sample of barred galaxies in our sample. Finally, in Section 5, we
summarise the main results and give an outline of the plans for the future.

##  2 Data and Methods

In this section we will discuss the criteria applied for the sample selection,
give details of the observations, list the steps followed in the data
reduction process and discuss the source finding and characterisation. We will
also describe the kinematic modelling for the sample galaxies which is
necessary to derive their rotation curves.

###  2.1 Sample selection

We use the Catalogue of Southern Ringed Galaxies (CSRG, Buta 2001), obtained
through the Vizier catalogue access tool (Ochsenbein et al., 2000) to select
our sample of ring galaxies. The CSRG was originally published by Buta (1995),
presenting properties of outer and inner rings, pseudo-rings and lenses of
3692 galaxies south of -17∘ in declination.

Our selection criteria was based on galaxy size, position, visual inspection
of each galaxy from CSRG, and a cross-check as to whether the galaxies were
detected in the H i All Sky Survey (HIPASS, Barnes et al. 2001; Meyer et al.
2004):

  * •

Our primary selection criteria was galaxy size, selecting galaxies with large
angular size, sufficient to resolve the H i in the rings, to be able create
reliable kinematic models. We selected galaxies with D${}_{\textrm{25}}$
$\geq$ 180′′, where D${}_{\textrm{25}}$ is the B-band 25 mag arcsec-2
isophotal diameter, from ESO-LV ”Quick Blue” IIa-O (Lauberts & Valentijn,
1989a).

  * •

The secondary constraint was declination of the galaxy, selecting galaxies
south of -25∘ in declination. This criteria ensures good observability with
the ATCA, avoiding an elongated synthesised beam.

  * •

Using Astroquery SkyView cutout service (Ginsburg et al., 2019a), we obtained
DSS images of all $\sim$3692 galaxies from the CSRG and visually inspected
them. This step was necessary to validate the galaxy selection process and it
was done before and after size/declination restrictions to the catalogue.
Next, with visual inspection of 101 ring galaxies – obtained after the
declination and size restrictions – we removed highly edge-on galaxies and
galaxies with no clear rings and/or inner (pseudo)rings. Galaxies with
unclear/faint images were re-checked in NED database. Additionally, two ring
galaxies – ESO 215-31 and ESO 179-IG-013, were added to our sample through
visual inspection.

  * •

As a final check, we inspected the HIPASS data to see if the selected galaxies
have an H i detection. This step was necessary to minimize risk of not
detecting H i through our new ATCA observations. With this check five galaxies
were removed due to H i non detection in HIPASS.

The above criteria left us with a sample of 27 ring galaxies of the $\sim$3692
from the original CSRG sample, as the majority of galaxies have an angular
size too small to be sufficiently resolved with the ATCA. Out of the 27, 12
galaxies were found to have archival ATCA data (in the 1.5km, 750m and
352/375m configurations). We have observed 15 ring galaxies with our ATCA
proposal (C3385). Ultimately, due to calibration issues and/or not being able
to derive reliable kinematic models, 3 of the 12 galaxies from the archive
were rejected, leaving us with a final sample size of 24 ring galaxies. Fig. 1
shows the spatial and velocity distribution of our sample of ring galaxies.

Figure 1: Spatial and velocity distribution of galaxies in the H i in Ring
Galaxies Survey (H i-RINGS). The grey points show all the ring galaxies from
the (Buta, 1995) catalogue, the stars show our sample of ring galaxies where
colour denotes their Heliocentric velocity.

###  2.2 Observations

The 15 galaxies in our sample, for which no archival data are available, were
observed with the ATCA in one of the default 750m and 1.5km array
configurations. These observations were carried out through the ATCA project
C3385 in 2020OCTS and 2021APRS semesters with a total observation time of
307.5h, and an average on-source time of $\sim 11$h per array configuration.
Combining the data from these two configurations will result in increasing
both sensitivity and spatial resolution of the synthesised data, which is key
to accurately modelling the kinematics of the sample galaxies (see Section
2.7). The ATCA, which is equipped with the Compact Array Broad-band Backend
(CABB, Wilson et al. 2011), allows us to observe the galaxies in both
continuum and spectral-line mode. The total bandwidth in the continuum band is
2048 MHz with a frequency resolution of 1 MHz, centered at 2.1 GHz. The
spectral line observations have a total bandwidth of 8.5 MHz with a channel
resolution of $0.5$ kHz, which corresponds to a velocity resolution of $\sim
1$ km s-1. We employed the following strategy for the observations – 55 min on
source followed by 5 min on the phase calibrator each hour. The primary
calibrators used are PKS 1934-638 and PKS 0823-500 depending on the time of
observations. The primary calibrator was observed for 10 min at the start and
end of the observing run. Other details of the observations for the sample are
listed in Table 1.

###  2.3 Archival ATCA data

As mentioned previously, 9 galaxies in our sample have archival ATCA data in
both the standard 1.5km and 750m array configurations. The data was retrieved
from the Australia Telescope Online Archive (ATOA). For a few galaxies we also
found observations in the 350m array configuration. The exact configurations
used for the data reduction for the various galaxies in our sample is listed
in Table 1.

Table 1: Details of the observations. Col (1): Common name; (2) ATCA baseline observation configuration; Col (3): The central frequency of the observation; Col (4): Primary flux calibrator used; Col (5): Phase calibrator used; Col (6): rms in the image cubes; Col (7): Synthesized beam size major and minor axis, respectively; Col (8): Notes about observations: This survey - ATCA project C3385; NGC 1433 data is from C305 (Ryder et al., 1996); NGC 7531 data is from C2196 (Koribalski et al., 2009); NGC 1808 data is from C585 (Dahlem et al., 2001); NGC1533 data is from C934 and C1003 (Ryan-Weber et al., 2003); NGC 7098 data is from C744 (Pisano et al., 2002) Name | Array | $\nu_{0}$ | Primary | Secondary | rms | Synthesized beam | Notes
---|---|---|---|---|---|---|---
| Configurations | [MHz] | Calibrator | Calibrator |  [mJy beam-1] | [arcsec] |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8)
ESO 215-31 | 750m & 1.5km | 1407 | 1934-638 | 0823-500 | 0.67 |  43$\times$30 | This survey
ESO 269-57 | 750m & 1.5km | 1405 | 1934-638 & 0823-500 | 1320-446 | 2.20 |  43$\times$30 | This survey
NGC 1326 | 750m & 1.5km | 1413 | 1934-638 | 0332-403 | 1.06 |  93$\times$25 | This survey
NGC 3358 | 750m & 1.5km | 1406.3 | 1934-638 | 1015-314 | 1.85 |  54$\times$31 | This survey
IC 5267 | 750m & 1.5km | 1412.2 | 1934-638 | 2311-452 | 1.09 |  49$\times$29 | This survey
NGC 1302 | 750m & 1.5km | 1412 | 1934-638 | 0237-233 | 1.53 |  73$\times$28 | This survey
NGC 1398 | 750m & 1.5km | 1413 | 1934-638 | 0413-210 | 1.26 |  68$\times$27 | This survey
NGC 2217 | 750m & 1.5km | 1412 | 1934-638 | 0614-349 | 2.13 |  74$\times$30 | This survey
NGC 7020 | 750m & 1.5km | 1405.2 | 1934-638 | 2117-642 | 1.67 |  38$\times$30 | This survey
NGC 1291 | 750m & 1.5km | 1416 | 1934-638 | 0332-403 | 1.62 |  48$\times$29 | This survey
NGC 1371 | 1.5km | 1413.2 | 1934-638 | 0413-210 | 2.58 |  50$\times$19 | This survey
ESO 179-IG-013 | 750m & 1.5km | 1416 | 1934-638 | 1613-586 | 2.09 |  36$\times$29 | This survey
NGC 1079 | 1.5km | 1413.2 | 1934-638 | 0237-233 | 2.41 |  57$\times$19 | This survey
NGC 5101 | 1.5km | 1411.7 | 1934-638 | 1308-220 | 2.71 |  54$\times$19 | This survey
IC 5240 | 750m & 1.5km | 1411.7 | 1934-638 | 2226-411 | 0.99 |  48$\times$28 | This survey
NGC 1350 | 750m & 1.5km | 1415.5 | 1934-638 | 0237-233 | 0.94 |  57$\times$23 | Murugeshan et al. (2019)
NGC 2369 | 750m & 1.5km | 1409 | 0823-500 | 0637-752 | 2.33 |  31$\times$24 | Murugeshan et al. (2019)
NGC 7098 | 375m & 750m | 1410 | 1934-638 | 2142-758 | 4.09 |  74$\times$58 | C744
NGC 1543 | 750m & 1.5km | 1419 | 1934-638 | 0420-625 | 1.25 |  37$\times$29 | Murugeshan et al. (2019)
NGC 1533 | 375m, 750m & 1.5km | 1417 | 1934-638 | 0407-658 | 8.2 |  35$\times$27 | C934 and C1003
NGC 1808 | 750m & 1.5km | 1414 | 1934-638 | 0521-365 | 1.00 |  50$\times$30 | C585
NGC 7531 | 352m & 1.5km | 1412.8 | 1934-638 | 2311-452 | 3.29 |  91$\times$62 | C2196
NGC 1433 | 750m & 1.5km | 1415 | 0438-436 | 0438-436 | 1.32 |  46$\times$34 | C305
NGC 6300 | 750m & 1.5km | 1415 | 1934-638 | 1718-649 | 0.96 |  36$\times$22 | Murugeshan et al. (2019)

###  2.4 Data Reduction

We used the MIRIAD software package (Sault et al., 1995) to reduce the ATCA
data. The 750m and 1.5km array data (and for some galaxies the 375m and 352m
data) were reduced separately and later combined to produce the final
visibility data set. The following steps are carried out in the data reduction
process – first we flag any spurious data using UVFLAG and PGFLAG. Once the
quality of the data looks good, we then perform the bandpass and phase
calibration using the tasks MFCAL, MFBOOT, GPCAL and GPBOOT. This is then
followed by continuum subtraction using UVLIN. We iteratively select line-free
channels in the data for the continuum subtraction process. Following this, we
use the INVERT task to fourier transform the UV data to form a dirty image.
The visibility data from the 1.5km and 750m configurations are combined using
the INVERT task after continuum subtraction has been carried-out on the
individual data sets. The INVERT task also estimates the expected theoretical
rms noise in the image cube. We use the 3$\sigma$ value of this estimate as
cutoff during the “cleaning” process. Next, we proceed to the de-convolution
step, which is performed using the CLEAN algorithm in MIRIAD. We set the
number of iterations for the minor CLEAN cycle to 2000 and wait for CLEAN to
reach the cutoff value. Following the de-convolution step, we restore the
final image cube using the RESTOR task. Finally, we apply the primary beam
correction to the restored image cube using the LINMOS task.

We set the robustness parameter to 0.5, as it offers a good compromise between
sensitivity and resolution. The semi-major axis values of the synthesized beam
(BMAJ) range from 31 to 93 arcsec, with a median beam size of $\sim 40$
arcsec. The channels were also smoothed to a width of 5 km s-1, thereby
increasing the signal-to-noise-ratio (SNR) in the image cubes. The median
3$\sigma$ rms noise level in the image cubes for our galaxy sample is $\sim
4.92$ mJy/beam per channel, which corresponds to a column density sensitivity
limit of N${}_{\textrm{H\,{\sc i}}}\sim 1.7\times$1019cm-2 assuming an
observed median beam size of $\sim 40$ arcsec.

###  2.5 Source Finding with SoFiA

We make use of the Source Finding Application (SoFiA; Serra et al. 2015;
Westmeier et al. 2021) to automatically detect sources in the final image
cubes. We have used the smooth + clip (S + C) algorithm in SoFiA to build the
source mask. The S + C finder works by spatially and spectrally smoothing the
data cube on multiple scales as defined by the user. In each smoothing
iteration, a user-specified flux threshold relative to the global rms noise
level after smoothing is applied to the data, and all pixels with an absolute
flux density exceeding that threshold is used to create the source mask
(Westmeier et al., 2021). The multi-scale smoothing and clipping helps in
identifying relatively faint and extended features in the data cubes that may
have been missed otherwise. After the source finding run, SoFiA returns a
number of data products including a catalogue with source-specific parameters
such as their RA, Dec, systemic velocity (V${}_{\textrm{sys}}$), integrated
flux (S${}_{\textrm{int}}$), $w_{20}$ and $w_{50}$ velocity widths and
kinematic position angle (P.A) based on the H i gas kinematics. In addition,
SoFiA also produces cutouts of the spectral cube, moment 0 (intensity), moment
1 (velocity) and moment 2 (dispersion) maps for all detected sources from the
inputed image cube. We make use of the spectral cubes and moment maps derived
from SoFiA for the rest of our analysis.

###  2.6 Interactive visualisation of the HI-RINGS galaxies

Using Streamlit, an open-source Python library, we have built a web
application for interactive visualization of our H i-RINGS survey111Link to
the interactive web app: https://hi-rings.streamlitapp.com/. The web
application allows the user to interactively explore our sample of ring
galaxies using H i intensity data together with data from several other
surveys including the DSS, WISE, 2MASS and GALEX. Within the app, it is
possible to overlay selected H i contours on top of image cutouts from the
different surveys. The app also allows image comparison with and without H i
contours, and image comparison with different survey images in the background.
We believe such an interactive tool will be both informative as well as
provide important visual context into the distribution of the H i and stellar
components of the different types of ring galaxies in our sample. We provide
further details on the web application in A. In addition, we also intend to
host the reduced H i data-products for the H i-RINGS sample.

###  2.7 Tilted Ring Modelling with 3DBarolo

We perform 3D tilted-ring fitting (Rogstad et al., 1974) to the H i emission
observed in the data cubes for our sample galaxies using the publicly
available software 3DBarolo (Di Teodoro & Fraternali, 2015). 3DBarolo takes in
a user-specified parameter file containng initial guesses for the fitting and
outputs the best fit tilted-ring model for the galaxy. For the fitting
strategy we follow the steps described in Murugeshan et al. (2019). 3DBarolo
fits $n$ independent tilted rings (where $n$ is the total number of resolution
elements across the major axis) of width equivalent to half the FWHM of the
resolution element (which varies from galaxy to galaxy in our sample). The
fitting is performed in two stages – in the first stage the user-defined
initial guesses for the galaxy centre (set to the optical/NIR centres obtained
from NED222http://ned.ipac.caltech.edu/), the systemic velocity
(V${}_{\textrm{sys}}$) and position angle (PA) are set to the values
determined from the SoFiA source finding run. The initial guess for the
inclination angle ($i$) is set to the value determined by Lauberts & Valentijn
(1989b) based on the $B$-band optical surface photometry of the galaxies. We
allow PA, $i$ and the rotation velocity ($v_{\textrm{rot}}$) to vary as free
parameters during the fit. Once the best fitting model is derived from the
first stage, the code uses these values in the second stage of the fitting.
Finally, the best fitting values for PA, $i$ and $v_{\textrm{rot}}$ for every
ring is written-out to a log file. The best fit model parameters such as
$v_{\textrm{rot}}$, $i$ and PA are shown for the full sample in Fig. 16 and
Fig. 17 in B.

We do note that since we are dealing with ring galaxies and deriving their
rotation curves from resolved H i observations, there is the danger of missing
the H i signal from the very centres of these galaxies (assuming the H i gas
is primarily co-located with the optical ring). Upon investigation, we find
that this is not the case for the majority of our sample galaxies. We find
that in most cases there is still H i emission from their centres and that the
signal-to-noise ratio (SNR) of the H i emission is high enough for 3DBarolo to
accurately fit the rotation curve in the inner regions. We find four galaxies
in our sample, namely, NGC 1326, NGC 1371, NGC 5101 and NGC 1534, for which
the SNR of the H i emission at the centre is too low for 3DBarolo to
effectively create a source mask and fit a kinematic model. We note that this
might potentially affect the resonance analysis we perform in Section 4 for
the above galaxies. This is because the prediction of the resonance locations
are dependent on their rotation curves. That being said, since the location of
the resonance rings are mainly dependent on the outer parts of the rotation
curve, we believe this should not significantly affect our analysis.

##  3 Results

We now discuss some general properties of the ring galaxies in our sample as
well as present the main results.

###  3.1 Flux comparison with HIPASS

Figure 2: The plot compares the integrated flux for the ring sample derived
from our ATCA observation (on the Y-axis) against their corresponding
integrated flux from HIPASS (X-axis). The dashed black line represents the
one-to-one correspondence.

As described in Section 2.5, we use SoFiA to automatically detect the sources
in the image cubes and also obtain a catalogue consisting of relevant H i
parameters including the integrated flux. We check the quality of the data and
as a validation step compare the integrated flux derived from the SoFiA run
for our sample galaxies with their HIPASS flux values. Fig. 2 shows a plot
comparing the integrated flux from the ATCA observations with their
corresponding HIPASS fluxes. For the majority of the sources, we find a good
one-to-one correspondence. About 7 galaxies in our sample ($\sim 29\%$) are
observed to have a low integrated flux compared to their HIPASS flux. We
attribute this to a combination of two effects: 1) Missing diffuse flux –
Three galaxies NGC 1371, NGC 1079 and NGC 5101 only have observations in the
1.5km array configuration. This means that these galaxies are very likely
missing the extended diffuse emission that the shorter baselines may have
otherwise captured. This is also evidenced by the fact that the HIPASS
$w_{20}$ values are typically $\sim 9\%$ broader than the ATCA $w_{20}$
values; 2) Confusion in the HIPASS flux – Galaxies NGC 1326 and NGC 3358 are
observed to have close companions spatially and/or in velocity space. Owing to
the larger $\sim 15$ arcmin HIPASS beam, it is likely that the HIPASS flux is
confused and overestimated for these two galaxies. This is also clearly seen
in Fig. 18, where we compare their HIPASS and ATCA spectra. We examine the rms
noise level in the image cubes and find them to be between 0.67 – 8.2 mJy/beam
per channel. This is close to the expected theoretical noise given the
integration time and UV-coverage of the individual observations. In the
following sections we present the main results from our study.

###  3.2 Sample Overview and HI properties

In this section we give an overview of both the global and resolved properties
of the ring galaxies in our sample. Figs. 7 – 10 in B show the optical images
of the sample galaxies with the H i contours overlaid on top, as well as their
H i velocity maps. From the optical images, we find that most of the galaxies
in our sample are early late-type galaxies with their morphology ranging from
(SA, SAB to S0). In addition to the presence of optical rings, we also notice
the H i gas distributed in the form of rings, in many cases (almost) co-
located with the optical rings (NGC 1326, NGC 1398, NGC 2217, IC 5240, NGC
1350, NGC 6300, ESO 215-31, NGC 1079, NGC 5101). We find this to be true
typically for galaxies possessing bars in our sample. However, it is also
interesting to note that the H i gas distribution in some ring galaxies is
more complex. For example, galaxies ESO 215-31 and NGC 1808 are observed to
have an excess of H i gas clumping at the ends of the bar and along the major
axis of the bar, which may be indicative of gas inflows (see for example Salak
et al. 2016 and references therein). From the velocity maps it is evident that
most galaxies in our sample exhibit ordered rotational motion.

We make use of the Carnegie-Irvine Galaxy Survey (CGS; Ho et al. 2011) and the
Third Reference Catalogue of Bright Galaxies (de Vaucouleurs et al., 1991) to
confirm presence or absence of a bar for our sample galaxies. We visually
observe that 14 galaxies in our sample have a strong bar, while the other 10
are observed to have a weak or no bar. For our analysis in the next sections,
we therefore split the galaxies into two groups – those with a strong bar and
those with a weak and/or no bar. The H i mass for our sample galaxies is
measured using the relation described in Walter et al. (2008),

$M_{\mathrm{HI}}\left[\mathrm{M}_{\odot}\right]=2.356\times
10^{5}D^{2}S_{\text{int }}$,

where $D$ is the distance in Mpc and $S_{\text{int }}$ is the ATCA integrated
flux in units of Jy km s-1. The uncertainties on $M_{\mathrm{HI}}$ are
computed by propagating the errors associated with $S_{\text{int }}$ as
described in Koribalski et al. (2004) and assuming a 10% uncertainty on the
distance measures. The H i mass of the sample galaxies is in the range
$8.61\leq\log(M_{\textrm{H\,{\sc i}}}/M_{\odot})\leq 10.19$, with a median H i
mass of $\log(M_{\textrm{H\,{\sc i}}}/M_{\odot})$ $\sim 9.36$.

We estimate the stellar mass for our sample using their 2MASS (Skrutskie et
al., 2006) $K_{s}$-band magnitudes as described by Eq. 3 in Wen et al. (2013).
The derived stellar mass for our sample is in the range
$9.77\leq\log(M_{\star}/M_{\odot})\leq 11.07$, with a median value of
$\log(M_{\star}/M_{\odot})\sim 10.56$. In Table 3 we list all the relevant
derived quantities such as the integrated flux, H i and stellar mass etc, for
our sample.

(a)

(b)

Figure 3: (a) The H i gas fraction vs the stellar mass of the ring galaxies
compared with the xGASS galaxies (blue points). The green stars represent
galaxies in our sample with a weak or no bar and orange stars represent
galaxies with a strong bar (b) The M${}_{\small\textrm{H\,{\sc i}}}-$ M⋆
relation for the rings sample, this time with the galaxies colour-coded in
terms of their $\Sigma_{5}$ local environment density values. The dark purple
line traces the rolling median of the xGASS galaxies with the shaded region
representing the 1$\sigma$ scatter.

Fig 3(a) shows the H i gas fraction plotted against the stellar mass for the
ring galaxy sample ($f_{\textrm{H\,{\sc i}}}-M_{\star}$ relation). As
mentioned, we segregate our ring galaxy sample into two – those without bars
and/or having a weak bar (green stars) and those observed to have a strong bar
(orange stars). We compare our ring galaxy sample with the representative
xGASS sample (Catinella et al., 2018) (blue points) in order to infer how the
H i gas fraction of the ring galaxies relates to their stellar mass. We
observe that 11 ($\sim 46$%) of our sample galaxies appear to be H i-deficient
compared to the xGASS sample (i.e., lying below the $1\sigma$ scatter about
the running median of the xGASS sample). In addition, we find that 9 of the 11
($\sim 82$%) galaxies identified to be H i-deficient are observed to have a
strong bar.

We quantify the degree of H i-deficiency of the barred and unbarred galaxies
in our sample using the xGASS galaxies as a representative sample. We compute
the H i deficiency as follows

DEF${}_{\textrm{H\,{\sc i}}}=\log(M_{\textrm{H\,{\sc
i},exp}})-\log(M_{\textrm{H\,{\sc i},obs}})$

Where $\log(M_{\textrm{H\,{\sc i},exp}})$ is the expected H i mass obtained
from the rolling median of the $M_{\textrm{H\,{\sc i}}}-M_{\star}$ scaling
relation for the xGASS sample, and $\log(M_{\textrm{H\,{\sc i},obs}})$ is the
observed H i mass calculated from our ATCA observations. We also measure the H
i deficiency values for the sample by replacing the ATCA-derived H i mass with
their HIPASS-derived H i mass. Higher the DEF${}_{\textrm{H\,{\sc i}}}$ value,
more H i-deficient the galaxy and vice-a-versa. In Table 2 we show the
computed median H i deficiency values for both the unbarred and barred
galaxies in our sample, and estimate the standard error on the computed median
values using a bootstrap re-sampling method. We find that the barred galaxies
are significantly more H i-deficient compared to the unbarred galaxies in our
sample.

Table 2: The median H i deficiency values computed for the unbarred and barred galaxies in our sample, using both their ATCA- and HIPASS-derived H i masses. The standard error on the median value is quoted within the brackets. |  DEF${}_{\textrm{H\,{\sc i}}}$ (ATCA) |  DEF${}_{\textrm{H\,{\sc i}}}$ (HIPASS)
---|---|---
Unbarred | -0.015 (0.131) | -0.087 (0.106)
Barred | 0.455 (0.154) | 0.298 (0.174)

H i deficiency in galaxies has been generally attributed to the effects of the
environment. In dense environments, a number of processes tend to strip the H
i gas from the discs of galaxies, thereby rendering them H i-deficient. Such
processes include, tidal interactions/stripping (Byrd & Valtonen 1990; Rots et
al. 1990), ram pressure stripping (Gunn & Gott 1972; Chung et al. 2009),
harassment (Moore et al., 1996) among others. To examine if environmental
processes are severely affecting the H i gas properties of the observed H
i-deficient ring galaxies in our sample, we measure the local environment
density for our sample galaxies. We make use of the 5th Nearest-Neighbour
projected environment density measure, $\Sigma_{5}$ (Mpc-2) (Baldry et al.,
2006), to probe the local density of the environment the sample galaxies are
residing in.

We follow the methods described in Murugeshan et al. (2020) to estimate the
$\Sigma_{5}$ values for our sample. The values are estimated based on the
2MASS Redshift Survey (2MRS) catalogue (Huchra et al., 2012). With more than
$\sim 43,000$ spectroscopic redshift measurements within the magnitude limits,
$K_{s}\leq 11.75$ mag and — b — ¿ 5o, the 2MRS catalogue is complete to 97.6%
and covers 91% of the entire sky. The below steps are followed to measure the
$\Sigma_{5}$ values:

  * •

First, we make the 2MRS catalogue volume-limited. We make a velocity cut to
the catalogue and include only galaxies with systemic velocities in the range
$200-8000$ km s-1. This is then followed by imposing an absolute $K_{s}$-band
magnitude cut of $M_{K}<-23.45$, corresponding to the survey’s limiting
apparent magnitude of 11.75 mag at the highest velocity edge (8000 km s-1),
thus making the 2MRS sample volume-limited.

  * •

The $\Sigma_{5}$ value is computed as $\Sigma_{5}=5/\pi D^{2}$, where $D$ is
the projected distance in Mpc to the 5th nearest-neighbour within
$\pm~{}500$km s-1of the target galaxy.

Fig. 3(b) shows the $f_{\textrm{H\,{\sc i}}}-M_{\star}$ relation for the ring
sample, but now the galaxies are colour-coded based on their $\Sigma_{5}$
values. We find that overall, the majority of our sample galaxies are from
lower-density environments such as the field and/or loose groups. However, few
galaxies in our sample are from high-density environments such as dense groups
and cluster. Based on the $\Sigma_{5}$ values and the H i and optical
morphology of the galaxies in our sample, we do not find any significant
evidence that the H i-deficiency is driven by environmental effects. This is
suggestive that other internal processes may be driving their H i-deficiency.
As mentioned earlier, since $\sim 82$% of the H i deficient galaxies in our
sample possess a strong bar, it may be the case that the bar is playing an
important role in driving the H i deficiency among the barred galaxies in our
sample.

The bar has been proposed to play a vital role in the evolution of galaxies,
specifically in the context of angular momentum (AM) transport and
redistribution of matter within the galactic disc (Sorensen et al. 1976;
Athanassoula & Bosma 1985; Knapen 2005b). Previous simulation studies have
shown that the gas in galactic discs experience torques due to the non-
axisymmetric potential of the bar, leading to the dissipation of energy due to
collisions in their orbit about the bar, consequently leading to loss in AM
and funneling towards the centre of the galaxy (Bournaud & Combes 2002; Knapen
2005a; Athanassoula 2013). In addition to driving gas to the centre within the
co-rotation radius, the bar also effectively transfers AM to the outer parts
of the disc, thereby increasing the stability of the H i gas in the outer
disc, which remains unavailable for collapse to form stars (Masters et al.,
2012). This duel effect i.e, funneling of the gas to the centre inducing star
formation and exhaustion of the H i gas (Lin et al. 2017), and the transport
of AM to the outer disc makes those galaxies having strong bars H i deficient
and redder due to quenching (Masters et al. 2012; Newnham et al. 2020;
Fernandez et al. 2021). Our sample galaxies are consistent with these
observations, wherein we find that the majority of the strongly barred
galaxies are more H i-deficient compared to those galaxies with no bar or
weakly barred.

We note that there may be other potential processes driving the H i-deficiency
in galaxies other than the effects of the bar. This may include interaction
driven tidal perturbations that may lead to starbursts (e.g., Knapen & James
2009; Bergvall et al. 2016) and consequent depletion of the H i gas; effects
of star formation-driven outflows and the Active Galactic Nuclei (AGN) in
ionising and expelling some fraction of the accreted H i gas out of the disk
and into the circumgalactic medium (e.g., Debuhr et al. 2012; Harrison et al.
2014; Cheung et al. 2016). All of these processes may be contributing to the
observed H i-deficiency in galaxies. However, as noted earlier in this
section, a majority of the galaxies in our sample are from low to intermediate
environments and so the effects of tidal interactions is likely to be minimal.
In addition, only three galaxies in our sample (NGC 2217, NGC 1808 and NGC
6300) are identified to have any AGN activity. We do note that the current
lack of nuclear activity in the majority of the galaxies in our sample does
not necessarily mean the AGN was not turned on at an earlier time. This is
because AGN activity in galaxies tend to turn on and off at shorter timescales
$\sim 10^{5}$ yrs (Schawinski et al., 2015). Even with this caveat in mind, it
is unlikely that AGN-driven winds have played a significant role in driving
the H i deficiency among our sample of ring galaxies, mainly because the
galaxies in our sample are typically high-mass i.e.,
$\log(M_{\star}/M_{\odot})\geq 10$, while AGN-driven H i depletion is
generally observed to impact low-mass galaxies i.e.,
$\log(M_{\star}/M_{\odot})\leq 10$ (see for example Bradford et al. 2018 and
references therein). Thus, the majority of our sample is likely not severely
affected by AGN-driven H i deficiency. In the next section, we investigate the
effects of the bar on the H i gas distribution among our sample galaxies, as
well as how the bars may be pivotal in both funneling H i gas to the centres
of galaxies as well as driving resonance rings in the inner and outer parts of
the galactic disc.

Table 3: Ring galaxies derived sample properties. Col (1): Common name; Cols (2) - (3): Right Ascension and Declination in J2000; Col (4): Sytemic velocity; Col (5): Distance; Col (6): Integrated flux derived from SoFiA; Cols (7) - (8): H i and stellar mass; Col (9): 5th nearest-neighbour density estimate; Col (10): Whether the galaxy has a strong, weak bar or does not have a bar. We refer to The Carnegie-Irvine Galaxy Survey (CGS; Ho et al. 2011) and the Third Reference Catalogue of Bright Galaxies (de Vaucouleurs et al., 1991) to confirm presence or absence of a bar. Name | RA | Dec |  V${}_{\textrm{sys}}$ | D |  S${}_{\textrm{int}}$ | $\log(M_{\textrm{H\,{\sc i}}})$ | $\log(M_{\star})$ | $\Sigma_{5}$ | Barred
---|---|---|---|---|---|---|---|---|---
| [J2000] | [J2000] |  [km s-1] | [Mpc] |  [Jy.km s-1] |  [M⊙] |  [M⊙] |  [Mpc-2] |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) | (9) | (10)
ESO215-31 | 11h10m34.79s | -49d06m09.4s | 2722 | 35.20 |  12.50 $\pm$ 0.27 |  9.56 $\pm$ 0.12 |  10.45 $\pm$ 0.09 | 0.05 | Yes
ESO269-57 | 13h10m04.43s | -46d26m14.4s | 3106 | 43.24 |  38.5 $\pm$ 0.81 |  10.19 $\pm$ 0.10 |  10.79 $\pm$ 0.10 | 0.59 | Yes
NGC1326 | 03h23m56.40s | -36d27m52.8s | 1360 | 14.95 |  7.67 $\pm$ 0.26 |  8.61 $\pm$ 0.11 |  10.32 $\pm$ 0.08 | 1.59 | Yes
NGC3358 | 10h43m33.02s | -36d24m38.5s | 2988 | 38.30 |  6.00 $\pm$ 0.33 |  9.32 $\pm$ 0.15 |  10.80 $\pm$ 0.07 | 0.33 | No
IC5267 | 22h57m13.57s | -43d23m46.1s | 1712 | 21.26 |  20.1 $\pm$ 0.56 |  9.33 $\pm$ 0.06 |  10.59 $\pm$ 0.07 | 0.18 | No
NGC1302 | 03h19m51.18s | -26d03m37.6s | 1710 | 11.21 |  15.70 $\pm$ 0.44 |  8.67 $\pm$ 0.09 |  9.87 $\pm$ 0.08 | 0.32 | Yes
NGC1398 | 03h38m52.13s | -26d20m16.2s | 1396 | 19.78 |  27.70 $\pm$ 0.53 |  9.41 $\pm$ 0.07 |  11.01 $\pm$ 0.05 | 0.47 | Yes
NGC2217 | 06h21m39.78s | -27d14m01.5s | 1619 | 21.87 |  37.50 $\pm$ 0.80 |  9.63 $\pm$ 0.06 |  10.79 $\pm$ 0.07 | 0.05 | Yes
NGC7020 | 21h11m20.09s | -64d01m31.2s | 3201 | 29.40 |  4.96 $\pm$ 0.29 |  9.01 $\pm$ 0.14 |  10.48 $\pm$ 0.07 | 0.24 | No
NGC1291 | 03h17m18.59s | -41d06m29.0s | 839 | 4.40 |  91.60 $\pm$ 1.09 |  8.62 $\pm$ 0.03 |  9.89 $\pm$ 0.07 | 0.28 | Yes
NGC1371 | 03h35m01.34s | -24d55m59.6s | 1463 | 23.79 |  45.50 $\pm$ 1.01 |  9.78 $\pm$ 0.18 |  10.60 $\pm$ 0.09 | 0.7 | Weak
ESO179-IG013 | 16h47m20.05s | -57d26m24.6s | 842 | 10.90 |  104.00 $\pm$ 1.44 |  9.46 $\pm$ 0.09 |  9.77 $\pm$ 0.04 | 0.11 | No
NGC1079 | 02h43m44.34s | -29d00m12.1s | 1452 | 25.20 |  18.60 $\pm$ 0.63 |  9.45 $\pm$ 0.08 |  10.44 $\pm$ 0.07 | 0.12 | Weak
NGC5101 | 13h21m46.24s | -27d25m49.9s | 1868 | 17.23 |  20.00 $\pm$ 0.81 |  9.15 $\pm$ 0.06 |  10.60 $\pm$ 0.12 | 0.08 | Yes
IC5240 | 22h41m52.38s | -44d46m01.8s | 1765 | 25.37 |  13.20 $\pm$ 0.33 |  9.30 $\pm$ 0.08 |  10.61 $\pm$ 0.04 | 0.24 | Yes
NGC1350 | 03h31m08.12s | -33d37m43.1s | 1905 | 18.83 |  12.50 $\pm$ 0.31 |  9.02 $\pm$ 0.08 |  10.62 $\pm$ 0.05 | 1.07 | Yes
NGC2369 | 07h16m37.73s | -62d20m37.4s | 3240 | 35.30 |  17.00 $\pm$ 0.56 |  9.70 $\pm$ 0.13 |  10.62 $\pm$ 0.08 | 0.18 | No
NGC7098 | 21h44m16.12s | -75d06m40.8s | 2381 | 29.10 |  40.20 $\pm$ 1.15 |  9.90 $\pm$ 0.10 |  11.07 $\pm$ 0.04 | 0.22 | Weak
NGC1543 | 04h12m43.25s | -57d44m16.7s | 1176 | 17.19 |  11.80 $\pm$ 0.38 |  8.92 $\pm$ 0.07 |  10.60 $\pm$ 0.11 | 0.11 | Yes
NGC1533 | 04h09m51.84s | -56d07m06.4s | 790 | 20.20 |  65.54 $\pm$ 2.31 |  9.80 $\pm$ 0.08 |  10.54 $\pm$ 0.09 | 1.59 | Yes
NGC1808 | 05h07m42.34s | -37d30m47.0s | 995 | 9.29 |  89.05 $\pm$ 0.89 |  9.26 $\pm$ 0.05 |  10.16 $\pm$ 0.07 | 5.64 | No
NGC7531 | 23h14m48.50s | -43d35m59.8s | 1596 | 22.22 |  85.44 $\pm$ 1.63 |  10.00 $\pm$ 0.10 |  10.20 $\pm$ 0.09 | 0.18 | No
NGC1433 | 03h42m01.55s | -47d13m19.5s | 1076 | 9.47 |  34.60 $\pm$ 0.61 |  8.86 $\pm$ 0.07 |  10.03 $\pm$ 0.05 | 0.34 | Yes
NGC6300 | 17h16m59.47s | -62d49m14.0s | 1109 | 12.26 |  34.50 $\pm$ 0.52 |  9.09 $\pm$ 0.06 |  10.44 $\pm$ 0.04 | 0.13 | Yes

##  4 Bar-driven resonance rings

In this section, we study the effects of bars in driving the ring structure
among the secularly evolving ring galaxies in our sample. There is a lot of
evidence from both observations and models that suggest that bars can lead to
the formation of resonance rings (Schwarz 1981; Buta 1986; Buta & Crocker
1991; Sellwood & Wilkinson 1993). The non-axisymmetric gravitational potential
of a bar can trigger density waves that leads to orbital resonances where gas
and stars can accumulate to form ring-like structures. The location of these
orbital resonances (also called Lindblad resonances) depend primarily on the
bar patter speed ($\Omega_{\textrm{bar}}$ ), the angular velocity ($\Omega$)
and the radial epicyclic frequency ($\kappa$) of the gas and stars orbiting
the galaxy (Lindblad 1964; Buta 1999). The resonances occur at the radius
where the following condition is met:

$\Omega_{\textrm{bar}}$ $=\Omega\pm\kappa/m$

where $m$ is an integer. The most commonly observed resonances in galaxies are
the Outer Lindblad Resonance (OLR), with $\Omega_{\textrm{bar}}$
$=\Omega+\kappa/2$, the Inner Lindblad Resonance (ILR), where
$\Omega_{\textrm{bar}}$ = $\Omega-\kappa/2$, the inner Ultra-harmonic
Resonance (iUHR) $\Omega_{\textrm{bar}}$ $=\Omega-\kappa/4$ and the outer
Ultra-harmonic Resonance (oUHR), where $\Omega_{\textrm{bar}}$
$=\Omega+\kappa/4$.

We measure the location of the various resonance radii based on the semi-major
axis radius ($a_{\textrm{bar}}$) and the angular pattern speed of the bar
($\Omega_{\textrm{bar}}$ ). Various methods have been proposed in the
literature to measure the length of bars in galaxies including visual methods
based on where the isophotes diverge from the typical boxy shape of the bar
(e.g., Garcia-Barreto & Momjian 2022); bar radius defined as the radius at
which the intensity along the major axis of the bar drops to half the
intensity of the peak (Salo et al., 2015); as well as based on fourier
analysis (Laurikainen & Salo 2002; Buta et al. 2006).

We measure the bar size visually following the methods described in Garcia-
Barreto & Momjian (2022). As the stellar population of the bar is mostly
dominated by older stars which are better represented in the IR (see for
example Buta 1999; Eskridge et al. 2000; Herrera-Endoqui et al. 2015 and
references therein), we make use of the DSS IR-band image to visually
determine the radius of the bar ($a_{\textrm{bar}}$) for our sample of barred
ring galaxies. This involves plotting isophotal contours on the DSS IR image
and defining the bar radius as the distance from the centre of the galaxy to
the last isophote that follows a boxy contour (see Fig. 4). We measure the
length of the semi-major axis radius of the bar from the centre to each ends
of the isophote along the major axis of the bar and take the bar radius to be
the average of the two, we term this value
$a_{\textrm{bar}}^{\textrm{uncorr}}$. We compute the standard error on the
average value of $a_{\textrm{bar}}^{\textrm{uncorr}}$ and propagate this error
when computing the de-projected bar radius $a_{\textrm{bar}}$.

Figure 4: DSS IR-band image of the zoomed-in central region of the ring galaxy
NGC 5101, showing the bar. The yellow curves are the IR-band isophotal
contours overlaid on the image. The red cross denotes the centre of the galaxy
and the orange arrows represent the length of bar semi-major axis radius
towards each end of the bar along the major axis of the bar.

The measured average uncorrected bar radius
($a_{\textrm{bar}}^{\textrm{uncorr}}$) is then de-projected to the kinematical
major axis of the galaxy, to obtain the corrected bar radius
($a_{\textrm{bar}}$). Again, this is done following the method described in
Garcia-Barreto & Momjian (2022) as:

$a_{\textrm{bar}}=a_{\textrm{bar}}^{\textrm{uncorr}}\sqrt{\cos^{2}(\phi)+\left(\sin^{2}(\phi)/\sin^{2}(i)\right)}$.

Here $\phi$ is the absolute difference between the median kinematic position
angle (P.A) of the disk of the galaxy and the position angle of the bar and
$i$ is the inclination angle of the optical disc of the galaxy as described in
Section 2.7. $a_{\textrm{bar}}$ is then converted to a physical size in units
of kpc using the distance ($D$) to the source. We propagate the errors
associated with $a_{\textrm{bar}}^{\textrm{uncorr}}$, $\phi$ and assume a
nominal 10% error on the distance measurements to compute the errors on
$a_{\textrm{bar}}$. We list the $a_{\textrm{bar}}$ values in Table 4.

Once we determine the corrected/de-projected bar radius, we assume the ratio
of co-rotation radius of the bar and bar radius, $R\equiv
R_{\textrm{CR}}/a_{\textrm{bar}}=1$, so that
$R_{\textrm{CR}}=a_{\textrm{bar}}$. A number of previous works have shown that
the value of $R$ ranges from 1 to 1.6 (e.g., Debattista & Sellwood 2000; Guo
et al. 2019; Schmidt et al. 2019). While other works have shown that $\sim
90$% of the barred galaxies have $1.0\leq R\leq 1.4$ ( Athanassoula 1992;
Debattista & Sellwood 2000; Starkman et al. 2018). Therefore, setting $R=1$ is
a reasonable approximation to make for our analysis. Indeed, we find that
setting $R=1.2$, which is the average of the range of reported values does not
significantly affect our results.

Next, we derive the angular velocity profile ($\Omega_{\textrm{H\,{\sc i}}}$ )
from the H i gas for our sample galaxies using the H i rotation curve
($v_{\textrm{rot}}(r)$) extracted from 3DBarolo (see Section 2.7). The
observed rotation curve is fit with a simple analytic function introduced in
previous works (e.g., Boissier et al. 2003; Obreschkow et al. 2016) and of the
form,

$v_{\textrm{rot}}(r)=v_{\max}\left(1-e^{-r/r_{\text{max}}}\right)$,

where $v_{\max}$ is the maximum rotation velocity corresponding to the flat
part of the rotation curve and $r_{\text{max}}$ is the radius where the
rotation curve transitions to the flat part. Both $v_{\max}$ and
$r_{\text{max}}$ are left as free parameters during the fit. We use the best
fitting function describing the rotation curve to then derive the angular
velocity profile of the H i gas, $\Omega_{\textrm{H\,{\sc i}}}$
$=v_{\textrm{rot}}(r)/r$, where $v_{\textrm{rot}}(r)$ is the circular rotation
velocity in km s-1and $r$ is the radius in kpc. We note that the choice of the
analytic function used to fit the rotation curves was motivated on the basis
of its simplicity as well as the fact that we were able to obtain reasonably
good fits for our sample. Following this we measure the bar pattern speed,
$\Omega_{\textrm{bar}}$ . $\Omega_{\textrm{bar}}$ is the angular velocity
corresponding to the bar co-rotation radius $R_{\textrm{CR}}$ (see panel (c)
in Fig. 5). To determine the various Lindblad resonance radii, we first need
to derive the epicyclic frequency ($\kappa$), which can be approximated as
follows:

$\kappa(r)=\sqrt{4\Omega_{\textrm{H\,{\sc i}}}^{2}+2r\Omega_{\textrm{H\,{\sc
i}}}\left(d\Omega_{\textrm{H\,{\sc i}}}/dr\right)}$

Figure 5: (a) H i moment 0 intensity map of NGC 7098. The synthesised beam is
shown in the left bottom corner. (b) DSS composite optical image of NGC 7098.
The blue, green, pink and orange ellipses represent the location of the ILR,
iUHR, oUHR and OLR, respectively. (c) The plot shows the angular velocity
($\Omega_{\textrm{H\,{\sc i}}}$ ;black line) derived from the rotation
velocity of the H i gas as a function of radius for the galaxy NGC 7098. The
blue, green, pink and orange profiles represent the ILR, iUHR, oUHR and OLR
Lindblad resonance curves respectively. The vertical black dashed line
represents the measured bar co-rotation radius (Rc), the horizontal black
dotted line represents the bar pattern speed ($\Omega_{\textrm{bar}}$ ). The
blue, green, pink and orange vertical dashed lines denote the location of the
ILR, iUHR, oUHR and OLR resonance radii respectively. (d) Top panel: Radial
profile of the H i surface density of the galaxy NGC 7098. Bottom panel:
Radial profile of the SFR surface density.

Panel (c) of Fig. 5 shows a plot of the angular velocity profile
$\Omega_{\textrm{H\,{\sc i}}}$ as well as the Lindblad curves,
$\Omega_{\textrm{H\,{\sc i}}}-\kappa/2$ , $\Omega_{\textrm{H\,{\sc
i}}}-\kappa/4$ , $\Omega_{\textrm{H\,{\sc i}}}+\kappa/4$ and
$\Omega_{\textrm{H\,{\sc i}}}+\kappa/2$ for the galaxy NGC 7098 (black, blue,
green, pink and orange curves respectively). The location of the co-rotation
radius ($R_{\textrm{CR}}$) of the bar is represented by the black dashed
vertical line in panel (c) of Fig. 5. The bar pattern speed
($\Omega_{\textrm{bar}}$ ) is represented by the horizontal black dotted line.
The ILR, iUHR, oUHR and OLR radii are determined as the radius at which
$\Omega_{\textrm{bar}}$ intersects the $\Omega_{\textrm{H\,{\sc i}}}-\kappa/2$
, $\Omega_{\textrm{H\,{\sc i}}}-\kappa/4$ , $\Omega_{\textrm{H\,{\sc
i}}}+\kappa/4$ and $\Omega_{\textrm{H\,{\sc i}}}+\kappa/2$ curves
respectively. The position angle of the inner rings (ILR and iUHR) is set to
be equal to that of the bar position angle, as many previous studies have
shown that the inner rings are typically aligned parallel to the major axis of
the bar (e.g., Buta & Combes 1996). The position angle of the outer rings
(oUHR and OLR) are set equal to the median kinematic P.A of the galaxy.

Panels (a) and (b) in Fig. 5 show the H i moment 0 map and the optical DSS
image respectively of NGC 7098 with the ILR, iUHR, oUHR and OLR ring locations
indicated by the blue, green, pink and orange ellipses. NGC 7098 is observed
to posses an inner ring as well as an outer ring (see DSS optical image in
panel (b) of Fig. 5) and we find that the predicted location of the oUHR (pink
ellipse) coincides nicely with the observed inner ring and the predicted OLR
(orange ellipse) traces the outer ring in very good agreement. Additionally,
we also observe an H i ring located at the OLR radius.

Table 4: Properties of the bar and Lindblad resonances. Col (1): Common name; Col (2): Average length of the semi-major axis of the bar in kpc; Col (3): Bar patter speed in units of km s-1kpc-1; Cols (4) – (7): measured radius of the inner Lindblad resonance (ILR), inner Ultra-harmonic resonance (iUHR), outer Ultra-harmonic resonance (oUHR) and outer Lindblad resonance (OLR) respectively in units of kpc Name | $a_{\textrm{bar}}$ | $\Omega_{\textrm{bar}}$ |  R${}_{\textrm{ILR}}$ |  R${}_{\textrm{iUHR}}$ |  R${}_{\textrm{oUHR}}$ |  R${}_{\textrm{OLR}}$
---|---|---|---|---|---|---
| [kpc] |  [km s-1kpc-1] | [kpc] | [kpc] | [kpc] | [kpc]
(1) | (2) | (3) | (4) | (5) | (6) | (7)
ESO215-31 |  10.6 $\pm$ 1.1 |  10 $\pm$ 1 | NA | 6.0 | 22.5 | NA
NGC1326 |  3.2 $\pm$ 0.3 |  52 $\pm$ 5 | 1.91 | 2.5 | 4.6 | 6.8
NGC1398 |  5.2 $\pm$ 0.5 |  47 $\pm$ 3 | NA | NA | 11.4 | 21.5
NGC2217 |  7.5 $\pm$ 0.8 |  27 $\pm$ 2 | 3.4 | 4.6 | 13.8 | 25.1
NGC1371 |  3.9 $\pm$ 0.4 |  57$\pm$ 3 | NA | NA | 6.9 | 11.0
NGC1079 |  4.3 $\pm$ 0.4 |  26 $\pm$ 1 | NA | NA | 13.0 | NA
NGC5101 |  6.7 $\pm$ 0.7 |  31 $\pm$ 3 | NA | NA | 12.0 | NA
IC5240 |  3.9 $\pm$ 0.4 |  44 $\pm$ 4 | 2.2 | 3.0 | 6.0 | 9.6
NGC1350 |  5.3 $\pm$ 0.6 |  39 $\pm$ 4 | NA | 3.2 | 9.2 | 15.7
NGC7098 |  6.1 $\pm$ 0.6 |  47 $\pm$ 10 | 4.1 | 5.1 | 9.2 | 14.9
NGC1533 |  2.7 $\pm$ 0.3 |  32 $\pm$ 1 | NA | NA | 5.8 | 10.2
NGC6300 |  2.2 $\pm$ 0.3 |  90 $\pm$ 13 | NA | NA | 2.6 | 3.6

In addition, we derive the star formation rate (SFR) surface density profiles
for our sample to examine if the SFR is enhanced at the locations of the
resonance rings. A number of previous works have highlighted that star
formation is triggered at the resonance locations in the disc where the H i
gas accumulates. We derive the SFR surface density profiles for the galaxies
in our sample using their WISE $W_{3}$ and $W_{1}$ images procured from the
NASA/IPAC IRSA333https://irsa.ipac.caltech.edu/applications/wise/ image cutout
service. The $W_{3}$ flux is shown to be a good tracer of the SFR in galaxies,
however, this flux is also contaminated by the contribution from older stellar
population ( Jarrett et al. 2011; Cluver et al. 2017). We correct for this
following the methods described in Cluver et al. (2017) and subtract $\sim
15.8\%$ of the $W_{1}$-band flux from the $W_{3}$ flux.

We first process both the $W_{3}$ and $W_{1}$ images to remove foreground
stars. We run the source finding code SExtractor (Bertin & Arnouts, 1996) to
identify point-like sources in the images, following this we use the IRAF
(Tody, 1986) task imedit to replace the bad pixels of the star with random
pixels surrounding the star. In the next step we place an annulus around each
galaxy and measure the median local sky background values from both the
$W_{3}$ and $W_{1}$ images, before subtracting this from the respective
images. Next, we measure both the $W_{3}$ and $W_{1}$ flux within concentric
ellipses, with position angle set to the median kinematic position angle of
the galaxy, and convert the corrected $W_{3}$ flux to SFR following the
methods described in Cluver et al. (2017).

In Panel (d) of Fig. 5 we plot both the azimuthally averaged H i surface
density ($\Sigma_{\textrm{H\,{\sc i}}}$) and the SFR surface density profile
($\Sigma_{\textrm{SFR}}$) for NGC 7098. The inner region (¡ 10 kpc) of this
galaxy is observed to have a dip in H i surface density
($\Sigma_{\textrm{H\,{\sc i}}}$), while showing a peak in SFR surface density.
At the location of the resonance radii, there is an enhancement in
$\Sigma_{\textrm{H\,{\sc i}}}$ and $\Sigma_{\textrm{SFR}}$. This is also
observed in most of the barred galaxies in our sample (see B for more details)
and is consistent with the theory that bars drive resonances, where the disk
stars and H i gas accumulate, and the gas eventually collapses to form new
stars, thereby resulting in increased SFR at the resonance ring locations
(Buta & Combes 1996). Such ring-like structures are likely to remain in
equilibrium for a few Gyrs as the resonant locations are regions in the
galactic disc where the net torque is zero (e.g., Buta 1999; Bournaud et al.
2005).

We perform the above resonance model analysis for the rest of the barred
galaxies in our sample to see if we are able to predict the location of the
resonance rings consistently. We report that there is good agreement between
the predicted locations of the resonances and the presence of actual ring-like
structures (in H i and/or in the optical) for the majority of the galaxies. We
note that depending on the angular velocity profiles, bar radius and pattern
speed of the individual galaxies, we are unable to always measure all four
resonances. We list the measured bar radius ($a_{\textrm{bar}}$), pattern
speed ($\Omega_{\textrm{bar}}$ ) and various resonance radii for the sample in
Table 4. Figs. 11 – 15 in B show the outcomes of the resonance model analysis
for the other barred galaxies in our sample.

The measured values of the bar pattern speed ($\Omega_{\textrm{bar}}$ ) for
our sample range from $\sim$ 10 – 90 km s-1kpc-1, with a median value of $\sim
44$ km s-1kpc-1. This range is consistent with previous works that have
measured $\Omega_{\textrm{bar}}$ values in galaxies (e.g., Corsini 2011;
Williams et al. 2021). In addition, we found a few galaxies in our sample for
which the $\Omega_{\textrm{bar}}$ measurements were available in the
literature and find that overall the values are in good agreement. For
example, Moore & Gottesman (1995) use VLA H i observations to determine
$\Omega_{\textrm{bar}}$ of NGC 1398 and estimate a value $\sim 46$ km
s-1kpc-1, which is in very good agreement with our measured value of $47\pm 3$
km s-1kpc-1. Garcia-Barreto et al. (1991) measure a pattern speed value of
$\sim 60$ km s-1kpc-1 for NGC 1326, while we measure a pattern speed of $52\pm
5$ km s-1kpc-1. We measure $\Omega_{\textrm{bar}}$ $90\pm 13$ km s-1kpc-1 for
NGC 6300, which is in disagreement with the value measured by Ryder et al.
(1996), who estimate a bar pattern speed of $27\pm 8$ km s-1kpc-1. However, we
emphasise that the higher $\Omega_{\textrm{bar}}$ value we measure predicts
the existence of an OLR at a radius of $\sim 3.57$ kpc, which coincides with
the location of both an H i and stellar ring in NGC 6300 (see Fig. 13 in B).
While Ryder et al. (1996) report that assuming the lower bar pattern speed of
$\sim 27$ km s-1kpc-1 would result in the location of an OLR at radius between
8 – 15 kpc, where they find no H i or stellar ring. The fact that our
$\Omega_{\textrm{bar}}$ measurement for NGC 6300 better matches the resonance
model instills confidence in our analysis.

Overall, from our analysis, we find that the bar driven resonance theory
explains the location of the inner and outer rings that we observe among
galaxies in our sample. It is indeed promising to note that we are able to
reliably measure the location of the various Lindblad resonances using just
the H i rotation curve and a visual measurement of the size of the bar. These
results come in support of the important role played by the bar in not only
funneling gas to the centers of galaxies but also in driving resonances and
redistributing the angular momentum in galactic discs. As noted in Section
3.2, where we observed strongly barred galaxies in our sample to be in general
more H i-deficient compared to the weak and non-barred galaxies, it seems
evident that bars accelerate the evolution of galaxies via their ability to
both use-up large reserves of the H i gas as well as making gas in the
outskirts unavailable for star formation due to a kick in angular momentum,
thus driving quenching in galaxies.

##  5 Summary and Future

We have introduced the new H i in Ring Galaxies Survey (H i-RINGS), comprising
of 24 ring galaxies for which we have obtained high-resolution H i
observations using the Australia Telescope Compact Array. The ring galaxies
were selected on the basis of their large angular size, their declination
range and a detection in the HIPASS catalogue. Our sample consists of a
variety of ring galaxies including what have been regarded in the literature
as resonance, collisional and interaction rings. We presented an overview of
the sample, including their H i, stellar and star formation properties. In
addition, we also highlighted their H i morphology and kinematics.

We segregated our sample of ring galaxies into two groups – those with a
strong bar and those with a weak or no bar. We find that in general the
strongly barred galaxies are more H i-deficient compared to the weakly barred
or unbarred galaxies in our sample. We also probed the local environment
densities of our sample galaxies in addition to looking at their optical and H
i maps to examine if the environment is playing a prominent role in driving
the H i deficiency in the galaxies, but find no strong evidence for the same.
We therefore hypothesize as suggested by previous works (e.g., Masters et al.
2012) that the bars are potentially driving the observed H i deficiency, via
their ability to funnel copious amounts of H i gas within the co-rotation
radius of the bar to the centers of galaxies. In addition, the bars also
transport angular momentum to the outer parts of the disc making the H i gas
in the outskirts of galaxies gain a kick in angular momentum and remain stable
against collapsing to form stars. This combined effect potentially leads to
bar quenching in galaxies as noted in other works.

We also employed Lindblad’s resonance theory to test if we are able to predict
the radius of the major resonances for a sub-sample of barred galaxies. To
determine the resonance radii, we measured the bar radius based on a visual
method and combined with the derived rotation curves of the galaxies, we
measure the bar pattern speed ($\Omega_{\textrm{bar}}$ ) before computing the
location of the ILR, iUHR, oUHR and OLR radius. The measured
$\Omega_{\textrm{bar}}$ values range from 10 – 90 km s-1kpc-1 in good
agreement with previous works. We find that the predicted location of the
resonance radii are in good agreement with the observed physical locations of
the H i/optical inner and/or outer rings in our galaxies. Furthermore, we
derived the azimuthally averaged H i ($\Sigma_{\textrm{H\,{\sc i}}}$) and SFR
surface ($\Sigma_{\textrm{SFR}}$) density profiles for these galaxies and
found that there is an enhancement in either $\Sigma_{\textrm{H\,{\sc i}}}$,
$\Sigma_{\textrm{SFR}}$ or both corresponding to one or more of the
resonances.

In forthcoming papers, we aim to study the global and resolved scaling
relations for our sample of ring galaxies to see how the different types of
ring galaxies behave in the various scaling relations. In addition, we also
aim to examine the properties of the collisional and interaction galaxies in
our sample. These are interesting systems as the ring structure is believed to
be driven by external perturbations. Specifically, as part of the H i-RINGS
survey, we have high-resolution H i observations for a recently discovered
collisional ring galaxy ESO179-IG013, also called “Kathryn’s Wheel” (Parker et
al., 2015), which is touted to be the nearest collisional ring galaxy and as
such offers a great opportunity to probe its H i and star formation
properties. In the near future we aim to address most of the questions posed
in the introduction pertaining to the evolution of ring galaxies.

We would like to thank the anonymous referee for their feedback which
significantly improved the quality of the paper. We would also like to extend
our thanks to Genevieve Batten, Vaishali Parkash and George Heald for their
useful comments and contributions.

M.E.C. is a recipient of an Australian Research Council Future Fellowship
(project No. FT170100273) funded by the Australian Government.

The Australia Telescope Compact Array is part of the Australia Telescope
National Facility (https://ror.org/05qajvd42) which is funded by the
Australian Government for operation as a National Facility managed by CSIRO.
We acknowledge the Gomeroi people as the traditional owners of the Observatory
site. This paper includes archived data obtained through the Australia
Telescope Online Archive (http://atoa.atnf.csiro.au). This research has made
use of the NASA/IPAC Ex- tragalactic Database (NED), which is operated by the
Jet Propulsion Laboratory, California Institute of Technology, under contract
with the National Aeronautics and Space Administration. This research has made
use of the VizieR catalogue access tool, CDS, Strasbourg, France. The original
description of the VizieR service was published in A&AS 143, 23.This
publication makes use of data products from the Two Micron All Sky Survey,
which is a joint project of the University of Massachusetts and the Infrared
Processing and Analysis Center/California Institute of Technology, funded by
the National Aeronautics and Space Administration and the National Science
Foundation. This publication makes use of data products from the Wide-field
Infrared Survey Explorer, which is a joint project of the University of
California, Los Angeles, and the Jet Propulsion Laboratory/California
Institute of Technology, funded by the National Aeronautics and Space
Administration. Part of this work was performed on the OzSTAR national
facility at Swinburne University of Technology. OzSTAR is funded by Swinburne
University of Technology and the National Collaborative Research
Infrastructure Strategy (NCRIS). This work was written on the collaborative
Overleaf platform https://www.overleaf.com. This research has made use of
python https://www.python.org and python packages: astropy (Astropy
Collaboration et al., 2013, 2018), matplotlib http://matplotlib.org/ (Hunter,
2007), APLpy https://aplpy.github.io/, pandas (McKinney, 2010), Jupyter
notebook https://github.com/jupyter/notebook, NumPy http://www.numpy.org/ (van
der Walt et al., 2011), SciPy https://www.scipy.org/ (Jones et al., 2001) and
CMasher https://github.com/1313e/CMasher (van der Velden, 2020). This research
used Streamlit https://streamlit.io/.

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##  Appendix A Interactive web application to visualise the HI-RINGS sample

Using Streamlit, an open-source Python library for turning data scripts into
shareable web applications (https://docs.streamlit.io/), we have built a web
application for interactive visualization of our H i-RINGS survey: https://hi-
rings.streamlitapp.com/. In Fig. 6 we show a screenshot of the web interface
of the application. The application allows the user to interactively explore
our sample of ring galaxies using H i intensity data in combination with
images from the following surveys: DSS, DSS2 (Blue, Red and Infra-red) WISE
(3.4$\mu$m, 4.6$\mu$m, 12$\mu$m and 24$\mu$m), 2MASS (J, H, K) and GALEX (Far-
UV and Near-UV), obtained using the python library astroquery (Ginsburg et
al., 2019b). There are two options for data display within the web
application:

  * •

Display a survey image and H i contours: this option is displayed upon opening
the application. The user will see (by default) the H i contour at the
$1\times 10^{20}$ cm-2 column density level overlaid on top of the galaxy’s
DSS image, with an image radius of 9 arcmin. The user can then choose another
galaxy, and/or adjust the display properties.

  * •

Compare two images: this option allows the user to visually compare two
personally adjusted images of the selected ring galaxy by dragging the image
slider from the centre of the images to the left and right. Image comparison
is possible with and without overlaid H i contours, as well as comparison with
different survey images in the background. Fig. 6 shows these options (with H
i contours), where the left side shows the DSS image while the right side
displays the GALEX Far-UV image of the galaxy NGC 1808.

In addition to selecting data display, it is possible to make further
adjustments: image radius, display color options – gray-scale or custom colour
map and image stretch. Next, the user can adjust the H i contour display
options, choosing which density contours they would like to display and which
colour map scheme to be used for the contours. We believe such an interactive
tool will be both informative as well as provide important visual context into
the distribution of the H i and stellar components of the different types of
ring galaxies in our sample. By allowing users to personally adjust displays
we are extending data visualisation accessibility. Furthermore, a code of this
web application is publicly available at:
https://github.com/ringgalaxies/HIRingGalaxies. Note: As with any other
application, it is highly likely that this application is going to be upgraded
with time, and the user interface may differ to the current one shown in Fig.
6.

Figure 6: A screenshot showing the user interface of the H i-RINGS Streamlit
web application (https://hi-rings.streamlitapp.com/) that allows users to
visualise our H i-RINGS galaxy sample in a multi-wavelength approach. The main
available options are: to select a ring galaxy, to show a survey image and H i
contours and to compare two images. The user can select a background image
from the listed surveys and adjust its display e.g., image radius, color -
grayscale or custom colourmap and image stretch. The user can also adjust the
H i contour display options, choosing which H i density contours they want to
be displayed and which colourmap to be used for the displayed contours. These
options are available for both a single image display and for a comparison of
two images from two different surveys.

##  Appendix B Notes on individual galaxies, 3DBarolo fit parameters and HI
spectra

In this section we specify notes on individual galaxies in our sample. As
mentioned, the objective of the H i-RINGS survey is to get a complete view of
the properties of ring galaxies. Our sample include what appear to be
resonance rings as well as collisional rings, ring galaxies with and without a
bar. In addition, as mentioned earlier we also comment on galaxies that posses
a bar and those that do not and probe the various mechanisms that may be
driving their ring-like structure. We note that all morphological classes of
the sample galaxies are based on the Carnegie atlas of galaxies (Sandage &
Bedke, 1994). We also discuss the location of the various Lindblad resonances
for the barred galaxies in our sample using the resonance method that was
described in Section 4. Additionally, we also note the location of the H i
ring, optical ring, spiral arms and show the H i and SFR surface density
profiles.

Fig. 16 and Fig. 17 display the parameters obtained from the 3D fitting using
3DBarolo as described in Section 2.7. Each panel shows the rotation curve, the
radial variation of the inclination ($i$) and position angles (PA) for the
individual galaxies in our sample. In addition, Fig. 18 and Fig. 19 show the
extracted ATCA and HIPASS spectra for our sample of ring galaxies.

Table 5: Ring galaxies in our sample categorised into those that are likely to be secularly evolving, those that are identified as collisional/interacting and also segregated on the basis of the presence or absence of bars. Individual notes on interesting systems is also made |  Secular Rings |  Collisional/Interaction Rings
---|---|---
Barred |  Galaxies ESO 215-31, NGC 1302, NGC 1398, ESO 269-57, NGC 1291, NGC 5101, NGC1543, NGC 1350, IC 5240, NGC 1433 and NGC 6300 are observed to show very similar optical morphologies. We find no strong evidence of interactions or external perturbations of the H i gas for these galaxies and therefore consider these galaxies to be following (semi-)secular evolution. All these galaxies are observed to posses a prominent bar, inner and outer (pseudo-)rings. We note that the H i disc of NGC 6300 is much more extended than its stellar disc unlike the other galaxies in this category, where the H i gas is observed to trace the outer ring and is not much more extended than the ring structure. For a subset of the barred galaxies for which we were able to derive accurate rotation curves, bar lengths and bar pattern speeds, we also test the resonance theory to predict the location of the various Lindblad resonances. Figs. 11 – 15 show all the relevant plots and images. It is observed that for all galaxies there is a good agreement between the predicted location of the resonances and the presence of ring-like structures either in the optical and/or the H i. In addition, we also show that for most galaxies there is an enhancement in the azimuthally averaged SFR ($\Sigma_{\textrm{SFR}}$) and/or H i surface density ($\Sigma_{\textrm{H\,{\sc i}}}$) profiles corresponding to the radius of the resonances. |  Galaxies NGC 1326 and NGC 2217 are observed to have H i tail-like extensions (towards the north for NGC 1326 and towards the north-east and south-west for NGC 2217, respectively, indicating some external perturbation and/or interaction). To what degree the external perturbation has affected the ring formation is uncertain. NGC 1533 – while the optical morphology of this galaxy is observed to be that of an early-type, the H i gas distribution on the other hand is observed to be much more extended and is seen to be linked to two nearby smaller companion galaxies (IC 2038 and IC 2039) to the northwest. Ryan-Weber et al. (2004) note that the amount of extended H i gas in the tail leading to IC 2038/IC 2039 from NCG 1533 is significantly more than what is expected from the size and morphology of IC 2038/IC 2039. Thus, they conclude that the copious amounts of H i gas may have been stripped from a low surface brightness galaxy which is not visible in the optical images. Either way, it is certain that NGC 1533 has recently undergone an interaction and accredited the H i from a low surface brightness galaxy. The H i gas is observed to be distributed in the form a ring, with no optical counterpart.
Unbarred |  Galaxies NGC 7020, NGC 2369, NGC 7531 and IC 5267 are observed to not possess bars. However, from their optical images we find that many of them possess an inner ring and all of them are observed to possess an outer (pseudo-)ring. Galaxies NGC 1079, NGC 1371 and NGC 7098 are observed to have weak bars and outer rings. We note that the H i gas distribution in NGC 1371 is much more extended compared to its stellar disc. A Bright linear inclined disc structure is observed in NGC 1808, but not resembling a bar. The H i gas is seen to trace the narrow linear bright central stellar distribution as well as the outer pseudo-ring. None of these galaxies show any significant signs of environmental effects as evidenced in their H i kinematic maps. Therefore, we attribute them to be secularly evolving. The origin of their rings therefore is not fully clear. |  NGC 3358 – This galaxy has a bright central nucleus and two tightly wound spiral arms extend to form a ring-like outer structure. Corresponding to the outer optical ring is the H i ring of gas. The velocity map of this galaxy shows no obvious deviation from ordered rotational motion. We note that this galaxy is part of an interacting system of at least seven galaxies part of the NGC 3347 group (Makarov & Karachentsev, 2011). This galaxy does not appear to possess a bar, therefore the origin of the ring in this galaxy may be attributed to external tidal perturbations. ESO 179-IG013 – This is an interesting candidate as it appears to have undergone an interaction based on both the optical and the H i maps (See upper right panels in Fig. 8) and is described as a collisional ring galaxy by Parker et al. (2015), who term it “Kathryn’s Wheel”. In terms of the galaxy morphology ESO 179-IG013 is classified as SB(s)m. The multiwavelength observations of ESO 179-IG013 by Parker et al. (2015) find a ring of star formation based on their H $\alpha$ and optical observations (see Fig. 3 in their paper). While de Vaucouleurs et al. (1991) classified this galaxy to posses a bar, Parker et al. (2015) argue that this elongated central structure may not actually be a bar but a residual late-type disc. In addition, Parker et al. (2015) identify the system to be the nearest collisional ring to the Milky Way with a systemic velocity of $\sim 840$ km s-1. They also identify the smaller ‘bullet’ galaxy that dropped through ESO 179-IG013, and find its star formation to be enhanced as well. In terms of the H i gas properties, this galaxy is considered the 60th brightest in the sample of 1000 brightest HIPASS galaxies (Koribalski et al., 2004), with an integrated flux of $\sim 104$Jy km s-1, corresponding to an estimated H i mass of $\log(\textrm{H\,{\sc i}}/M_{\odot})\sim 9.46$ assuming a distance of 10.9 Mpc. A high-resolution H i follow-up of this galaxy has been lacking until now. As part of the H i-RINGS survey, we present for the first time high-resolution H i follow-up observations for ESO 179-IG013. The H i velocity map of the system shows clear signs of perturbation with H i tails extending towards the south-east and north-west ends of the galaxy
Figure 7: The panels to the left show the H i contours overlaid on top of the
composite DSS images of the ring galaxies. The contour levels are at the
N${}_{\textrm{H\,{\sc i}}}\sim$ 0.1, 0.5, 1, 5 and 7$\times 10^{20}$ cm-2. The
panels to the right show the moment 1 (velocity) maps of the H i gas in the
rings galaxies. The colour-bars show the velocity range for each individual
galaxy. Figure 8: Same as Fig. 7 Figure 9: Same as Fig. 7 Figure 10: Same as
Fig. 7 Figure 11: Top row: Image to the left shows the moment 0 H i intensity
map and image to the right shows the composite DSS image of the galaxy NGC
1326. Second row: Left panel shows the angular velocity
($\Omega_{\textrm{H\,{\sc i}}}$) profile in black, while the blue, green, pink
and orange profiles denote the various Lindblad curves. The radius of co-
rotation ($R_{c}$) is denoted by the black vertical dashed line. The blue,
green, pink and orange ellipses (and vertical dashed lines) always represent
the ILR, iUHR, oUHR and OLR locations. The bottom two rows are similar to the
first two rows but shown for the galaxy NGC 1398. Figure 12: Top row: Image to
the left shows the moment 0 H i intensity map and image to the right shows the
composite DSS image of the galaxy NGC 2217. Second row: Left panel shows the
angular velocity ($\Omega_{\textrm{H\,{\sc i}}}$) profile in black, while the
blue, green, pink and orange profiles denote the various Lindblad curves. The
radius of co-rotation ($R_{c}$) is denoted by the black vertical dashed line.
The blue, green, pink and orange ellipses (and vertical dashed lines) always
represent the ILR, iUHR, oUHR and OLR locations. The bottom two rows are
similar to the first two rows but shown for the galaxy IC 5240. Figure 13: Top
row: Image to the left shows the moment 0 H i intensity map and image to the
right shows the composite DSS image of the galaxy NGC 1350. Second row: Left
panel shows the angular velocity ($\Omega_{\textrm{H\,{\sc i}}}$) profile in
black, while the blue, green, pink and orange profiles denote the various
Lindblad curves. The radius of co-rotation ($R_{c}$) is denoted by the black
vertical dashed line. The blue, green, pink and orange ellipses (and vertical
dashed lines) always represent the ILR, iUHR, oUHR and OLR locations. The
bottom two rows are similar to the first two rows but shown for the galaxy NGC
6300. Figure 14: Top row: Image to the left shows the moment 0 H i intensity
map and image to the right shows the composite DSS image of the galaxy ESO
215-31. Second row: Left panel shows the angular velocity
($\Omega_{\textrm{H\,{\sc i}}}$) profile in black, while the blue, green, pink
and orange profiles denote the various Lindblad curves. The radius of co-
rotation ($R_{c}$) is denoted by the black vertical dashed line. The blue,
green, pink and orange ellipses (and vertical dashed lines) always represent
the ILR, iUHR, oUHR and OLR locations. The bottom two rows are similar to the
first two rows but shown for the galaxy NGC 1079. Figure 15: Top row: Image to
the left shows the moment 0 H i intensity map and image to the right shows the
composite DSS image of the galaxy NGC 5101. Second row: Left panel shows the
angular velocity ($\Omega_{\textrm{H\,{\sc i}}}$) profile in black, while the
blue, green, pink and orange profiles denote the various Lindblad curves. The
radius of co-rotation ($R_{c}$) is denoted by the black vertical dashed line.
The blue, green, pink and orange ellipses (and vertical dashed lines) always
represent the ILR, iUHR, oUHR and OLR locations. The bottom two rows are
similar to the first two rows but shown for the galaxy NGC 1371. Figure 16:
Each panel shows the plots of the rotation curve (v${}_{\textrm{rot}}$) and
the radial variation of the inclination ($i$) and position angles (PA) from
the 3DBarolo fits to the galaxies. Figure 17: Each panel shows the plots of
the rotation curve (v${}_{\textrm{rot}}$) and the radial variation of the
inclination ($i$) and position angles (PA) from the 3DBarolo fits to the
galaxies. Figure 18: The HI spectra for the H i-RINGS sample. The grey lines
are spectra from HIPASS with shaded rms of 13.3mJy and the purple lines are
spectra from ATCA observations.  Figure 19: The HI spectra for the H i-RINGS
sample. The grey lines are spectra from HIPASS with shaded rms of 13.3mJy and
the purple lines are spectra from ATCA observations.