Doctor of Philosophy

BIBLIOGRAPHY

152 Chapter 4

WALLABY Early Science - III. An H I Study of the Spiral Galaxy NGC 1566

ABSTRACT

This paper reports on the atomic hydrogen gas (HI) observations of the spiral galaxy NGC 1566 using the newly commissioned Australian Square Kilometre Array Pathﬁnder (ASKAP) radio telescope. This spiral galaxy is part of the Dorado loose galaxy group, which has a halo mass of

13.5 1 10 M . We measure an integrated HI ﬂux density of 180.2 Jy km s− emanating from this ∼ galaxy, which translates to an HI mass of 1.94 1010 M at an assumed distance of 21.3 Mpc. Our × observations show that NGC 1566 has an asymmetric and mildly warped HI disc. The HI-to-stellar mass fraction (MHI/M ) of NGC 1566 is 0.29, which is high in comparison with galaxies that have ∗ the same stellar mass (1010.8 M ). We also derive the rotation curve of this galaxy to a radius of

50 kpc and fit different mass models to it. The NFW, Burkert and pseudo-isothermal dark matter halo profiles fit the observed rotation curve reasonably well and recover dark matter fractions of

0.62, 0.58 and 0.66, respectively. Down to the column density sensitivity of our observations

19 2 (NHI = 3.7 10 cm ), we detect no HI clouds connected to, or in the nearby vicinity of, the × − HI disc of NGC 1566 nor nearby interacting systems. We conclude that, based on a simple analytic model, ram pressure interactions with the IGM can affect the HI disc of NGC 1566 and is possibly the reason for the asymmetries seen in the HI morphology of NGC 1566.

4.1 Introduction

The formation and evolution of a galaxy is strongly connected to its interactions with the surrounding local environment (Gavazzi and Jaffe, 1986; Okamoto and Habe, 1999; Barnes and Hernquist,

153 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

1991; Antonuccio-Delogu et al., 2002; Balogh et al., 2004; Avila-Reese et al., 2005; Blanton et al., 2005; Hahn et al., 2007; Fakhouri and Ma, 2009; Cibinel et al., 2013; Chen et al., 2017; Zheng et al., 2017). The term environment is broad and not only refers to the neighbouring galaxies but also to the tenuous gas and other material between these systems, the so-called intergalactic medium (IGM). These interactions include gas accretion (Larson, 1972; Larson et al., 1980; Tosi, 1988; Sancisi et al., 2008; de Blok et al., 2014b; Vulcani et al., 2018; Rahmani et al., 2018), ram pressure stripping due to the interaction with the IGM (Dickey and Gavazzi, 1991; Balsara et al., 1994; Abadi et al., 1999; Boselli and Gavazzi, 2006; Westmeier et al., 2011; Yoon et al., 2017; Jaffé et al., 2018; Ramos-Martínez et al., 2018), or mergers and tidal harassments (Toomre and Toomre, 1972; Farouki and Shapiro, 1982; Dubinski et al., 1996; Alonso-Herrero et al., 2000; Alonso et al., 2004; Brandl et al., 2009; Blanton and Moustakas, 2009; Moreno et al., 2015; Lagos et al., 2018; Elagali et al., 2018b,a). There is a plethora of observational evidence on the relationship between galaxies and their local surrounding environment. For instance, the morphology-density relationship (Oemler, 1974; Dressler, 1980; Dressler et al., 1997; Postman et al., 2005; Fogarty et al., 2014; Houghton, 2015), according to which early-type galaxies (ellipticals and lenticulars) are preferentially located in dense environments such as massive groups and clusters of galaxies, whereas late-type galaxies (spirals, irregulars and/or dwarf irregulars) are located in less dense environments, e.g., loose groups and voids. Furthermore, galaxies in dense environments are commonly HI deﬁcient in comparison with their ﬁeld counterparts (Davies and Lewis, 1973; Haynes and Giovanelli, 1986; Magri et al., 1988; Cayatte et al., 1990; Quilis et al., 2000; Schröder et al., 2001; Solanes et al., 2001; Omar and Dwarakanath, 2005; Sengupta and Balasubramanyam, 2006; Sengupta et al., 2007; Kilborn et al., 2009; Chung et al., 2009; Cortese et al., 2011; Dénes et al., 2016; Yoon et al., 2017; Brown et al., 2017; Jung et al., 2018). This is because dense environments promote all sorts of galaxy-galaxy and/or galaxy-IGM interactions, which have various consequences on the participant galaxies including their colour and luminosity (Loveday et al., 1992; Norberg et al., 2001; Croton et al., 2005; Kreckel et al., 2012; McNaught-Roberts et al., 2014) as well as their star formation rates (Balogh et al., 1997; Poggianti et al., 1999; Lewis et al., 2002; Gómez et al., 2003; Kauffmann et al., 2004; Porter et al., 2008; Paulino-Afonso et al., 2018; Xie et al., 2018). Hence, a complete picture of the local environments surrounding galaxies is essential to distinguish between the different mechanisms that enable galaxies to grow, or quench (Thomas et al., 2010; Peng et al., 2010; Putman et al., 2012; Marasco et al., 2016; Bahé et al., 2019).

The effects of galaxy interactions with the surrounding environment are most noticeable in the outer discs of the participant galaxies. However, HI discs are better tracers of the interactions with

154 4.1. INTRODUCTION

Figure 4.1: Panel (a) & (b): The SINGG Hα and R-band image of NGC 1566. Panel (c): The GALEX FUV image of NGC 1566. Panel (d): The three-color RGB image of NGC 1566; Red represents the Hα image, Green the R-band image and Blue the FUV image. The Hα, and R-band images are convolved to 600, to match the resolution of the FUV image (Wong, 2007; Meurer et al., 2006). the local environment, because HI discs are commonly more extended and diffuse than the optical discs and, as a consequence, more sensitive to external inﬂuences and susceptible to hydrodynamic processes unlike stars (Yun et al., 1994; Braun and Thilker, 2004; Michel-Dansac et al., 2010).

Probing the faint HI gaseous structures in galaxies is an onerous endeavour as it requires high observational sensitivity to low surface brightness features, which means long integration times.

Such observations are now possible on unprecedentedly large areas of the sky. These wide ﬁeld HI surveys will be conducted using state-of-the-art radio telescopes that have subarcminute angular resolution, for instance the Australian Square Kilometre Array Pathﬁnder (ASKAP; Johnston et al., 2007, 2008), the Karoo Array Telescope (MeerKAT; Jonas and MeerKAT Team, 2016) as well as the APERture Tile In Focus (APERTIF; Verheijen et al., 2008).

155 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

The Wideﬁeld ASKAP L-band Legacy All-sky Blind surveY (WALLABY) is an HI imaging survey that will be carried out using ASKAP radio telescope to image 3/4 of the sky out to a redshift of z 0.26 (Koribalski, 2012). ASKAP is equipped with phased-array feeds (PAFs, Hay ∼ and O’Sullivan, 2008), which can deliver a ﬁeld-of-view of 30 square degrees, formed using 36 beams at 1.4 GHz (Koribalski, 2012). WALLABY will provide HI line cubes with relatively high

1 sensitivity to diffuse emission with a root-mean-square (rms) noise of 1.7 mJy beam− per 4 km

1 s− channel and will map at least 500, 000 HI emitting galaxies over its entire volume (Duffy et al.,

2012). Hence, WALLABY will help revolutionise our understanding of the behaviour of the HI gas in different environments and the distribution of gas-rich galaxies in their local environments. The prime goal of this work is to provide the community with an example of the capabilities of the wideﬁeld HI spectral line imaging of ASKAP and the scope of the science questions that will be addressable with WALLABY. We present the ASKAP HI line observations of the spiral galaxy NGC 1566 (also known as WALLABY J041957-545613) and validate these with archival single-dish and interferometric HI observations from the Parkes 64 m telescope using the 21-cm multibeam receiver (Staveley-Smith et al., 1996) and the Australia Telescope Compact Array

(ATCA), respectively. Further, we use the sensitivity and angular resolution of the ASKAP HI observations to study the asymmetric/lopsided HI gas morphology and warped disc of NGC 1566 and attempt to disentangle the different environmental processes affecting the gas and the kinematics of this spectacular system. NGC 1566 is a face-on SAB(s)bc spiral galaxy (de Vaucouleurs, 1973; de Vaucouleurs et al., 1976), and is part of the Dorado loose galaxy group (Bajaja et al., 1995; Agüero et al., 2004; Kilborn et al., 2005).

Figure 4.1 shows the Survey for Ionisation in Neutral Gas Galaxies (SINGG; Meurer et al., 2006)Hα and R-band images, the GALEX FUV image and a three-color RGB image of NGC 1566, where the Hα, R-band, and FUV images present the Red, Green and the Blue, respectively. This galaxy has a weak central bar (north-south orientation with a length of 32 .5, Hackwell and ∼ 00 Schweizer, 1983) and two prominent star-forming spiral arms that form a pseudo-ring in the outskirts of the optical disc (Comte and Duquennoy, 1982). The Hα emission map of NGC 1566 is dominated by a small but extremely bright Hα complex region located in the northern arm that emits approximately a quarter of the total disc’s Hα ﬂux when excluding the emission from the nucleus (Pence et al., 1990). NGC 1566 hosts a low-luminosity Seyfert nucleus (Reunanen et al., 2002; Levenson et al., 2009; Combes et al., 2014; da Silva et al., 2017) known for its variability from the X-rays to IR bands (Alloin et al., 1986; Glass, 2004). The origin of the variability in AGNs is still controversial and is hypothesised to be caused by processes such as accretion prompted by

156 4.1. INTRODUCTION

Table 4.1: The properties of NGC 1566 adopted from the literature and reported in this paper.

Property NGC 1566 Reference Right ascension (J2000) 04:20:00.42 de Vaucouleurs(1973) Declination (J2000) 54:56:16.1 de Vaucouleurs(1973) Morphology SAB(rs)bc− de Vaucouleurs(1973) 1 HI systemic velocity (km s− ) 1496 This work 11 LR (L ) 1.2 10 Meurer et al.(2006) × 10 LFIR (L ) 2.5 10 Sanders et al.(2003) 2 × µR (ABmag arcsec− ) 19.3 Meurer et al.(2006) mI (ABmag) 8.7 Walsh(1997) m3.6µm (ABmag) 10.15 Laine et al.(2014) mV (ABmag) 9.8 Walsh(1997) M (M ) 6.5 1010 This work ∗ × Hα Equivalent width (Å) 39 3 Meurer et al.(2006) ± MFUV (ABmag) 20.66 Wong(2007) − 10 MHI (M ) 1.94 10 This work × 9 MH2 (M ) 1.3 10 Bajaja et al.(1995) × 6 MBH (M ) 8.3 10 Woo and Urry(2002) × Position angle (degree) 219 4 This work ± Inclination (degree) 31 7 This work 1 ± SFRHα(M yr− ) 21.5 Meurer et al.(2006)

D25 (kpc) 35 Walsh(1997) Distance (Mpc) 21.3 Kilborn et al.(2005) disc instabilities, surface temperature ﬂuctuations, or even variable heating from coronal X-rays (Abramowicz et al., 1986; Rokaki et al., 1993; Zuo et al., 2012; Ruan et al., 2014; Kozłowski et al., 2016). Although, NGC 1566 has been subjected to numerous detailed multiwavelength studies, as cited above, this work presents the ﬁrst detailed HI study of this spiral galaxy. Table 4.1 presents a summary of the relevant properties of NGC 1566 from the literature and from this paper.

This paper is organised as follows: in Section 4.2, we describe the WALLABY early science observations and reduction procedures along with the previously unpublished archival data obtained from the ATCA online archive. Section 4.3 describes our main results, in which we provide detailed analysis of the HI morphology and kinematics of this galaxy. In Section 4.4, we ﬁt the observed rotation curve to three different dark matter halo models, namely, the pseudo-isothermal, the Burkert and the Navarro-Frenk-White (NFW) halo proﬁles. Section 4.5 discusses the possible scenarios leading to the asymmetry in the outer gaseous disc of NGC 1566 and presents evidence that ram pressure could be the main cause of this asymmetry. In Section 4.6, we summarise our main

ﬁndings. For consistency with previous HI and X-ray studies of the NGC 1566 galaxy group, we adopt a distance of 21.3 Mpc (Kilborn et al., 2005; Osmond and Ponman, 2004), which is based on

1 1 a Hubble constant of H0 = 70.3 km s− Mpc− , though we note that this is at the upper end of

157 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

-52°00'

-54°00' Dec (J2000) Dec -56°00'

-58°00'

4h48m00.0s 36m00.0s 24m00.0s 12m00.0s 00m00.0s RA (J2000)

Figure 4.2: The Dorado observations footprint, where blue and red circles represent footprints A and B, respectively. The inner eight black circles represent the beams around NGC 1566 used in this paper. The background image is the blue band Digitised Sky Survey (DSS) image, and the green circle highlights the location of NGC 1566. values quoted in the NASA/IPAC Extragalactic Database (NED). For more details refer to Section 4.4.5.

4.2 Data

4.2.1 WALLABY Early Science Observations

ASKAP is situated in the Murchison Radioastronomy Observatory in Western Australia, a remote radio-quiet region about 305 km North-East of Geraldton in Western Australia. ASKAP is one of the new generation of radio telescopes designed to pave the way for the Square Kilometre Array

(SKA; Dewdney et al., 2009). This SKA precursor consists of 36 separate 12-metre radio dishes

1 that are located at longitude 116.5◦ east and latitude 26.7◦ south . Each 12-metre antenna has a single reﬂector on an azimuth-elevation drive along with a third axis (roll-axis) to provide all-sky coverage and an antenna surface capable of operation up to 10 GHz. The antennas are equipped with

Mark two (MK II) phased array feeds (PAFs), which provide the antennas with a 30 square degree ﬁeld-of-view, making this radio telescope a surveying machine (DeBoer et al., 2009; Schinckel and

Bock, 2016). During the ﬁrst ﬁve years of operation, ASKAP will mainly carry out observations for

1https://www.atnf.csiro.au/projects/askap/index.html

158 4.2. DATA ten science projects, one of which is WALLABY2.

2https://wallaby-survey.org/

159 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566 WALLABY early science observations: Dorado ﬁeld Table 4.2: Observation Dates28 Time Dec on 2016 Source (hrs)29 Dec 2016 Bandwidth (MHz)30 Dec 2016 Central Frequency (MHz)31 11.1 Dec 2016 Number01 of 11.2 Jan Antennas 201702 11.1 Footprint Jan 201703 12.6 Jan 201723 Sep 11.1 201724 Sep 12.0 201725 Sep 3.8 201726 Sep 12.0 201727 12.0 Sep 201728 4.0 Sep 192 201715 12.0 Dec 201703 Jan 12.0 201804 Jan 12.0 2018 9.1 12.0 240 1344.5 9.0 240 1368.5 10 10 10 9 1320.5 10 12 10 B 9 A 12 B A 12 16 12 A 12 B B 12 A A 16 16 A A B A B B A

160 4.2. DATA

In October 2016, ASKAP early science observations program started using 12 PAF-equipped

ASKAP antennas (out of 36 antennas) to pave the way and improve the data reduction and analysis techniques while commissioning ASKAP to full speciﬁcations. Over 700 hours of early science observations were dedicated to WALLABY, during which four ﬁelds were observed to the full survey

1 1 sensitivity depth, rms noise sensitivity per 4 km s− channel of 1.7 mJy beam− (Lee-Waddell et al., 2019; Reynolds et al., 2019). One of these fields is the Dorado early science field. Figure 4.2 shows the two interleaves of the square 6 6 beam footprint of the observations overlaid on the Digitised × Sky Survey (DSS) image of this field. The blue and red beams are referred to as footprint A (centred at α = 04:18:35, δ = 54:51:43 ;J2000) and B (centred at α = 04:21:44.7, δ = 55:18:33.8 ; − − J2000), respectively. Mosaicking the two interleaves reduces the noise level at the edge of the beams and produces a smoother noise pattern. The inner eight black circles represent the beams around NGC 1566 used to make the HI cube that forms the basis of this work. The green circle shows the location of NGC 1566. The observations of this field started in December 2016 and were completed in January 2018 using at the beginning only 9 - 10 ASKAP antennas, however the most recent observations were completed using 16 ASKAP antennas. The baseline range of the ASKAP array for the current observations is between 22-2305 m, which ensured excellent uv-coverage and a compromise between the angular resolution and the surface brightness sensitivity of the observations. The bandwidth of the Dorado observations ranges between 192 and 240 MHz, due to upgrades to the correlator capacity during the early science period, and has a channel width

1 of 18.5 kHz (4 km s− ). For each day of observation, the primary calibrator, PKS1934-638, is observed at the beginning for two to three hours and is positioned at the centre of each of the 36 beams. The total on-source integration time is 167.0 hr; 83.2 hr for footprint A and 83.8 hr for footprint B. Refer to Table 4.2 for a summary of the Dorado early science observations.

We use ASKAPSOFT3 to process the Dorado observations. ASKAPSOFT is a software processing pipeline developed by the ASKAP computing team to do the calibration, spectral line and continuum imaging, as well as the source detection for the full-scale ASKAP observations in a high-performance computing environment. This pipeline is written using C++ and built on the casacore library among other third party libraries. A comprehensive description of ASKAPsoft reduction pipeline is under preparation in Kleiner et al. (in prep.). The reader can also refer to Lee-Waddell et al.(2019) or Reynolds et al.(2019) for a similar brief description of the reduction and pipeline procedures. For each day of observation, we ﬂag and calibrate the measurement data

3https://www.atnf.csiro.au/computing/software/askapsoft/sdp/docs/current/ index.html

161 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

Figure 4.3: Optical DSS image (blue band) on a linear stretch with contours from ASKAP continuum map overlaid at 2, 2, 5, 10, 15, 20, 25 multiplied by 1 σ (0.3 mJy beam 1). The 2 σ level is − − − displayed with dashed contours. The restored beam of ASKAP observation is in the bottom left.

Table 4.3: ASKAP and ATCA HI Observations Results

Parameter ASKAP value ATCA value rms noise (Jy beam 1 per 4 km s 1 channel ) 1.7 10 3 2.4 10 3 − − × − × − Synthesised Beam Size (arcsec arcsec) 42 35 45 40 × × × Synthesised Beam Size (kpc kpc) 4.3 3.6 4.6 4.1 × × × Beam PA (degree) 80.0 6 1 Channel width (km s− ) 4 4 Channel map pixel size (arcsec arcsec) 6 6 6 6 × × × NGC 1566 HI total ﬂux (Jy km s 1) 180.2 16.3 177.0 14.0 − ± ± NGC 1566 peak ﬂux density (Jy) 1.35 0.09 1.24 0.02 ± ± NGC 1566 HI mass (M ) (1.94 0.18) 1010 (1.91 0.15) 1010 ± × ± × w Line Width (km s 1) 208 10 201 8 50 − ± ±

162 4.2. DATA

set on a per-beam basis using ASKAPSOFT tasks CFLAG and CBPCALIBRATOR, respectively. Using the CFLAG utility, we flag the autocorrelations and the spectral channels affected by radio frequency interference (RFI) by applying a simple flat amplitude threshold. Then, we apply a sequence of Stokes-V flagging and dynamic flagging of amplitudes, integrating over individual spectra. We process the central 8 beams (4 in each footprint) of each observation using an 8 MHz bandwidth

1 (432 channels), between 1410 to 1418 MHz (velocity range between 511 to 2196 km s− ) to save computing time and disc space on the Pawsey supercomputer. We then process the calibrated visibilities to make the continuum images using the task IMAGER and self-calibrate (three loops) to remove any artefacts or sidelobes from the continuum images.

Prior to the spectral-line imaging stage, we subtract radio continuum emission from the visibility data set using the best-ﬁt continuum sky model produced in the previous step. Thereafter, we combine the data set for each beam, seven nights in footprint B and nine nights in footprint A, in the uv domain and image using the ASKAPSOFT task IMAGER with a robust weighting value of

+0.5 and a 30 arcsec Gaussian taper. We clean the combined image for each beam using a major and minor cycle threshold of 3σ, three times the theoretical rms noise for the seven observations combined. To measure the theoretical rms noise of each epoch, we use the on-line sensitivity calculator 4, with antenna efﬁciency value of 0.7 and system temperature of 50 K. For instance, the

1 theoretical rms noise for ten hours using 10, 12 or 16 antennas is 5.03, 4.15 and 3.08 mJy beam− , respectively. Then, we subtract the residual continuum emission from the restored cube using the

ASKAPSOFT task IMCONTSUB and mosaicked the 8 beams using the ASKAPSOFT task LINMOS.

Even though mosaicking different ASKAP beams with LINMOS can introduce correlated noise to the

final HI cube, this has no effect on the final flux scale, and only a minor effect on the rms noise. This image domain continuum subtraction is necessary to obtain a higher dynamic range HI cube and remove any remaining residual artefacts, which aids source finding and parameterisation. Figure 4.3 shows the optical DSS image (blue band) overlaid with contours from the ASKAP continuum map.

The total 21 cm continuum ﬂux density of NGC 1566 from ASKAP observations is 199 3 mJy, ± in agreement with the value of 204 28 mJy (at 21.7 cm ) reported by Ehle et al.(1996) using ± the ATCA. The restored synthesised beam has a size of θFWHM = 42 35 with a position 00 × 00 1 1 angle of 80◦. The HI line cube has rms noise per 4 km s− channel of 1.7 mJy beam− , which

19 2 translates to a 3σ column density sensitivity of NHI = 1.5 10 cm . The rms noise in a × − 1 1 wider channel width (20 km s− ) equals 0.84 mJy beam− , and corresponds to a 3σ column density

19 2 sensitivity of NHI = 3.7 10 cm . The properties of the ﬁnal cube are summarised in Table 4.3. × − 4http://www.atnf.csiro.au/people/Keith.Bannister/senscalc/

163 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

164 4.2. DATA Observations: Instrumental Parameters I H 375B 1.5D 750B 1.5C ATCA Archival Table 4.4: ParameterObservation DatesOn-source integration time (hrs)Shortest Baseline (m) 9Longest Baseline (m)Central Frequency (MHz)Bandwidth (MHz) 1994 April 05 31 1994 May 26 1413 9 5969 1994 June 01 1994 June 17 Array Conﬁguration 8 1413 107 4439 9 1413 4500 61 8 9 1413 4500 77 8 8

165 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

4.2.2 ATCA Observations

NGC 1566 was observed using the ATCA in four epochs between April 1994 and June 1994 with each epoch being 12 h in duration. Four different ATCA array configurations were used to observe NGC 1566, namely the 1.5C, 375, 750B as well as the 1.5D configuration (Walsh, 1997). These configurations have baseline distances in the range between 31 and 5969 m. In each of the four epochs, the observation was centred at 1413 MHz (the redshifted 1420 MHz line frequency for NGC 1566) for a duration of nine hours on the source (a total of 36 hrs, see Table 4.4). The bandwidth of these observations is 8 MHz, over 512 velocity channels. Hence, each channel

1 corresponds to 15.625 kHz in width, and a 3.3 km s− velocity resolution. The ATCA primary calibrator, PKS1934 638 (ﬂux density S1396 MHz = 14.9 Jy), was observed before each epoch of − observations for a duration of 30 minutes and used as the bandpass calibrator. The phase calibrator

PKS0438 436 (S1396 MHz = 6.09 Jy) was observed each hour during the four epochs for a duration − of 15 minutes to ensure the precision of the calibration.

We follow the standard procedures described in Elagali et al.(2018b) to ﬂag, calibrate and image the HI line observations using the MIRIAD package (Sault et al., 1995). We use the MIRIAD

UVLIN task (Sault, 1994; Cornwell et al., 1992) to subtract the radio continuum from the visibility data set. Then, we use MIRIAD INVERT task to Fourier-transform the continuum subtracted visibilities to a map with robust weighting parameter of +0.5 and a symmetric taper of 15 arcsec to have an optimal sidelobe suppression and intermediate weighting between uniform and natural. We

1 also re-sample to the ASKAP resolution of 4 km s− at this stage, and apply the CLEAN task down to three times the theoretical rms noise. As a ﬁnal step, we apply the primary beam correction using the LINMOS task. The synthesised beam-size is θ FWHM = 45 40 with PA = 6 . The 00 × 00 ◦ 1 1 cube has rms noise per 4 km s− channel of 2.4 mJy beam− , close to the theoretical rms noise of

1 1 our observations (2.2 mJy beam− ). The 3σ column density sensitivity per 4 km s− channel is

19 2 1 1 NHI = 1.7 10 cm . Over a 20 km s channel width the rms noise equals 1.2 mJy beam × − − − 19 2 and the corresponding 3σ column density sensitivity is NHI = 4.3 10 cm . × −

166 4.3. GAS MORPHOLOGY AND KINEMATICS

Figure 4.4: The HI line proﬁle of NGC 1566 from our ASKAP observations (black line), from the ATCA (blue dashed line) and Parkes single-dish observations (red line). The vertical line delimits 1 the HI systemic velocity of NGC 1566 Vsys = 1496 km s− derived from our kinematics analysis (refer to Section 4.3).

4.3 Gas Morphology and Kinematics

4.3.1 HI Morphology and Distribution in NGC 1566

Figure 4.4 presents the integrated HI spectrum of NGC 1566 as obtained from ASKAP observations (black line), unpublished archival ATCA observations (blue dashed-line) and re-measured Parkes single-dish spectrum (red line) from Kilborn et al.(2005). We measure a total flux value of 180.2 16.3 Jy km s 1 from ASKAP observations. This flux value corresponds to a total HI ± − mass of 1.94 1010 M , assuming that NGC 1566 is at a distance of 21.3 Mpc. The integrated × HI flux density of NGC 1566 from ASKAP early science observations is within the expected error of the value measured from the ATCA observations (177.0 14.0 Jy km s 1) and from Parkes ± − observations (175.0 7.4 Jy km s 1). We note that the integrated flux density of NGC 1566 derived ± − from our ATCA data reduction is within error of the value reported in the thesis of Walsh(1997).

Figure 4.5 presents the individual HI channel maps of NGC 1566 from ASKAP observations in the

1 1 velocity range between 1628-1393 km s− and with a step-size of 16 km s− . Figure 4.6 shows the HI column density map of NGC 1566 as obtained from ASKAP observations (top panel), the archival ATCA observations (middle panel), and a difference between the two maps (bottom panel).

To produce the difference map, we have convolved the ASKAP and ATCA maps to the same beam

167 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

Figure 4.5: The individual HI channel maps of NGC 1566 from ASKAP observations in the velocity 1 1 range between 1628-1393 km s− and with a step-size of 16 km s− . The velocity of each channel is shown in the top-left, whereas the restored beam is shown in the bottom-right (blue ellipse). Contours are at 3, 3, 5, 10, 20, 40 times the 1 σ noise (0.85 mJy beam 1 per 16 km s 1 channel). − − − The yellow cross marks the kinematic centre of NGC 1566 derived using ROTCUR (refer to Section 4.3).

168 4.3. GAS MORPHOLOGY AND KINEMATICS

Figure 4.6: The HI column density map of NGC 1566 as obtained from ASKAP observations (top panel), the archival ATCA observations (middle panel), and a difference between the two maps (bottom panel). To produce the difference map, we have convolved the ASKAP and ATCA maps to the same beam angular resolution. The ellipse shows the restored beam of ASKAP observations (θFWHM = 42 35 ) and the ATCA observations (θFWHM = 45 40 ). 00 × 00 00 × 00

169 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

Figure 4.7: DSS blue band image of NGC 1566 with column density contours from ASKAP observations overlaid at (0.6, 1.2, 2.4, 4.8, 9.6, 14.0, 16.2) 1020 cm 2. The ellipse shows the × − restored beam of ASKAP observations. Left-to-right in the ﬁgure corresponds to east-to-west of the galaxy. The blue-dashed line corresponds to the rectangular cut through the galaxy, with a width of 5 kpc, that is used to measure the cut in the column density distribution across NGC 1566 in Figure 4.8.

170 4.3. GAS MORPHOLOGY AND KINEMATICS

Figure 4.8: A column density cut across NGC 1566 with respect to right ascension offset from the centre of the galaxy. The width of this cut is 5 kpc and is shown by the blue-dashed line in Figure 4.7. Left-to-right in the ﬁgure corresponds to east-to-west.

angular resolution.5 ASKAP observations are as sensitive as the ATCA observations to low surface brightness features surrounding the outer disc of NGC 1566. This figure emphasises the ability of the ASKAP instrument to probe relatively low column density levels while mapping large areas of the sky. The small flux difference between ASKAP and ATCA in the centre is likely due to the smaller baseline of the ASKAP configuration in comparison with that of the ATCA, however this difference is within the error of the two instruments. Hence, ASKAP will provide the HI community with unprecedented amount of high quality HI line cubes. We note that Reeves et al.(2016) presented the moment zero and velocity maps of NGC 1566 using more recent ATCA observations but with less integration time ( 17 hrs) in comparison with the archival ATCA data used in this paper ∼ ( 36 hrs). The study of Reeves et al.(2016) was mainly focused on the intervening H I absorption ∼ in NGC 1566 along with other nine nearby galaxies.

Figure 4.7 shows the optical DSS blue band image of NGC 1566 with column density contours from ASKAP observations overlaid at (0.6, 1.2, 2.4, 4.8, 9.6, 14.0, 16.2) 1020 cm 2. This map × − shows an HI disc that extends beyond the observed optical disc especially around the northern

5We make use of the Source Finding Application (SOFIA, Serra et al., 2015) to produce these maps. We apply a 4σ detection limit and a reliability of 95 per cent.

171 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

Figure 4.9: The ASKAP line-of-sight velocity (left) and velocity dispersion (right) of NGC 1566. The ellipse marks the optical disc of NGC 1566 (Walsh, 1997). Contours are centred at the systemic 1 velocity of NGC 1566 (1496 km s− ; black-dashed contour) and are equally separated by 20 km 1 s− interval; white and black contours show the approaching and receding sides, respectively. The black ﬁlled ellipse shows the restored beam of the observations.

and western parts of NGC 1566 (also refer to Figure 4.15). The HI gas is very concentrated in the inner arms and gradually decreases following the outer arms of the disc. The HI also highlights the difference between the two outer arms better than in the optical (Figure 4.1). The eastern outer arm forms a regular arc shape that extends between PA = 50◦ until where the PA = 260◦; here the PA is estimated from the north extending eastwards to the receding side of the major axis. However, the western arm is signiﬁcantly shorter extending from PA = 210◦ to PA = 330◦, and appears less regular or disturbed between PA = 270 to 330 . The 1.5 1019 cm 2 column density contour ◦ ◦ × − extends for a diameter of almost 160, which at the adopted distance of NGC 1566 translates to a diameter of 99 kpc. Figure 4.8 shows a column density cut across NGC 1566 with respect to ∼ right ascension offset from the centre and measured at the declination of the centre of the galaxy.

The width of this cut is 5 kpc and is shown by the blue-dashed line in Figure 4.7. We use the KARMA visualisation tool KVIS to generate this column density cut across NGC 1566 (Gooch, 1996). The column density of the HI gas slowly drops with radius, mainly due to presence of the outer arms in

NGC 1566 at radius r 20kpc; the column density falls off at r 20 kpc by only a factor of two ∼ ∼ from the peak value at r 6 kpc. Further, Figure 4.8 shows an asymmetry in the distribution of ∼ the HI gas in this galaxy; the eastern part of the HI disc sharply declines after 30 kpc from the ∼ centre as opposed to the western part which extends beyond 30 kpc and smoothly declines up ∼ to a radius of 50 kpc. This asymmetry is more evident in Figure 4.7, in which the HI contours ∼ show crowding on the east side of the galaxy and are spread out on the west side. The column density asymmetries present in NGC 1566 could be signs of ram pressure interaction, gas accretion and/or past ﬂyby interaction(s) with the other members of NGC 1566 galaxy group. We discuss the

172 4.3. GAS MORPHOLOGY AND KINEMATICS possibility of these interaction scenarios in Section 4.5.

4.3.2 HI Kinematics

Previous studies of the kinematics of NGC 1566 utilised optical spectroscopy to map the numerous emission line regions in this galaxy. For example, Pence et al.(1990) used the Fabry-Pérot interferometer at the 3.9 m Anglo-Australian Telescope to map the Hα emission in NGC 1566 and were able to measure gas velocities out to a radius of 10 kpc (c.f. Figures 4 & 8 in Pence et al., 1990). Figure 4.9 shows the ASKAP line-of-sight velocity and the velocity dispersion of NGC 1566. The shape of the isovelocity contours of the approaching side (white contours) of the galaxy is slightly different in comparison with the receding side (black contours) which suggests the presence of kinematic asymmetry in NGC 1566. The western side of the velocity field shows significant asymmetry in comparison with the eastern side. This is also seen in both the moment zero map and the velocity channel map in Figures 4.5 and 4.6. In the individual channel maps, the western side is more extended than the eastern side, especially at velocities between 1534 1456 km s 1. On the other hand, the dispersion map shows a very high peak in − − the inner regions which is likely a result of beam smearing. There is a noticeable increase in the velocity dispersion associated with the inner arms of NGC 1566, similar to other grand design spirals like M83 (Heald et al., 2016). However the dispersion velocity decreases quickly in the outer arms and the outskirts of the HI disc. We discuss in detail the possibility of an interaction scenario and other external influences that may lead to the lopsidedness in NGC 1566 in Section 4.5.

We follow the standard procedure described in Elagali et al.(2018b) to derive the H I rotation curve of NGC 1566 using a tilted ring model. This model assumes that the gas moves in circular orbits. Each tilted ring is ﬁtted independently as a function of radius and has 6 deﬁning kinematic parameters: the central coordinate (xc,yc), the systemic velocity Vsys, the circular velocity Vrot, the inclination angle i, as well as the position angle PA. According to this model, the observed line-of-sight velocity V(x,y) is given by:

V (x, y) = Vsys + Vrot sin i cos θ, (4.1) where the angle θ is a function of position angle and inclination. We use the Groningen Image

Processing System (GIPSY; Allen et al., 1985; van der Hulst et al., 1992) ROTCUR task (Begeman, 1989) to apply the tilted ring model to the observed velocity ﬁelds of NGC 1566. We run the task

173 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

Figure 4.10: The tilted ring model solutions for the inclination (a), the position angle (b), and the rotation velocity Vrot (c) of NGC 1566 using ROTCUR. Solid line, dashed line and ﬁlled blue circles denote the approaching, receding and both sides, respectively. The green horizontal lines mark the best ﬁt values for the i ( 31◦) and PA ( 219◦), respectively, which are used to estimate the rotation curve for both sides in the panel (c). The vertical dash-dotted line shows the location of the optical radius of NGC 1566 (Walsh, 1997).

174 4.3. GAS MORPHOLOGY AND KINEMATICS

Figure 4.11: The position-velocity diagram of NGC 1566 along the major axis, overlayed is the rotation curve derived using ROTCUR. The red contours are at 2, 4, 8, 16, 20 times the 1 σ noise level (1.7 mJy beam 1 per 4 km s 1 channel), while the blue contour is at 2 σ. − − − in an iterative fashion to determine the above mentioned free kinematic parameters for each ring, and use a ring width that equals half the restored beam-size ( 20 arcsec), i.e., we fit two rings per ∼ beam-size. Since the minor axis provides no information on the rotation curve, we apply cos θ | | weighting function to minimise the contribution of points far from the major axis. Firstly, we fit the systemic velocity Vsys and the dynamical centre out to the edge of the optical disc simultaneously by fixing the PA (221◦) and i (27◦) to their optical values (Pence et al., 1990). We next fit the inclination and position angle simultaneously, keeping the systemic velocity and the dynamical centre fixed to the determined values from the previous step. Then, we smooth the PA and i profiles with a radial boxcar function, estimate their average values and derive an optimum solution for

1 the Vrot. The best ﬁt for the systemic velocity of NGC 1566 is Vsys = 1496 7 km s , and the ± − derived dynamical HI centre agrees with the optical centre. The position angle and inclination values estimated using ROTCUR are PA = 219 4 and i = 31 7 , respectively, and are within ◦ ± ◦ ◦ ± ◦ the error of the optical values determined by Pence et al.(1990). We derive the rotation curve for both the receding and approaching sides, to check for possible departures from symmetry and highlight any systematic uncertainties associated with our ﬁnal results.

Figure 4.10 shows ROTCUR results for the inclination, the position angle and the rotation velocity of NGC 1566. The variation of the inclination with radius implies the existence of a

175 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566 mild warp in the HI disc of this galaxy. The inclination of the inner regions of the disc can be averaged to the value 19◦ (r =< 15 kpc), while the outer parts of the disc (r > 15 kpc) have an average inclination value of 37◦. On the other hand, the position angle derived using ROTCUR is constrained across the HI disc. The green horizontal lines mark the best ﬁt values for the i ( 31◦) and PA ( 219◦), respectively, which are used to estimate the rotation curve for both sides in the lower panel. The rotation of the approaching side is slightly different in comparison with the receding side’s rotation. Figure 4.11 presents the position-velocity diagram of NGC 1566 along the major axis, with its rotation curve overlaid. We note that the low-level emission apparent in the forbidden quadrants (gas with forbidden velocities) is not real and is due to sidelobes. The

1 rotation of both sides reaches a velocity Vrot = 161 13 km s at a radius r = 50 kpc. This ± − velocity translates to, assuming a spherically-symmetric mass distribution, a total dynamical mass

11 of Mtot = (11.30 1.91) 10 M enclosed within this radius. As the ﬁtted inclination of 31◦ ± × is outside the typical range where ROTCUR is thought to be reliable (Begeman, 1989), we also use the the Fully Automated 3D Tilted Ring Fitting Code (FAT; Kamphuis et al., 2015) to derive the rotation curve of NGC 1566. This software works directly on the data cube, thus ﬁtting in 3D, and is more robust against certain instrumental effects and hence is thought to be reliable to lower inclinations (Kamphuis et al., 2015). The results from FAT are consistent with those derived using

ROTCUR. Hence, for brevity, we decide to only show and use the ROTCUR results.

4.4 Dark Matter Content and Mass Models

The gravitational potential of any galaxy is a function of its combined gaseous, stellar and dark matter mass components. In this section, we investigate the distribution of the dark and baryonic matter in NGC 1566, and ﬁt different mass models to the observed rotation curve of this system using the GIPSY task ROTMAS. ROTMAS ﬁts the following equation to the observed rotation (Vobs):

2 2 2 2 VDM(r) = Vobs(r) V (r) Vgas(r), (4.2) − ∗ − where V*, Vgas, and VDM are the contributions of the stars, gas and the dark matter components to the total rotation curve of NGC 1566, respectively. The velocity contributions of the gaseous and the stellar mass components are estimated from their mass radial surface density distribution (Σ(r)) using the GIPSY task ROTMOD (Casertano, 1983). Below, we present the gaseous and the stellar mass radial surface density distributions along with the different dark matter models used to ﬁt the total rotation curve in equation 4.2.

176 4.4. DARK MATTER CONTENT AND MASS MODELS

Figure 4.12: The gaseous and stellar radial mass surface density distributions in NGC 1566 shown in blue and red, respectively.

4.4.1 Gaseous Distribution

To estimate the gaseous mass surface density proﬁle, we use the HI column density map obtained from our ASKAP observations (Figure 4.7). The gas surface density is measured in tilted rings using similar parameters (i, PA, dynamical centre) to those used to derive the total rotation curve with

ROTCUR. We use the task ELLINT in GIPSY to derive the radial HI column density proﬁle NHI(r).

We then convert the HI column density profile to gaseous mass surface density. We scale the gas mass surface density by a factor of 1.4 to account for the presence of helium and metals and assume that it is optically thin. It is important to note that we neglect the contribution of the molecular hydrogen (H2) to the total gaseous distribution in NGC 1566, which can be consider as a limitation to our halo mass model. Figure 4.12 shows the gas mass surface density profile of NGC 1566. We use the radial gas mass surface density profile of NGC 1566 to estimate the corresponding gas rotation velocities and assume that the gas is mainly distributed in a thin disc.

4.4.2 Stellar Distribution

To derive the stellar mass surface density proﬁle in NGC 1566, we use the infrared (IR) photometry obtained from the Infrared Array Camera (IRAC) on board the Spitzer Space Telescope (Werner

177 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566 et al., 2004; Fazio et al., 2004). We convert the IR radial ﬂux density distribution (Sλ(r)) to stellar radial mass density distribution (Σ (r)) using the following equation: ∗

Σ (r) Sλ(r)Yλ, (4.3) ∗ ∼ where λ indicates the wavelength band and Yλ is the mass-to-light ratio with respect to that band. Here, we derive the surface density profile in two different bands, namely, the IRAC 3.6 and 4.5 µm bands. We follow the approach in Oh et al.(2008) to derive the mass-to-light ratio of the IRAC two near-infrared bands; for NGC 1566 the values are Y = 0.648 0.073 and 0.611 0.072 for the ± ± 3.6 and 4.5 µm bands, respectively. Similar to the gaseous profile, we use ELLINT to measure the near-infrared flux density in tilted rings with the same parameters (i, PA, dynamical centre) used to derive the rotation curve of NGC 1566.

1 We then convert the near-infrared flux density from the IRAC pipeline flux units (MJy sr− ) to solar units and apply aperture correction for the 3.6 and 4.5 µm flux density. The stellar mass radial surface density Σ (r), is calculated by the following equation: ∗

Sλ(r) Σ (r) = CλYλfλ , (4.4) ∗ ZP where fλ is the aperture correction factor, Cλ is the conversion factor and ZP is the zero point magnitude for each band. We adopt aperture correction values of f3.6µm = 0.944 and f4.5µm = 0.937 and zero point magnitude ﬂuxes of 280.9 and 179.7 Jy for the 3.6 and 4.5 µm (Reach et al., 2005), respectively. Following the calculations in Oh et al.(2008) based on the spectral-energy

2 distributions of the Sun, we use C3.6µm = 0.196 and C4.5µm = 0.201 M pc− . Figure 4.12 shows the stellar mass radial surface density profile of NGC 1566. We measure a total stellar mass of (6.5 0.4) 1010 M , using the averaged 3.6 and 4.5 µm bands stellar mass radial surface ± × density profiles, which is within the error of the value reported in Laine et al.(2014). The stellar mass radial profile can be described by a simple exponential function (Σ exp( r/h)), for ∗ ∼ − which the the radial scale-length (h) equals 3 kpc and thus the scale-height of the disc (z0), using h/z 5 (Kregel et al., 2002; van der Kruit and Searle, 1981), equals 0.6 kpc. Using ROTMOD, 0 ' we construct the stellar velocity component from the stellar mass radial surface density profile of NGC 1566, assuming that the stellar disc has a vertical sech2 scale-height distribution with z0 = 0.6 kpc (van der Kruit and Searle, 1981).

178 4.4. DARK MATTER CONTENT AND MASS MODELS

4.4.3 Dark Matter Halo Proﬁles

Pseudo-isothermal Dark Matter Proﬁle

This proﬁle is the simplest and most commonly used in the studies of galaxies’ rotation curves

(Kent, 1987; Begeman et al., 1991). This model assumes a central constant-density core (ρ0 ) and a density proﬁle given by: ρ ρ r 0 , ( ) = 2 (4.5) 1 + (r/rc) where rc is the core radius. The corresponding rotation velocity (v(r)) to the pseudo-isothermal (ISO) potential is:

2 2h rc r i v (r) = 4πGρ0rc 1 arctan , (4.6) − r rc

(Kent, 1986).

Burkert Dark Mtter Proﬁle

This profile adopts the following definition for the dark matter density profile Burkert(1995):

ρ r3 ρ r 0 c , ( ) = 2 2 (4.7) (r + rc)(r + rc )

Similar to the ISO proﬁle, rc and ρ0 denote the core radius and the central density of the dark matter halo, respectively. This proﬁle resembles the distribution expected for a pseudo-isothermal sphere at the inner radii (r << rc) and predicts central density ρ0. At larger radii, the Burkert density

3 proﬁle ρ(r) is roughly proportional to r− . The velocity corresponding to this proﬁle is given by (Salucci and Burkert, 2000):

6.4Gρ r3 h r 1 r2 r i v2 r 0 c . ( ) = ln 1 + + ln 1 + 2 arctan (4.8) r rc 2 rc − rc

To ﬁt the rotation curve of this halo proﬁle, we have two free parameters, namely, the core radius and the central density of the dark matter halo.

NFW Dark Matter Proﬁle

Navarro et al.(1996, 1997) used N-body simulations to explore the equilibrium density profile of the dark matter halos in a hierarchically clustering universe and found that these profiles have the same shape, regardless of the values of the cosmological parameters or the initial density fluctuation spectrum. The density profile in this case can be described by the following equation:

179 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

250 Pseudo-iosthermal ] 1

− 200

150

100

Velocity [km s 50

0 250 Burkert ] 1

− 200

150

100

Velocity [km s 50

0 250 NFW ] 1

− 200

150

100

Velocity [km s 50

0 0 5 10 15 20 25 30 35 40 45 50 Radius [kpc]

Figure 4.13: The mass models fit to the observed rotation curve of NGC 1566 using the ISO, Burkert and NFW dark matter profiles. The black circles and grey line show the observed rotation curve and the best fit modelled rotation curve of NGC 1566, respectively. The red, blue and dashed-black lines show the contributions of the gaseous, stellar and dark matter components to the overall rotation curve of NGC 1566, respectively.

180 4.4. DARK MATTER CONTENT AND MASS MODELS

δcρcrit ρ(r) = , (4.9) 2 r 1 + r rs rs where rs is the scale radius, δc is a critical dimensionless density of the halo and ρcrit is the critical density for closure. The NFW halo proﬁle is similar to the Burkert proﬁle; the only difference is at

1 r << rs, in which the NFW halo density ρ(r) r instead of a constant core density value as it ∼ − is the case for the Burkert proﬁle. The velocity of this proﬁle is given by:

v2 ln(1 + cx) cx v2(r) = 200 − 1+cx , (4.10) x ln(1 + c) c − 1+c where x is the radius (x = r/r200) in virial radius units, c is the halo concentration c = r200/rs and v200 is the circular velocity at r200:

r GM200 v200 = = 10H0r200. (4.11) r200

Here, H0 is the Hubble constant and M200 is the virial mass. To ﬁt the rotation curve of the NFW halo proﬁle, we have two free parameters, namely, the scale radius and the virial radius of the dark matter halo.

4.4.4 Mass Model Results

To derive the dark matter contribution to the total rotation curve of NGC 1566, we fit equation 4.2 using ROTMAS. Figure 4.13 shows the mass model results using the ISO, Burkert and the NFW halo profiles. Table 4.5 lists the results of our mass models. In all models, the stellar disc dominates the rotation curve up to a radius r 18 kpc and starts to sharply decline at r 20 kpc. On the other ∼ & hand the dark matter rises linearly with radius reaching the maximum at r 10 kpc and remains ∼ fairly constant at larger radii. Based on the ISO, Burkert and the NFW halo profiles, we estimate dark matter fractions in NGC 1566 to be 0.66, 0.58 and 0.62, respectively. The three dark ' ' ' matter profiles result in reasonable fits. However, due to the lack of angular resolution in the inner regions of NGC 1566 ( 2 kpc), we can not differentiate between the three dark matter density ∼ profiles. To distinguish between different dark matter halo profiles, higher angular resolution observations of few hundred parsec scales are required (see for example Oh et al., 2015; de Blok, 2010; Bolatto et al., 2002; de Blok et al., 2001). This question will soon be addressable for large samples of nearby galaxies using the 36 ASKAP antennas (longest baseline of 6 km) and MeerKAT telescope (de Blok et al., 2016), which will aid constraining the mass distributions of both the dark and baryonic matter in large numbers of galaxies.

181 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

Table 4.5: Mass models ﬁt results for NGC 1566.

Parameter Pseudo-isothermal Burkert NFW rc (kpc) 1.9 0.6 4.3 0.8 - 3 ± ± ρ0 (M pc− ) 0.092 0.050 0.082 0.037 - ± ± rs (kpc) - - 9.52 1.78 ± r200 (kpc) - - 80.84 3.71 11 ± MDM (10 M ) 7.60 5.38 6.29 2 χred 0.34 1.36 0.65 fDM 0.66 0.58 0.62

4.4.5 Tully-Fisher Distance of NGC 1566

Here, we use the Tully-Fisher (TF) relation (Tully and Fisher, 1977) in an attempt to measure a more accurate distance for NGC 1566 than currently reported in the literature. We use the I-band apparent magnitude (mI ) and line width (w50) values reported in Table 4.1& 4.3, respectively, and the kinematic inclination derived in Section 4.3.2. The distance of NGC 1566 using the I-band TF

+7.5 relation (Masters et al., 2006) is 16.9 4.1 Mpc, which is smaller than the value we adopt but within − the errors. Unfortunately, the error in the TF distance does not allow us to deﬁnitively exclude the far distance value reported in NED nor claim a superior distance measurement for this galaxy. We have therefore left the nominal distance of this galaxy as 21.3 Mpc throughout this paper. The large error in the TF relation is due to the low inclination of NGC 1566, which is also reﬂected by the large uncertainties in the distance measurements for this galaxy in the literature.

4.5 Discussion

4.5.1 Possible Origins for the HI Disc Asymmetries

Many disc galaxies are asymmetric, and have lopsided stellar and/or gaseous components (Baldwin et al., 1980; Bournaud et al., 2005; Mapelli et al., 2008). Both theoretical and observational studies suggest three different environmental mechanisms that can cause such asymmetries, namely, ram pressure interactions with the IGM, galactic interactions as well as gas accretion from hosting/neighbouring ﬁlaments (Oosterloo et al., 2007; McConnachie et al., 2007; Reichard et al., 2008; Mapelli et al., 2008; Yozin and Bekki, 2014; de Blok et al., 2014b; Vulcani et al., 2018). In this subsection, we explore the possibility of each of these scenarios given the available data for NGC 1566. Below, we combine the discussion of the galactic interactions and gas accretion scenarios in one

182 4.5. DISCUSSION subsection for brevity while discussing the ram pressure stripping scenario separately.

Interactions and/or Accretion Scenarios:

Tidal interactions between NGC 1566 and neighbouring galaxies could lead to its asymmetries and warps. The interacting galaxy pair NGC 1596/1602 (Chung et al., 2006) is a possible candidate for such a tidal encounter, the physical projected separation between this galaxy pair and NGC 1566 is

700 kpc (2 on the sky). Figure 4.14 shows the column density map of NGC 1566 mosaic field. ∼ ◦ We detect six galaxies including NGC 1566 in this mosaic. A detailed description of these galaxies and the remainder of the Dorado group galaxies will be presented in Elagali et al. (in prep.) and Rhee et al. (in prep.). Even though tidal interactions between NGC 1566 and the neighbouring galaxies may have occurred, we do not see any HI tail/bridge that would result from such an interaction. The centre of the NGC 1566 group is shown by the blue filled circle in Figure 4.14, as defined by where the two most massive and bright galaxies of this group are located, namely NGC

1553/1549 (Kormendy, 1984; Kilborn et al., 2005; Kourkchi and Tully, 2017). We detect no HI emission above the rms noise of ASKAP observations from this interacting galaxy pair. The upper

HI mass limit of NGC 1553/1549 based on the 3σ noise level of this observations and over 40 km

1 6 1 s− channel widths (ten channels) is 3.7 10 M beam− at the distance of this galaxy pair ' × (17.6 Mpc, Kilborn et al., 2005).

Figure 4.15 presents a deep optical image of NGC 1566 taken by David Malin at the Australian Astronomical Observatory. We also examine this deep image and see no faint stellar substructures around NGC 1566, nor a dramatic system of streams/plumes that may have formed through a tidal interaction or minor merger with neighbouring galaxies; see for example Kado-Fong et al.(2018); Pop et al.(2018); Elagali et al.(2018b); Martínez-Delgado et al.(2015); de Blok et al.(2014a);

Martínez-Delgado et al.(2010). We note that these faint structures can form on timescales of 108 years and exist for few gigayears after the interaction (Hernquist and Quinn, 1988, 1989). Hence, a recent interaction scenario is less likely to be the reason for the asymmetries seen in the outer

HI disc of NGC 1566. Alternatively, the lopsidedness of the HI distribution of NGC 1566 can be a result of gas accretion. de Blok et al.(2014b) present an example of this scenario, in which a low column density, extended cloud is connected to the observed main HI disc of NGC 2403. In the case of NGC 1566, we do not detect any clouds or ﬁlaments connected to, or in the nearby vicinity of, the HI disc of this galaxy. We note that our 3σ column density sensitivity measured

1 19 2 over a 20 km s width is NHI = 3.7 10 cm . Therefore, it is difﬁcult to say any more − × − quantitative/conclusive statements about the accretion effects on the HI disc of NGC 1566 and that

183 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

Figure 4.14: The HI column density map of NGC 1566 mosaic ﬁeld as revealed by ASKAP early science observations. The synthesised beam size is θFWHM = 4200 3500. There are six galaxies × 1 detected in this mosaic, within a velocity range between 950 2040 km s− . The projected physical distance between NGC 1566 and the galaxy pair NGC 1596/1602− is shown by the blue horizontal line. The blue ﬁlled circle shows the centre of the group, and the location of the two interacting galaxies NGC 1553/1549 (Kilborn et al., 2005; Kourkchi and Tully, 2017).

more sensitive HI observations of this galaxy are needed to rule out an accretion scenario.

Ram Pressure Stripping Scenario:

Ram pressure is widely observed in massive galaxy clusters (White et al., 1991; Abadi et al., 1999; Acreman et al., 2003; Randall et al., 2008; Vollmer, 2009; Merluzzi et al., 2016; Ruggiero and Lima Neto, 2017; Sheen et al., 2017), such as the nearby Virgo cluster (Chung et al., 2007; Yoon et al., 2017), and is the main reason for the stripping and removal of gas in galaxy clusters especially closer to their dynamical centres. Even though ram pressure stripping is not as prevalent in galaxy groups, a few cases have been reported in the literature (Kantharia et al., 2005; McConnachie et al., 2007; Westmeier et al., 2011; Rasmussen et al., 2012; Heald et al., 2016; Vulcani et al.,

2018). Many authors also ascribe ram pressure as a potential cause for HI deﬁciency in groups (see for example Sengupta and Balasubramanyam, 2006; Sengupta et al., 2007; Freeland et al., 2010;

184 4.5. DISCUSSION

Figure 4.15: Deep optical image of NGC 1566 observed by David Malin and available at the Australian Astronomical Observatory database. The red and orange contours are at HI column 20 2 20 2 density values of NHI = 3.7 10 cm and 0.6 10 cm , respectively. The black ellipse × − × − shows the restored beam of the ASKAP observations.

Dénes et al., 2016). Here, we investigate the asymmetries present in the outskirts of NGC 1566 and its connection to the gaseous halo undergoing ram pressure stripping as a consequence of its interaction with the IGM. As a galaxy passes through the IGM with an inclined orientation, the HI gas will be compressed in the leading edge of the outer disc while the gas at the lagging edge get stripped and pulled away as a result of the ram-pressure forces (refer to Figure 1 and 3 in Quilis et al., 2000). This will produce an asymmetry in the HI column density distribution similar to that observed in NGC 1566, in which the south-eastern edge of the HI disc sharply declines after

30 kpc from the centre, while the north-western edge is more extended and smoothly declines ∼ with radius up to 50 kpc (refer to Figures 4.7& 4.8). ∼

To examine the link between the observed asymmetries in the gas disc of NGC 1566 and ram pressure stripping, we follow the approach proposed by Gunn and Gott(1972) and compare the restoring force by the gravitational potential at the outer disc and the pressure from the IGM. The gas in the outer parts will remain intact to the halo as long as the ram pressure force is lower in magnitude than the pressure from the gravitational potential of the halo. For a disc galaxy with a gravitational potential φ(r) at a distance r from the centre, the restoring force due to this

185 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

3.0 (a) 2.5 2.0

) 1.5 ram

/P 1.0 grav

(P 0.5 10 150 log 0.0 200 0.5 − 300 5 3 1.0 nIGM = 5x10− cm− 400 − 1.5 − 0 5 10 15 20 25 30 35 40 45 50 55 Radius [kpc]

3.0 (b) 2.5 2.0

) 1.5 ram

/P 1.0 grav

(P 0.5 10 5 10− log 0.0 5 2x10− 0.5 5 − 5x10− 1 4 1.0 Vrel = 300 km s− 10− − 1.5 − 0 5 10 15 20 25 30 35 40 45 50 55 Radius [kpc]

Figure 4.16: A comparison between the gravitational pressure and ram pressure ratio in the disc of NGC 1566 for different radii. (a) the pressure ratio assuming an IGM density value of nIGM = 5 3 5 10− cm− and relative velocities for NGC 1566 through the IGM of vrel = 150, 200, 300 and × 1 1 400 km s− . (b) the pressure ratio for constant velocity (vrel = 300 km s− ) and four different IGM 5 5 5 4 3 density values (nIGM = 1.0 10− , 2 10− , 5 10− and 1.0 10− cm− ). The horizontal dashed line represents the equality,× above× which the× gaseous disc is dominated× by gravity and below where ram pressure dominates and the disc is subjected to ram pressure stripping. The vertical line shows the radius at which the HI column density of NGC 1566 is lopsided; At r = 40 kpc the column density of the south-eastern part of the HI disc is nearly zero, while the north-western edge is more extended and smoothly declines with radius up to 50 kpc. ∼

186 4.5. DISCUSSION gravitational potential will exert a pressure given by the following equation (Köppen et al., 2018; Roediger and Brüggen, 2006; Roediger and Hensler, 2005):

∂Φ(r) Pgrav = Σgas(r) , (4.12) ∂z max where the derivative denotes the maximum value of the restoring force at height z above the disc. The height that corresponds to the maximum force is estimated by equating the second derivative

∂2Φ(r) of the gravitational potential ∂z2 , to zero. On the other hand, the IGM will apply a pressure on the outer-most part of the disc that is given by the formula:

2 Pram = ρIGMvrel, (4.13) where ρIGM is the density of the IGM and vrel is the relative velocity of the galaxy. The ram pressure is capable of removing gas from the galaxy as it passes through the IGM only when Pram > Pgrav. For our calculation, we make two assumptions. First, we only calculate the gravitational force due to the dark matter and neglect the restoring force from the stars and gas, i.e., we assume that the contribution from the baryonic matter is negligible. This is only true at the outer radii (refer to Figure 4.13) where the contribution of the dark matter dominates the total halo mass of NGC 1566 (Westmeier et al., 2011). The second necessary assumption is that NGC 1566 moves with an intermediate or “face-on” vector into the IGM, when ram pressure stripping is efﬁcient (Quilis et al., 2000; Vollmer et al., 2001; Roediger and Hensler, 2005) but not directly in the line-of-sight, where morphological effects would be harder to discern 6. Based on the HI morphology and kinematics of this galaxy an encounter with the IGM at an intermediate inclination is favoured. To estimate the restoring force acting on the disc of NGC 1566, we use the NFW dark matter halo gravitational potential which is described by the equation:

GM ln 1 + r 200 rs φNFW(r) = h i. (4.14) −r ln(1 + c) c − 1+c Figure 4.16 presents the results of our simple analytic model, in which the ratio between

Pgrav and Pram is calculated at different radii. In Figure 4.16a, we calculate this ratio assuming

5 3 that the IGM density is constant (nIGM = 5 10 cm ) and use different relative velocities × − − for NGC 1566. This IGM density is similar to the Local Group gas density value (Rasmussen and Pedersen, 2001; Williams et al., 2005). Hence, we note that this adopted IGM density is

6We remind the reader that the inclination of NGC 1566 with respect to the observer is known but the orientation of the disc relative to the direction of motion through the IGM is unknown.

187 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566 a lower limit to the IGM density of NGC 1566 galaxy group. The NGC 1566 group has a halo

13.5 mass (Mhalo 10 M ; Kilborn et al., 2005) larger than the Local Group and consequently ∼ hotter/denser IGM gas is expected (Eke et al., 1998; Pratt et al., 2009; Barnes et al., 2017). In Figure

4.16b, we derive the same ratio adopting a constant relative velocity for NGC 1566 (vrel = 300 km

1 s− ) and different values for the IGM gas density. We use N-body simulations to predict the probability distribution of the relative velocity of NGC 1566 in the IGM following the orbital libraries described in Oman et al.(2013); Oman and Hudson(2016) and based on the projected coordinates of NGC 1566 (angular and velocity offsets) from the group centre (refer to Kilborn et al., 2005, for more information on this group). The relative velocity of a subhalo (galaxy) with a similar mass to NGC 1566 falling into a host halo (galaxy group), with a similar mass to NGC 1566 group, based on this analysis lies in the range between 250 500 km s 1 at 99 per cent conﬁdence. − − Figure 4.16 shows that the outer part of the HI disc of NGC 1566 (r & 40 kpc), for certain values 5 3 for vrel and nIGM, can be affected by ram pressure winds in particular for nIGM 5 10 cm & × − − 1 and vrel & 200 km s− . The highest IGM density adopted for NGC 1566 group and used in Figure 4 3 4.16b (nIGM & 10− cm− ) is consistent and in lower bound of the IGM density values reported for loose galaxy groups and derived from x-ray luminosities in Sengupta and Balasubramanyam (2006); Freeland et al.(2010). As expected for the inner regions the pressure due to the gravitational potential is much higher than the ram pressure force. We note that the ratio in the inner radii is higher than shown in the ﬁgure since, as already noted, we neglect the contributions from the stellar and the gaseous gravitational potentials (refer to Figures 4.12-4.13).

Even though our simple analytic approach suggests that ram pressure interaction with the IGM is the likely reason for the lopsidedness of the gas morphology of NGC 1566, the result is tentative.

This is for two reasons. First, our ASKAP early science observation is not sensitive enough to

18 2 probe HI column densities below 10 cm− , which means that we can not rule out gas accretion as a plausible reason for the asymmetries seen in NGC 1566. The second caveat is that all the environmental processes that affect the HI gas in galaxy groups such as tidal stripping by the host-halo (group), galaxy-galaxy encounters and ram pressure interaction operate at similar radial distances from the group centre. This is to say, as a galaxy approaches the centre of the host-halo and is within a distance of d/Rvir < 0.5, where d is the physical distance from the group centre and

Rvir is the virial radius of the group, it can experience ram pressure stripping from the IGM and/or can also tidally interact with nearby satellites (Marasco et al., 2016; Bahé et al., 2013). This adds to the complexity of disentangling the contributions of different external processes on the HI gas content and morphology in galaxy groups. However, we think that sensitive HI observations of

188 4.5. DISCUSSION large samples of galaxies, similar to those that WALLABY, APERTIF and MeerKAT will deliver in the next few years, will provide the HI community with the opportunity to systematically study the HI gas in group environments and help disentangling the contributions of these environmental processes.

4.5.2 Gas Content and Star Formation Rate in NGC 1566

To demonstrate the high atomic gas content of NGC 1566, we compare the HI-to-stellar mass

9 11.5 fraction of this galaxy to a sample of 25, 000 galaxies within 10 M . M . 10 M and ∗ 0.02 z 0.05 obtained from the Sloan Digital Sky Survey. Brown et al.(2015) reported the ≤ ≤ HI-to-stellar mass fractions of these galaxies using HI data from the Arecibo Legacy Fast (ALFA) survey (Giovanelli et al., 2005). The stellar mass of NGC 1566 is M = 6.5 1010 M , and has ∗ ×

HI-to-stellar mass fraction of log(MHI/M ) = 0.53 0.10. Figure 4.17a presents the HI-to-stellar ∗ − ± mass fraction vs. the stellar mass for Brown+15’s sample and for NGC 1566 (blue circle). It is evident that NGC 1566 has a relatively high HI-to-stellar mass fraction in comparison with its counterparts that have a stellar mass of M 1010.81 M . The average logarithmic HI-to-stellar ∗ ∼ mass fraction of galaxies with M 1010.81 M is 0.97. We note that the uncertainty in the ∗ ∼ − HI-to-stellar mass of NGC 1566 is within the upper bounds of the scatter relation. This galaxy continues to have a relatively higher HI-to-stellar mass fraction than the average at ﬁxed stellar mass even when a smaller distance (16.9 Mpc; the TF distance) is adopted (red circle in Figure 4.17a). Even though, the outskirts of NGC 1566 maybe subjected to ram-pressure interaction with the IGM, this scenario is not inconsistent with a high atomic gas fraction. Ram pressure stripping in group environments is likely to be subtle in comparison with galaxy clusters (Westmeier et al., 2011), in which the density of the intra-cluster medium is orders of magnitude higher than in the IGM (Eke et al., 1998; Pratt et al., 2009; Barnes et al., 2017).

Figure 4.17b shows the relation between the star formation rate and stellar mass for 105 ∼ galaxies in the nearby Universe (z < 0.2) observed in the Sloan Digital Sky Survey (SDSS) and reported in Brinchmann et al.(2004). The dashed-line shows the location of the main sequence of star formation for over 105 galaxies in the local Universe (median redshift z = 0.2) obtained by the Galaxy And Mass Assembly (GAMA) survey and reported in Davies et al.(2016). The blue

1 circle shows the star formation rate of NGC 1566, log(SFRHα[M yr− ]) = 1.33 0.1. This ﬁgure ± highlights the high SFR in NGC 1566. The star formation in this galaxy (Figure 4.18) is mainly concentrated in the nucleus and the inner spiral arms and declines gradually following the outer

189 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566 spiral arms, in a similar fashion to the HI gas. Further, this galaxy has a speciﬁc star formation rate

1 (sSFR = log (SFR/M [yr− ])) of 9.48 0.05, which again places NGC 1566 above the average ∗ − ± with respect to galaxies that have the same stellar mass, but within the scatter (Pan et al., 2018; Abramson et al., 2014). We note that the location of NGC 1566 in the mass-SFR parameter-space is subject to the distance adopted for this galaxy. For instance, if we use the TF distance (16.9 Mpc) instead of the adopted distance in this work (21.3 Mpc), the SFR and the stellar mass of NGC 1566 will be a factor of 0.63 smaller. Thus, the SFR of NGC 1566 will become relatively closer to the median value per ﬁxed stellar mass (red circle in Figure 4.17b), and NGC 1566 will appear less extreme in this parameter-space.

4.6 Conclusion

In this work, we present our ASKAP HI observations of NGC 1566, a grand design spiral in the Dorado group. Our major results and conclusions from this analysis are as follows:

• We measure an HI mass of 1.94 1010 M , assuming a distance of 21.3 Mpc, mainly concen- × trated in the spiral arms of NGC 1566. The HI gas is distributed in an almost regular circular disc that extends well beyond the observed optical disc especially around the northern and

western parts of NGC 1566. The HI gas distribution also highlights the difference between the two outer arms better than in the optical. The eastern outer arm forms a regular arc shape that is more extended than the western arm, which is less regular or disturbed between

PA = 270◦ to 330◦. The HI disc of NGC 1566 is asymmetric: the south-eastern part of the

HI disc sharply declines beyond 30 kpc from the centre, whereas the north-western edge is ∼ more extended and smoothly declines with radius, up to a radius of 50 kpc. ∼

• We measure the rotation curve of NGC 1566 out to a radius of 50 kpc and estimate the dark matter content in this galaxy based on the ISO, NFW and Burkert dark matter halo

proﬁles. We report dark matter fractions of 0.66, 0.58 and 0.62 based on the ISO, NFW and

Burkert proﬁles, respectively. Using our current ASKAP observations, we can not differentiate

between these dark matter density proﬁles as the central region ( 2 kpc) of NGC 1566 is ∼ not resolved. Higher angular resolution observations (few hundred parsec scales) are required for such analysis and for any conclusive ﬁndings (see for example Oh et al., 2015; de Blok et al., 2008; de Blok and Bosma, 2002). Such high angular resolution observations will be

190 4.6. CONCLUSION

Figure 4.17: (a)The HI-to-stellar mass fraction vs. the stellar mass for Brown+15’s galaxy sample (grey) and for NGC 1566 (blue circle). The red circle shows the location of NGC 1566 in the HI-to-stellar mass fraction vs. M plane measured at a distance of 16.9 Mpc (the TF distance). The grey shaded regions delimit the∗ y-axis scatter in the Brown+15’s galaxy sample. (b) The star formation rate vs. stellar mass for SDSS galaxies with z < 0.2; the grey circles show the median value per stellar mass bin and the grey shaded region shows the scatter from the mean values along the SFR axis at a given stellar mass (Brinchmann et al., 2004). The dashed-line shows the location of the main sequence of star formation for the GAMA galaxy sample (Davies et al., 2016). The blue circle shows the location of NGC 1566 in the SFR-M plane. The red circle shows location of NGC 1566 in the SFR-M plane measured at the TF distance.∗ ∗

191 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

Figure 4.18: The SFR map based on the Hα luminosity of NGC 1566 calculated following the SFR calibration of Meurer et al.(2009). The star formation is mainly concentrated in the inner spiral arms with some large HII regions distributed in the outer arms.

achieved in the next coming years using the 36 ASKAP antennas (longest baseline of 6 km) and MeerKAT telescope (de Blok et al., 2016) for nearby galaxies. This will increase the number statistics of highly resolved rotation curves, and consequently better constraints on both the dark and baryonic matter distributions within these galaxies.

• We study the asymmetric HI morphology of NGC 1566 and attempt to discriminate between three major environmental mechanisms that can cause asymmetries in galaxies, namely, ram pressure interactions with the IGM, galactic interactions as well as gas accretion from hosting/neighbouring ﬁlaments. We detect no nearby companion galaxy that may induce

tidal forces on the HI disc of NGC 1566 or tidal tails/plumes that are suggestive of such an

encounter within the last 3.1 Gyr. The HI mass detection limit of the ASKAP observations

1 7 1 based on the 3σ noise level and over 40 km s− channel widths is 2.2 10 M beam− ∼ × at the assumed distance of NGC 1566 (21.3 Mpc). We show, based on a simple analytic

model, that ram pressure stripping can affect the HI disc of NGC 1566 and is able to remove

gas beyond a radius of 40 kpc, using lower-limit values for the gas density of the IGM and the relative velocity of this galaxy. Further, we do not detect any clouds or ﬁlaments

connected to, or in the nearby vicinity of, the HI disc of NGC 1566. However, we are unable

192 4.6. CONCLUSION

to completely rule out gas accretion from the local environment at lower column densities.

Future HI surveys with the SKA precursors and with large single dish telescopes, such as the Five-hundred-meter Aperture Spherical radio Telescope (FAST; Nan et al., 2011; Li and Pan, 2016; Zhang et al., 2019), will help probe the environment around galaxies and quantify the prevalence of gas accretion, interactions and ram pressure stripping in large sample of galaxies and their effects on the atomic gas morphology and kinematics.

• NGC 1566 has a relatively high HI-to-stellar mass fraction in comparison with its coun-

terparts that have the same stellar mass. The average logarithmic HI-to-stellar mass frac-

10.81 tion of galaxies with M 10 M is log(MHI/M )= 0.97. while for NGC 1566 is ∗ ∼ ∗ −

log(MHI/M ) = 0.53 0.1. Further, NGC 1566 possesses a speciﬁc star formation rate ∗ − ± 1 (sSFR = log SFR/M [yr− ]) of 9.48 0.05, which is again above the average with respect ∗ − ± to galaxies that have the same stellar mass, but within the scatter (Pan et al., 2018; Abramson et al., 2014). However, the location of NGC 1566 in the mass-SFR parameter-space is dependent on the assumed distance, for which there remains signiﬁcant uncertainty.

Acknowledgements

We thank the anonymous referee for their positive and constructive comments which greatly improved the presentation of the results in this manuscript. AE is thankful for Kyle Oman for providing the theoretical predictions for the relative velocity PDF of NGC 1566 from his N-body simulations and libraries. AB acknowledges ﬁnancial support from the CNES (Centre National d’Etudes Spatiales, France). JW thank support from the National Science Foundation of China (grant 11721303). This research was supported by the Australian Research Council Centre of Excellence for All-sky Astrophysics in 3 Dimensions (ASTRO 3D) through project number CE170100013.

The ATCA is part of the Australia Telescope National Facility (ATNF) and is operated by CSIRO. The ATNF receives funds from the Australian Government. This work was supported by resources provided by the Pawsey Supercomputing Centre with funding from the Australian Government and the Government of Western Australia, including computational resources provided by the Australian Government under the National Computational Merit Allocation Scheme (project JA3). PS has received funding from the European Research Council (ERC) under the European Union’s Horizon 2020 research and innovation program (grant number 679629;name FORNAX). This paper used archival HI data of NGC 1566 available in the Australia Telescope Online Archive

193 CHAPTER 4. WALLABY EARLY SCIENCE - III. AN H I STUDY OF THE SPIRAL GALAXY NGC 1566

(http://atoa.atnf.csiro.au). ASKAP is part of the ATNF and is operated by CSIRO. The Operation of ASKAP is funded by the Australian Government with support from the National Collaborative

Research Infrastructure Strategy. ASKAP uses the resources of the Pawsey Supercomputing Centre. Establishment of ASKAP, the Murchison Radio-astronomy Observatory and the Pawsey Supercom- puting Centre are initiatives of the Australian Government, with support from the Government of Western Australia and the Science and Industry Endowment Fund. We acknowledge the Wajarri Yamatji people, the custodians of the observatory land. This work used images of NGC 1566 available in the NASA/IPAC Extragalactic Database (NED) and the Digitised Sky Surveys (DSS) website. NED is managed by the JPL (Caltech) under contract with NASA, whereas DSS is managed by the Space Telescope Science Institute (U.S. grant number NAG W-2166). We also used infrared and ultraviolet images of NGC 1566 from the Spitzer Space Telescope and the NASA Galaxy Evolution Explorer websites, both space missions were managed by JPL under contract with NASA.

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Z. Zheng, H. Wang, J. Ge, S. Mao, C. Li, R. Li, H. Mo, D. Goddard, K. Bundy, H. Li, P. Nair, L. Lin, R. J. Long, R. Riffel, D. Thomas, K. Masters, D. Bizyaev, J. R. Brownstein, K. Zhang, D. R. Law, N. Drory, A. Roman Lopes, and O. Malanushenko. SDSS-IV MaNGA: environmental dependence of stellar age and metallicity gradients in nearby galaxies. MNRAS, 465:4572– 4588, Mar. 2017. doi: 10.1093/mnras/stw3030.

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216 Chapter 5

WALLABY Early Science - VI. Imaging the NGC 1566 Galaxy Group with

ASKAP

ABSTRACT

This paper presents Australian Square Kilometre Array Pathfinder, ASKAP, atomic gas (HI) observations of the Dorado 30 deg2 field taken as part of the early science of the Widefield ASKAP

L-band Legacy All-sky Blind surveY (WALLABY). We detect a total of 24 HI sources in this field, seven of which are new HI sources that have no previous known redshift in the literature. All the sources have optical counterparts, down to the 5σ HI column density limit of 0.6 1020 cm 2 × − 7 1 ( 1.2 10 M beam− ). We derive the rotation curves for the five best resolved galaxies in our ' × sample and estimate their corresponding dynamical masses. In comparison with the ALFALFA sample of galaxies, the gas fractions and the star formation rates of the detected HI sources lie within the scatter except for a few outliers. The specific angular momentum of the majority of the detected galaxies follows the analytic relation in Obreschkow et al. (2016), meaning they have

13.5 normal HI content despite being in a relatively massive group (Mhalo = 10 M ). The HI sources that break the relation, exhibit high speciﬁc angular momentum but have low gas mass fractions, are the outlier galaxies in the HI gas fractions vs stellar mass scaling relation. Based on their

HI gas morphology and kinematics, we conclude that these outlier galaxies are affected by their environment through either past tidal interactions or/and possible ram pressure stripping.

217 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP 5.1 Introduction

The current widely-accepted Λ Cold Dark Matter (ΛCDM) cosmological paradigm predicts that the structure of the Universe assembles in a hierarchical fashion, so that massive structures assemble from smaller sub-structure through mergers and accretion events (Press and Schechter, 1974; White and Rees, 1978). At the present time, structures such as galaxy groups and clusters are the most massive objects in this hierarchy and have formed through the infall of galaxies and smaller groups. It is well known that galaxies in the ﬁeld or in voids are different than their counterparts in groups or clusters. These differences are vast and include morphology (e.g., Oemler, 1974; Dressler, 1980; Fogarty et al., 2014), optical colour (e.g., Tempel et al., 2011; McNaught-Roberts et al., 2014), star formation rate (e.g., Poggianti et al., 1999; Kauffmann et al., 2004; Peng et al., 2010), metallicity (e.g., Cooper et al., 2008; Bahé et al., 2017), as well as the atomic gas content (e.g., Haynes and Giovanelli, 1986; Chung et al., 2009). This distinctive difference suggests that there are certain mechanisms in play that affect ﬁeld galaxies while infalling to form massive structures. For instance, infalling galaxies are subjected to mass loss via tidal interactions or harassments (e.g., Moore et al., 1998; Gnedin, 2003; Smith et al., 2015; Rhee et al., 2017), the loss of the gas via ram pressure stripping while approaching the hot centre of their host haloes (e.g., Gunn and Gott, 1972; Abadi et al., 1999; McCarthy et al., 2008; Arthur et al., 2019), via thermal evaporation (Cowie and Songaila, 1977), via turbulent viscous stripping (Nulsen, 1982) and via tidal truncation (Byrd and Valtonen, 1990; Cayatte et al., 1990, 1994). Furthermore, galaxies in clusters lose their extended gaseous halo to this “hostile” environment, which ceases gas accretion into their disc and consequently their star formation is hampered as a result in a process known as starvation (e.g., Larson et al., 1980; Balogh et al., 2000; van de Voort et al., 2017).

Currently, the ‘nature-nurture’ debate, i.e. the evolution of galaxies through secular intrinsic processes (e.g., Boselli et al., 2001) versus external environmental processes (e.g., Boselli and Gavazzi, 2014), is one of the open topics in modern astrophysics. There has been extensive studies of few nearby galaxy clusters that indicate the signiﬁcant role of the environment in transforming galaxies such as the work in the Coma cluster (e.g., Yagi et al., 2010; Fossati et al., 2012; Jáchym et al., 2017; Gavazzi et al., 2018; Cramer et al., 2019), Virgo (e.g., Kenney et al., 2008; Chung et al., 2009; Cortese et al., 2010, 2011; Boselli et al., 2016; Cortese et al., 2016; Boselli et al., 2018), Norma (Sun et al., 2010; Zhang et al., 2013; Jáchym et al., 2014, 2019) and Abell 1367 (Scott et al., 2012, 2013; Gavazzi et al., 2017; Yagi et al., 2017; Scott et al., 2018). For instance, Chung et al.

(2009) conducted a high resolution HI survey for galaxies in Virgo cluster and mapped 53 late-type

218 5.1. INTRODUCTION galaxies distributed in different locations with respect to the cluster’s centre. They found that galaxies that are close to the cluster’s core have truncated (or displaced) HI discs and are signiﬁcantly smaller than their optical discs. This indicates that these galaxies have lost their gas as a result of the stripping occurring due to Virgo’s hot and dense intracluster medium (ICM). In addition to gas stripping, cluster environment is also able to remove dust from infalling cluster galaxies (Cortese et al., 2010; Pappalardo et al., 2012; Corbelli et al., 2012). Cortese et al.(2010) reported that both the dust extent and content of the HI stripped galaxies in Virgo cluster are signiﬁcantly reduced compared to non-stripped galaxies in the same cluster. In contrast, less multiwavelength work is done on high resolution studies of massive galaxy groups and only statistical studies are available. For instance, Brown et al.(2017) studied a large sample of satellite galaxies from the Arecibo Legacy FastALFA survey (ALFALFA, Giovanelli et al., 2005) and the Sloan Digital Sky Survey (SDSS, Abazajian et al., 2009), that reside in different halo mass groups to investigate systematic environmental trends in this sample. Their study reveals that satellites residing in groups with halo

12 masses & 10 M show a noticeable decrease in both the gas content and speciﬁc star formation rate at ﬁxed stellar mass. This emphasise the role of group environments in transforming galaxies and hence detailed studies of groups are essential. Especially since most of the galaxies in the local

Universe reside in groups, ranging from low mass groups, 1012 M , to massive ones, 1014 M . ∼ ∼

We might have reached a consensus of the different mechanisms in play that affect galaxies in different environments. However, the current challenge is to distinguish between the effects of each of these mechanisms and their role in causing the environmental trends seen in observations. This is difficult because some of these mechanisms are coincidental and operate in the same phase-space. For example, tidal stripping or satellite interactions and ram pressure stripping are all in play at the central part of massive groups/clusters (Marasco et al., 2016). Further, different mechanisms are efficient in removing cold gas from galaxies on their own after certain timescales; some are more rapid and violent while others are not. For instance, thermal evaporation is a fast but less violent process and its effects can only be observed if not interrupted by another violent gas removal mechanism such as ram pressure stripping or mergers (Bahé and McCarthy, 2015). To improve on our understanding of the environmental role in the transformation of galaxies and to move forward in resolving the ‘nature-nurture’ debate (see for example, Knobel et al., 2015; Casado et al., 2015; Bremer et al., 2018; Janowiecki et al., 2017; Poudel et al., 2017), large all sky multi-wavelength observational surveys are essential. One of the ongoing surveys that is specifically designed to tackle such questions is the Widefield ASKAP L-band Legacy All-sky Blind surveY (WALLABY) which will be conducted using the Australian Square Kilometre Array Pathfinder (ASKAP). WAL-

219 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

LABY is planned to deliver atomic hydrogen (HI) cubes for more than 600, 000 HI rich galaxies with relatively high angular resolution (3000) mapping about 75 per cent of the sky (Koribalski, 2012; Duffy et al., 2012). This survey will play a key role in enhancing our understanding of the environmental mechanisms affecting galaxies in all environments (clusters, groups and voids), especially because the HI gas is known to be a good tracer of recent interactions in comparison with other probes (e.g., Yun et al., 1994; Verdes-Montenegro et al., 2001b; Michel-Dansac et al.,

2010). Early science observations of WALLABY were completed in 2018, mapping four different ﬁelds. Those are the NGC 7232 group (Kleiner et al., 2019; Lee-Waddell et al., 2019; Reynolds et al., 2019), Fornax group (Murugeshan et al. in prep), M83 group (For et al. submitted) and the Dorado group (see the early results on NGC 1566 galaxy in Elagali et al., 2019).

There are several ways to measure and quantify the environmental effects on galaxies. An example of this is the HI deficiency parameter (e.g., Haynes and Giovanelli, 1984; Magri et al., 1988; Sanchis et al., 2004; Kilborn et al., 2009; Sorgho et al., 2016), which is the difference between the expected and observed HI mass of a galaxy. The expected HI mass of a galaxy is derived from its optical properties, i.e. from the morphology and optical diameter. The HI deficiency parameter strongly correlates with the distance from clusters’ cores, HI deficient galaxies are found closer to the centre, but weak/no correlation exists with respect to other cluster properties such as the X-ray luminosity and temperature (e.g., Solanes et al., 2001). The HI gas morphology and kinematics of galaxies can also serve as a tracer for environmental effects (e.g., Bekki, 2008; Pisano et al., 2011; Elson et al., 2011; Lockman et al., 2012; Mihos et al., 2012; de Blok et al., 2014; Bosma, 2017; Ramírez-Moreta et al., 2018; Elagali et al., 2018; Džudžar et al., 2019; Elagali et al.,

2019). This is because interactions often induce obvious signatures, i.e., HI tails, lopsidedness and warped kinematics, between the participant galaxies which can be used to understand the interaction scenario and develop accurate dynamical models (Bekki, 2008; Mihos et al., 2012; Elagali et al., 2018). In this paper we report on the full ﬁeld of view (FOV) of the Dorado observations and study the properties of the HI sources detected. We use our sensitive observations to investigate the effects of the local environments on the HI gas content, morphology and kinematics of these galaxies.

We also explore properties such as the HI gas fraction, star formation rate and speciﬁc angular momentum in these galaxies. We pay especial emphasis to the latter examining the correlation between HI mass fraction and the global stability parameter, proposed in Obreschkow et al.(2016), and the reasons why some of the Dorado galaxies break this analytic model.

220 5.1. INTRODUCTION ). The 2 − cm 20 10 × 6 . . The mosaic is centred 0 ASKAP column density level of the final cube ( I H detections for galaxies that have no known redshift in the literature. σ I 5 field of view of the Dorado early science field obtained with ◦ 5 . 5 × . The blue contour shows the ◦ 1 5 − . 5 km s 1700 − 700 column density distribution in the I Right hand: The H detections are numbered according to the list in Table 5.1 , numbers with asterisks indicate new H I Figure 5.1: on NGC 1566 and showsH galaxies that lie between line marks the velocity of the brightest galaxy in the group, NGC 1533. virial radius of NGC 1566 galaxy group estimated in Kourkchi and( 2017 ). Tully Left hand: The velocity distribution of the galaxies in the Dorado field, the solid vertical The red filled circle shows the location of the central and two brightest galaxies in this galaxy group, namely NGC 1553/1549. The black circle delimits the projected

221 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

5.1.1 NGC 1566 Group

The NGC 1566 galaxy group is located in the southern Dorado constellation, and it is part of the so-called Dorado loose group. The Dorado group is centred at α = 04:15, δ = 55:41 (J2000) − and has been studied by several authors (e.g., de Vaucouleurs, 1975; Huchra and Geller, 1982; Maia et al., 1989; Ferguson and Sandage, 1991; Kilborn et al., 2005; Firth et al., 2006). It is an unvirialised, extended group of at least 79 galaxies that covers over 10 10 on the sky between ◦ × ◦ the velocity range of 500 1700 km s 1 (Firth et al., 2006). Carrasco et al.(2001) also reported − − the detections of 69 low surface brightness dwarfs, within the V band surface brightness range − of 22.5 25.5 mag arcsec2, that could be part of the Dorado group and are located within the − 2.4 2.4 central part of the group. In addition to the NGC 1566 sub-group, the Dorado group ◦ × ◦ has two further smaller sub-groups, namely the NGC 1433 and NGC 1672 groups (Maia et al., 1989). The two additional subgroups contain eleven and fourteen members, respectively (Kourkchi and Tully, 2017). The Dorado group is thought to be part of the so called Fornax-Dorado-Eridanus ﬁlament (da Costa et al., 1988; Willmer et al., 1989; Omar and Dwarakanath, 2005).

1 The NGC 1566 group contains 31 galaxies, has a projected velocity dispersion of 242 km s− ,

13.5 a projected virial radius of 0.654 Mpc and a halo mass of Mhalo = 10 M (Kourkchi and Tully,

2017). This galaxy group is dominated by two early-type members (brightest members), namely the

S0 galaxy NGC 1553 (Ks = 6.34 mag) and the elliptical E1 galaxy NGC 1549 (Ks = 6.88 mag), refer to Kourkchi and Tully(2017) for more information on the group properties. These two galaxies have faint sharp shells surrounding their stellar discs and a noticeable stellar plume in the western part of NGC 1553, which indicates past interactions between the two systems and/or neighbouring galaxies (Malin and Carter, 1983). There are other interesting and interacting galaxies in this group. For instance, the triplet NGC 1533, IC2038 and IC 2039 have been tidally interacting with each other (Ryan-Weber et al., 2003, 2004; Cattapan et al., 2019). NGC 1533 is an S0 galaxy with a ring of HI gas surrounding it and is the fourth brightest member of the Dorado group

(Ks = 7.74 mag). The other two galaxies are a gas rich late-type galaxy (IC 2038) and a gas poor early type galaxy (IC 2039; the nearest to NGC 1533). Different formation scenarios for the HI ring around NGC 1533 are discussed in Ryan-Weber et al.(2003) and the most likely formation origin is via past minor merger with a low surface brightest gas rich satellite (Ryan-Weber et al., 2004). Ryan-Weber et al.(2003) also report the detection of some H α regions associated with the south-eastern part of the HI ring of NGC 1533. In the most recent study of this triplet using deep images with the ESO VLT survey telescope, Cattapan et al.(2019) show that IC 2039 is undergoing

222 5.2. DATA tidal interactions with both IC 2038 and NGC 1533. Later in the paper, we will discuss in detail this triplet, how the ASKAP map reveals a more extended HI ring associated with NGC 1533 in comparison with Ryan-Weber et al.(2003) as well as the existence of an H I bridge connecting this triplet which conﬁrms the ﬁndings of Cattapan et al.(2019). Other interacting galaxies in this group also include the pair NGC 1602/1596 (Chung et al., 2006) and IC 2058/PGC 75125 (Pearson et al., 2016).

This paper is organised as follows: Section 5.2 briefly describes the ASKAP observations, reduction procedures of the Dorado field as well as the ancillary data used in the paper. We describe the source finding method in Section 5.3. In Section 5.4, we present the main results of the paper, demonstrate the large field imaging capability of ASKAP and discuss the HI gas morphology and kinematics of the interesting HI sources in this field. We also derive the rotation curves for the

five but resolved galaxies by our observations. Section 5.5 discusses properties such as the HI gas fractions, star formation rates and specific angular momenta for the detected HI sources. A summary of our findings is present in Section 5.6.

5.2 Data

5.2.1 HI Observations and Reduction

The Dorado field is one of four early science fields that were observed during the WALLABY early science period. Observations of this field started in December 2016 and were completed in January

2018, utilising between 10 16 ASKAP antennas. Refer to Table 2 in Elagali et al.(2019) for more − details on the observations of this field. We use ASKAPSOFT1 to reduce the data and do the imaging of these HI observations. ASKAPSOFT is an in-house developed pipeline to flag, calibrate, and perform spectral line and continuum imaging of ASKAP observations. Details of this software are available in Kleiner et al.(2019); Lee-Waddell et al.(2019); Reynolds et al.(2019). We follow the same procedures adopted in Elagali et al.(2019) to reduce and image the Dorado field. For brevity, we only mention the important steps in this paper. We flag the data using the ASKAPSOFT task

CFLAG and calibrate using the the task CBPCALIBRATOR. We reduce the entire observation footprint

72 beams (36 in A and B footprints), however only a small part of the frequency bandwidth range is reduced (8 MHz) covering the frequency range of the NGC 1566 galaxy group (1410 1418 MHz; − 511 2196 km s 1). Then, we image the visibilities using the ASKAPSOFT IMAGER task, with a − − 1https://www.atnf.csiro.au/computing/software/askapsoft/sdp/docs/current/ index.html

223 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

3000 taper and a robust weighting value of +0.5. The continuum subtraction of the ASKAP data is done in two steps: the subtraction of the continuum emission from the uv visibilities and the subtraction of the remaining continuum emission in the image domain using the ASKAPSOFT IMCONTSUB task from the ﬁnal HI cube. Finally, we use ASKAPSOFT LINMOS task to mosaic the 72 beams. The ﬁnal HI cube has a restored beam size of 42 36 and rms noise of 0.82 mJy beam 1 per 00 × 00 − 20 km s 1 velocity channel, which is equivalent to a 5σ HI column density value of 0.6 1020 cm 2. − × −

5.2.2 Ancillary Data

We make use of ancillary archival data of the galaxies in the Dorado ﬁeld to determine their stellar masses (M ) and star formation rates (SFR). To begin with, we use the infrared J and Ks-bands ∗ from the VISTA Hemisphere Survey (VHS, McMahon et al., 2013) to estimate the stellar masses presented in Table 5.1. The VISTA J and Ks-band ﬂuxes (FJ and FKs ) can be converted to stellar mass according to the following equations (see Reynolds et al., 2019, and references there in):

2.27 1.135 M = 0.013D FJ , (5.1) ∗ and

2.04 1.02 M = 0.207D FK , (5.2) ∗ s where M is the stellar mass in solar units and D is the distance in megaparsecs. We adopt the ∗ weighted average distance for the NGC 1566 galaxy group D = 17.69 Mpc for all the galaxies in the ﬁeld (Kourkchi and Tully, 2017). The stellar mass reported in Table 5.1 is the average between the J and Ks-band stellar masses derived using the above two equations.

We use the WISE W4 band (22 µm) to derive the star formation rates of our galaxy sample following the calibrations presented in Cluver et al.(2017), which assume a Kroupa(2001) initial mass function (IMF), and using the following equation:

1 log SFR (M yr− ) = (0.915 0.023) log L22µm (L ) (0.801 0.020) (5.3) ± − ± where L22µm is the spectral luminosity ν Lν (22µm). The W4 band absolute magnitude of the Sun is 3.25 mag (Jarrett et al., 2012). Refer to Table 5.1 for the infrared derived SFRs for this galaxy sample.

224 5.3. SOURCE FINDING

5.3 Source Finding

We use the Source Finding Application (SOFIA, Serra et al., 2015) to search and parameterise HI detections in the final cube. In SOFIA, we set the detection threshold to 4σ of the local noise within window size of 100. We run SOFIA in two steps. First we use the HIPASS HI catalogue of the NGC 1566 group (Kilborn et al., 2005) as a prior to search for galaxies in the 30 deg2 field of view of the Dorado early science field. We detect all the galaxies that were previously reported in the HI catalogue of Kilborn et al.(2005) except for one galaxy, IC 2049, which lies outside the

ﬁeld of view of our ASKAP observations. The integrated HI ﬂux density values measured with the ASKAP observations are within the instrumental 10 per cent error of the values detected by

Parkes and reported in Kilborn et al.(2005). The second step after the identification of known H I galaxies in the field, is to conduct a blind search to identify any uncatalogued HI sources. We find ten new HI detections that were not reported in Kilborn et al.(2005), seven of which have no H I or optical redshifts known in the literature. Figure 5.1 shows the HI column density map of the

30 deg2 footprint of the Dorado ﬁeld. In total, we detect twenty-four HI sources in this ﬁeld, with the seven new source marked with asterisks in Figure 5.1 and their spectra shown in Figure 5.2.

Table 5.1 lists some of the HI parameters derived from the source finder among other properties derived from the WISE and VISTA surveys. Galaxies that belong to NGC 1566 group according to Kilborn et al.(2005); Kourkchi and Tully(2017) are labeled with the number “1” in the last column of Table 5.1, whereas those that are not members are given the label “0”. We note that the SOFIA is unable to separate the HI profiles of the galaxy pair NGC 1533/IC 2038 and the triplet NGC 1533, IC2038(2039). These sources are undergoing tidal interactions and their velocity profiles are separated by one (or less than one) velocity channel. We create masks for each of these individual sources and run SOFIA for a second time using the masks to parameterise the HI line properties of these galaxies.

5.4 Results

5.4.1 HI Gas Morphology & Kinematics

Figure 5.3 shows three HI sources identified by SOFIA in the Dorado field, each row corresponds to the same source. The first column presents the HI column density distribution of the ASKAP observations overlaid on the VHS J-band image; the contours are at 2n 5σ level (5σ = 0.6 1020 cm 2; × × − n is a positive integer). The second and third columns are the velocity field and velocity dispersion,

225 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

20 2 respectively, both clipped to show pixels of column density values NHI 0.6 10 cm for each & × − of the these sources. The remaining galaxies are presented in Appendix A.

The ﬁrst galaxy in Figure 5.3 is the interacting triplet NGC 1533, IC 2038 and IC 2039. This triplet has been studied in Ryan-Weber et al.(2003, 2004) and Cattapan et al.(2019). The ASKAP image of the S0 galaxy NGC 1533 reveals a more extended HI distribution in comparison with the ATCA map presented in Ryan-Weber et al.(2003, 2004). We note the existence of an H I bridge connecting the HI distributions of NGC 1533 and IC 2038. Cattapan et al.(2019) report that IC 2039 is also interacting with both NGC 1533 and IC 2038. We see signs of this ongoing interaction in the ASKAP HI map where the aforementioned HI bridge connects these three galaxies.

However, we do not detect HI gas emanating from IC 2039, except for that associated with the

HI bridge. This maybe due to the gas-poor nature of IC 2039. Based on the 5σ sensitivity of

7 1 the ASKAP observations, we estimate an upper HI mass of 1.2 10 M beam− for IC 2039 ' × (measured over 40 km s 1). The HI distribution of this interacting triplet extends for 20 , which is − ∼ 0 equivalent to a physical length of 100 kpc. The second and the third row show the velocity ﬁeld and dispersion maps. The velocity ﬁeld of IC 2038 and NGC 1533 show deviations from symmetry. The isovelocity contours of both galaxies reveal that the approaching side (white contours) is different and more extended in comparison to the receding side (black contours). These asymmetries are expected given the ongoing interaction of this system.

The second galaxy pair is the interacting system, consisting of IC 2058, an edge-on spiral, and its companion dwarf PGC 75125. A tidal HI bridge exists between these two galaxies, which suggests an on-going tidal harassment of PGC 75125 by the larger system IC 2058 (Toomre and Toomre, 1972; Barnes and Hernquist, 1998; Besla et al., 2010; Pearson et al., 2016). The projected physical separation between these two galaxies is 9.2 kpc and they are only separated by 4 km

1 s− in velocity space. Pearson et al.(2016) suggest that due to the proximity of this pair as well as their large stellar mass ratio, the companion dwarf is likely affected by the larger halo of IC 2058 and maybe losing gas to this galaxy via ram pressure stripping and/or tidal stripping. The

HI distribution and the velocity ﬁeld reveal a certain degree of lopsidedness in IC 2058, see also Kregel et al.(2004a). This asymmetry is likely due to the tidal interaction with PGC 75125. The velocity dispersion of this edge-on galaxy is highest in the inner regions and along the minor axis in the direction of its companion dwarf. Overall, the dispersion velocity in the receding side is lower in comparison with the approaching side, which is in agreement with the stellar kinematics (Kregel et al., 2004b). In the outer parts of the HI disc, the velocity dispersion drops quickly.

226 5.4. RESULTS

The third galaxy in Figure 5.3 is a low-surface brightness galaxy ﬁrst detected in (Ferguson and Sandage, 1990), and also reported by Chung et al.(2006). This dwarf galaxy is part of the interacting galaxy pair NGC 1602/1596, refer to Figure 5.1. The physical separation between this dwarf and the galaxy pair NGC 1602/1596 is 25 kpc and the line-of-sight velocity difference ∼ is 73 km s 1. There is an HI bridge connecting this dwarf and the aforementioned pair, which ∼ − suggests that this low surface brightness galaxy is part of the interacting pair. In fact, Chung et al. (2006) speculate that this dwarf has formed as a result of the interaction between the pair NGC 1602/1596, i.e., source seven is a tidal dwarf galaxy which has formed in the tidal tail or the debris resulting from the collision of the pair NGC 1602/1596. Refer to Ploeckinger et al.(2018) and

Lee-Waddell et al.(2012) for more details on the formation of tidal dwarfs. The H I morphology of this dwarf is regular with no signs of asymmetries, while the isovelocity contours shows some signs of warped distribution in the central part (dashed line in the middle panel).

227 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP ] membership ∗ /M I H . 1 log [M -- − ∗

km s log M 1700 ks m − J 700 mag mag M 1 − yr

SFRM m m µ 22

log L L I H

log M M 1 − I H F Jy km s 1 − 20 w km s 1 − 50 w km s 1 − hel (J2000) V δ (J2000) α hh:mm:ss dd:mm:ss km s detections in the Dorado early science ﬁeld between the velocity range of I The H Table 5.1: ID Name123 WALLABYJ043718-5555314a WALLABYJ042252-561340 ESO157-0444b WALLABYJ040705-551920 GALEXASCJ042251.91-561339.7 WALLABYJ040941-5605375a WALLABYJ040854-555932 IC20325b NGC1533 04:22:52 WALLABYJ041755-5555506 IC2038 WALLABYJ041806-5555477 IC2058 -56:13:408 Other WALLABYJ042713-572546 Name PGC 751259 WALLABYJ042730-545558 1347 LSBGF157-08110 WALLABYJ042752-550313 FS9007711 WALLABYJ040607-524018 WALLABYJ040554-540011 NGC160212 04:37:18 43.9 WALLABYJ041059-562906 NGC152213 NGC1515 WALLABYJ041238-57444314 -55:55:31 NGC1536 WALLABYJ041436-56034315 NGC1543 60.8 04:09:41 WALLABYJ041957-545613 04:07:0416 1586 NGC1546 WALLABYJ042708-53500217 NGC1566 04:08:54 WALLABYJ042531-565128 -56:05:3718 04:27:13 -55:19:20 04:18:06 SSEP24J042708.5-535027 04:17:54 WALLABYJ043919-541245 1.39 118.219 ESO157-G027 770.9 -55:59:32 WALLABYJ040956-563115 106520 -57:25:46 ESO157-G047 -55:55:47 WALLABYJ043029-543920 -55:55:5021 GALEXASCJ040955.27-563113.3 04:27:30 738.0 124.6 WALLABYJ042253-563738 118622 04:27:51 1381 232.3 LSBGF157-043 1376 WALLABYJ043045-544813 102.9 04:27:07 8.01 04:06:07 04:09:55 LSBGF157-085 04:05:53 -54:55:58 WALLABYJ041614-554150 -55:03:12 190.0 SSEP24J043045.4-544816 2.01 04:10:59 66.2 319.6 -53:50:02 43.0 -52:40:18 -56:31:15 GALEXASCJ041615.27-554150.4 1496 193.0 04:12:37 115.8 -54:00:11 1565 9.89 04:14:36 -56:29:06 1143 909 200.0 894 04:16:14 1164 04:19:56 -57:44:43 65.9 04:25:31 81.4 209.2 30.6 5.40 45.0 1299 8.17 -56:03:43 04:30:45 84.2 04:39:19 -55:41:50 1168 2.83 -54:56:13 36.8 0.06 -56:51:28 68.8 30.0 360.4 1260 04:30:28 1193 16.72 -54:48:13 -54:12:45 1.22 117.8 1500 46.8 0.66 04:22:53 10.50 985 9.70 154.5 185.6 8.60 -54:39:20 1353 1647 63.2 377.6 328.1 103.6 86.6 -56:37:38 8.32 102.8 15.72 154.8 208.2 905 1.27 9.09 25.30 11.04 192.2 102.6 0.23 7.95 1202 10.15 7.69 41.3 132.4 15.78 14.70 0.61 336.8 5.25 126.0 0.46 10.01 2.62 299.6 7.56 11.03 41.8 6.70 122.6 10.18 35.9 9.27 166.4 9.97 0.70 7.97 58.8 22.66 0.11 0.51 13.92 9.04 180.8 7.65 8.59 0.45 0.05 0.76 7.54 58.1 8.29 14.0 3.68 0.69 10.23 40.0 - 8.69 0.46 0.11 9.22 11.99 9.73 0.07 9.83 8.33 14.26 11.77 7.58 10.13 10.0 11.11 1.11 13.75 0.31 9.78 7.75 14.26 11.01 0.13 8.43 12.24 1 12.52 -0.16 15.93 7.53 12.44 13.92 5.24 10.03 8.20 15.05 10.13 3.28 - 12.17 0.06 0.81 15.43 10.07 0.08 8.37 7.91 10.16 15.15 0.66 7.36 9.08 9.79 12.69 15.98 -0.43 0.40 7.63 13.39 0.08 7.84 9.67 -0.05 1 12.67 13.19 9.86 17.96 0.09 0.01 11.71 - 13.68 0.11 0.32 19.86 9.60 8.88 13.44 9.50 10.09 -0.15 17.44 0.05 11.81 9.06 13.70 19.07 1 18.52 8.60 1 10.18 9.96 6.79 9.54 9.27 8.44 0.06 0.39 9.03 13.63 1 6.32 17.86 13.76 10.01 1 -1.15 10.23 - -0.01 13.77 16.43 1 0.86 6.60 10.44 13.66 1 -0.98 -0.90 1.65 8.43 16.18 -0.79 -1.54 14.30 8.44 -0.31 0.98 1 7.35 14.94 -0.68 1 0 -0.01 16.28 0 8.04 1 0 0.18 1 1 16.17 1 1 0 -0.13 7.37 0 0 -0.01 0 0 0

228 5.4. RESULTS

5.4.2 HI Rotation Curves

We derive the rotation curves for ﬁve galaxies in the Dorado ﬁeld, namely IC 2058, NGC 1566,

IC 2032, NGC 1546 and NGC 1543. These galaxies are resolved (with 4 beams across the ≥ major axis) by our ASKAP observations. We use the tilted ring model approach to estimate the rotation velocity of the HI gas in concentric rings of varying inclination and position angle from the dynamical centre of each of the galaxies. This model is founded on the basis that the gas in each ring is in homogeneous circular rotation and disregards outﬂows and small-scale random motions. Refer to Rogstad et al.(1974) for a thorough description of the tilted ring model and assumptions. We will exploit this model to study large and small kinematic structures in the aforementioned resolved galaxies, explore the reasons for any non-negligible small-scale random motions (if existing) and link that with the interaction history of the galaxies. We use the 3D-Based

Analysis of Rotating Object via Line Observations (3DBAROLO, Di Teodoro and Fraternali, 2015) to derive the kinematic parameters of the five resolved galaxies from their HI spectral line cubes. We provide the algorithm with the position angle, systemic velocity, inclination, and dynamical centre values for each galaxy estimated by SOFIA as priors or initial guesses. Then, 3DBAROLO generates the best kinematic model for these galaxies and outputs the best fit for the rotation curves and other kinematic parameters. The rotation curve, position angle, inclination, modelled moment maps and residuals for four of the five galaxies are provided below. The rotation curve for NGC

1543 has already been derived in Murugeshan et al.(2019), using observations from the ATCA that have similar brightness sensitivity and angular resolution to this ASKAP observations. Our 3D

BAROLO results of NGC 1543 are in agreement with Murugeshan et al.(2019) results, hence for brevity we refer the interested reader to their paper.

The residual maps presented below for the four resolved galaxies are a good diagnostic to identify any asymmetries present in these galaxies due to internal processes as well as any signatures of past interactions with their environments. The ﬁrst galaxy presented is IC 2032 (Figure 5.4). This is a gas rich peculiar dwarf galaxy of the class IAB(s)m (Carrasco et al., 2001; Lee et al., 2003). The optical image of IC 2032 reveals a very irregular ‘comet-like’ morphology; the western part contains knots of star-forming regions whereas the eastern part has fainter extended tail (Lee et al., 2003). There are clear signs of asymmetries in the velocity ﬁeld of this galaxy; the isovelocity contours seen in IC2032 (source 3 in Figure 5.11) of the receding and approaching sides are different, which suggests kinematic asymmetries. Overall, the rotation velocity in the receding side is lower than the approaching side (refer Figure 5.4). The HI morphology is slightly more

229 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

Figure 5.2: The integrated HI spectra of the new seven low surface brightness galaxies detected using ASKAP. The vertical line shows the systemic velocity of each galaxy estimated by SOFIA. The seven spectra are numbered in accordance with the list in Table 5.1.

230 5.4. RESULTS

Figure 5.3: Three HI sources identified by SOFIA in the Dorado field; each row corresponds to the same source. The galaxies are number in accordance with Table 1.1. The first column presents the HI column density distribution of the ASKAP observations overlaid on the VHS J-band; the n 20 2 contours are at 2 5σ level (5σ = 0.6 10 cm− ; n is a positive integer). The second and third columns are the velocity× field and velocity× dispersion, respectively, both clipped to show pixels 20 2 of column density values NHI & 0.6 10 cm− for each of the three sources. The remaining galaxies are discussed in Appendix A.×

231 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP extended in the receding side of the galaxy; the HI column density contours are more extended in the eastern part of the galaxy as oppose to the western part. Further, the residual ﬂux and velocity maps show a non-negligible amount of gas deviating from circular rotation, especially along the minor axis of the galaxy. This is more evident in the position-velocity (PV) diagram taken along the minor axis, where the observation (blue contours) and the model (red contours) slightly differ.

This unusual optical and HI morphology could not be a result of interactions with neighbouring group members, as the nearest group members are the interacting triplet NGC 1533/IC2038/39.

IC2032 is separated by over 200 kpc from this triplet, which makes this scenario less favourable. A different scenario is that this dwarf has passed through the centre of the group and experienced ram pressure winds that have created this ﬁreball of star formation at one end and a fainter end on the other side. The maximum rotation velocity of the gas at r 6 kpc is 67.5 km s 1, which under the ∼ − assumption of spherical symmetry corresponds to a total dynamical mass of 6.3 109 M . Hence, × the HI gas mass to total mass of IC 2032 makes up only 0.06 of the total mass. This means that this galaxy is dark matter dominated, which is expected since this IC 2032 is a dwarf galaxy.

The second galaxy is IC 2058 (Figure 5.5), an edge-on spiral galaxy of the class Scd (Lauberts, 1982; Kregel and van der Kruit, 2005). As discussed in the previous section, this galaxy is interacting with a neighbouring dwarf and an HI bridge connects the two galaxies (Pearson et al., 2016). In the derivation of the kinematics of this pair, we only consider the main galaxy IC 2058 and mask the dwarf companion. Overall, the rotation velocity in the receding side is lower in comparison with the approaching side, which is in agreement with the stellar kinematics of IC 2058 (see also Kregel et al., 2004b,a). Both the ﬂux and velocity residual maps show a considerable amount of gas deviating from circular motion. This is expected because IC 2058 is ongoing tidal interaction with its companion. Further, this galaxy pair is located close to the group centre and the large early type galaxy NGC 1553. Pearson et al.(2016) speculate that the IC 2058 truncated H I disc is likely due to its proximity to the group centre and that this pair may have been subjected to ram pressure stripping. IC 2058 has a maximum rotation curve of 111.5 km s 1, at r 11.1 kpc, − ∼ which translates to a total dynamical mass of 3.2 1010 M . The HI mass contribution to the total × dynamical mass in IC 2058 is only 4 per cent. The third galaxy NGC 1546 (Figure 5.6), is a highly inclined spiral galaxy of the type Sbc (Sandage and Brucato, 1979). The optical image of this galaxy reveals a dust lane across its edge. The HI morphology and velocity ﬁelds seem regular except for the 2 kpc tail that extends from the outer disc along the south-eastern parts. However, there is a considerable fraction of gas in the inner parts of the galaxy that has random non-circular motion. This is clearer in the PV diagram taken across the minor axis, where the model fails to account

232 5.4. RESULTS

Figure 5.4: Upper panel: The observed and modelled moment zero as well as the residual maps (observed - model). Middle panel: The observed and modelled velocity ﬁelds as well as the residual maps. Bottom panel (left): The inclination (i), position angle (PA) and the rotation curve (Vrot) derived using 3DBAROLO. Bottom panel (right): the position-velocity (PV) diagram along the major and minor axes of the galaxy. The contours are at (2, 4, 8, 16, 32) σ level, where blue contours show the observations and red contours the model. The magenta contours× represent the 2σ noise level in the observed data. −

233 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

Figure 5.5: Same as Figure 5.4.

234 5.5. DISCUSSION for all the observed gas. The tail seen in NGC 1546 could be the result of tidal interaction with a neighbouring dwarf galaxy, namely AM 0413-560. This dwarf is located along the same direction of the tail and is between this spiral galaxy and NGC 1533 (refer to Figure 2 in Carrasco et al.,

2001). The rotation velocity at the outermost radius of NGC 1546, r 6 kpc, equals 159.7 km s 1 ∼ − and corresponds to a total mass of 3.6 1010 M . ×

The last galaxy NGC 1566 (Figure 5.7), is a face-on spiral galaxy of the type SAB(s)bc (de Vaucouleurs, 1973, 1975). Elagali et al.(2019) studied the asymmetric gas morphology of this galaxy and found evidence that the lopsidedness of its HI disc is due to ram pressure stripping from the intergalactic medium of this group. There are regions with large velocity residuals in the edge of HI disc of this galaxy, and are most noticeable in the eastern side of the disc. This side, according to Elagali et al.(2019), is the most affected by ram pressure winds from the intergalactic medium. Further, the residual ﬂux map shows non-negligible amount of gas in the nucleus and the inner spiral arms of NGC 1566. This residual gas is characterised by non-circular motions

20 2 and has column density that reaches NHI 3.0 10 cm (equivalent to a surface densities of . × − 2 . 3.0 M pc− ) in the nucleus and the inner spiral arms (Figure 5.7). This is nearly 10 per cent of the local gas densities in these inner regions of the galaxy. This large amount of gas could be a result of outﬂows resulting from the Seyfert nucleus in NGC 1566. Such ouﬂows are detected in CO observations in this galaxy (for more details refer to Combes et al., 2014; da Silva et al., 2017). An alternative scenario for the presence of non-circular gas motions in the inner regions of NGC 1566 is through past gas accretion from dwarf companion(s).

5.5 Discussion

5.5.1 Star Formation and HI Gas Scaling Relations

To explore the properties of the galaxies in the Dorado ﬁeld, we compare them with representative sample of galaxies in the nearby Universe. Figure 5.8 (upper panel) shows the stellar mass versus the speciﬁc star formation rate (sSFR=SFR/M ) for galaxies in the ALFALFA survey in blue color ∗ and the contours show the 1.0 σ and the 2.0 σ of the distribution (Giovanelli et al., 2005; Haynes et al., 2011, 2018). The black circles show the extended GALEX Arecibo SDSS Survey galaxy sample (xGASS, Catinella et al., 2018). The star formation and the stellar mass of these galaxies were derived from the GALEX FUV and NUV -band and the SDSS ugriz bands spectral energy distributions assuming the Kroupa(2001) IMF. For more details refer to Brinchmann et al.(2004)

235 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

Figure 5.6: Same as Figure 5.4.

236 5.5. DISCUSSION

Figure 5.7: Same as Figure 5.4.

237 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

Figure 5.8: Upper panel: the speciﬁc star formation rates (sSFR=SFR/M ) versus M for the ALFALFA galaxy sample (blue circles), xGASS galaxy sample (black circles)∗ and the∗ Dorado galaxies (red stars). Lower panel: the star formation efﬁciency (SFE=SFR/MHI) versus M for the ∗ three aforementioned galaxy samples using the same color code. The contours are the 1.0 σ and the 2.0 σ boundaries of the ALFALFA distribution.

238 5.5. DISCUSSION and Huang et al.(2012). We note that both the star formation rates and stellar masses for the the ALFALFA, xGASS and Dorado galaxy samples were estimated using the same IMF. The majority of our galaxy sample lie within the 2.0 σ scatter, with four outliers that possess higher sSFR but within the scatter of the total ALFALFA sample. These sources are, in ascending order of stellar mass, 22, 15, 9 and 13. The first two galaxies are low surface brightness dwarfs and the latter are an S0 and a peculiar type galaxy, respectively (Kilborn et al., 2005). Figure 5.8 (lower panel) shows the stellar mass versus the star formation efficiency (SFE=SFR/MHI) for the same xGASS and ALFALFA samples, with the contours showing the 1.0 σ and the 2.0 σ boundaries of the distribution. Half of the Dorado galaxies (red stars) exhibit star formation efficiency more than

2.0 σ above the ALFALFA sample. Five of these galaxies are dwarfs with M 108 M and high ∗ ≤ atomic gas masses. The late-type galaxy NGC 1546 possesses the highest SFE in this sample and a

1 SFR 17 M yr− . ∼

Figure 5.9 (upper panel) shows the gas fraction (MHI/M ) versus M for the galaxies in the ∗ ∗ ALFALFA sample overlaid with the 24 galaxies detected with ASKAP. The contours show the 1, 2 and 3 σ scatter of the ALFALFA sample. The majority of the Dorado HI sources are within the

3 σ scatter of the ALFALFA sample, albeit ﬁve lie near the lower limit. These HI sources are 11 (type Sbc peculiar), 12 (type (R)Sb), 16 (dwarf irregular), 18 (dwarf irregular) as well as 5b (dwarf irregular). Source 11 (or NGC 1536) is in close proximity to the interacting triplet NGC 1533/IC

2038(2039), which makes gas removal via tidal interactions a possible scenario. The HI disc of

NGC 1536 shows no signs of tails and appears slightly truncated with the HI disc as extended as the optical disc in the North and South directions (vertical in Figure 5.11) and more extended in the East and West directions (left-right). Source 12 (or NGC 1543) has a ring of HI gas around its nucleus, and very faint HI emission emanating from the centre and/or the bar. This HI gas traces a low-surface-brightness ring located at the edge of the bar. The origin of this peculiar HI morphology is likely a drop-through interaction with a lower mass group member (Murugeshan et al., 2019). The three remaining HI deﬁcient dwarfs are companions of (or in close proximity to) a more massive galaxy, which hints for gas losses via tidal interactions. Figure 5.9 (lower panel) shows the location of the Dorado galaxies in the gas fraction versus M plane, together with the ∗ average HI gas fractions for satellite galaxies that reside in different halo masses (Brown et al.,

2017). We remind the reader that the halo mass of NGC 1566 group is 1013.5 M (Kourkchi and

Tully, 2017). The scatter of the NGC 1566 group galaxies in this parameter space compared to the average gas fractions for satellites that reside in a similar halo mass is signiﬁcant. Several group members have orders of magnitude lower (or higher in the case of NGC 1566) gas fractions

239 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP compared to the average gas fraction of their counterpart satellite galaxies at ﬁxed stellar mass. This emphasises the importance of high resolution observations in order to fully understand the physical processes that cause this scatter, which the HI stacking experiment of Brown et al.(2017) is not able to offer. Such high resolution HI imaging of groups for the entire sky will be available in few years thanks to the imaging capability of ASKAP as well as the other SKA precursors that have currently started operating.

According to Verdes-Montenegro et al.(2001a), there are three evolutionary phases for the H I gas morphology and kinematics transformation in galaxy groups. In the ﬁrst phase the HI gas in the group members is unperturbed and is within the discs of these galaxies. In the second phase more asymmetries and HI features, such as tails and bridges, are present between the group members; in the third phase most of the HI gas is stripped from the group members and is found (if detectable) in tails between these galaxies or in the group medium. Hence, it is expected based on the unvirialised dynamical nature of this group that only one third of the Dorado group members show clear signs of

HI disturbances and only ﬁve are below the 3 σ contours of the ALFALFA sample. In comparison with the the three stages proposed in Verdes-Montenegro et al.(2001a), the H I gas in the Dorado group is possibly in phase two of its transformation.

5.5.2 The Stability Parameter and the Atomic Gas Mass Fraction Relation

Obreschkow et al.(2016) found a correlation between the atomic gas (H I+He) mass fraction (fatm) for isolated disc galaxies and their speciﬁc angular momentum; galaxies that have higher atomic gas mass fractions tend to have higher speciﬁc angular momentum as opposed to galaxies with lower atomic gas mass fractions. This correlation emerges from the fact that if the atomic gas in disc galaxies is Toomre stable, this gas will not collapse and form molecular gas and stars.

Hence, in isolation a galaxy can sustain as much HI as its speciﬁc angular momentum allows.

Naturally, galaxies with high angular momentum can then sustain larger HI reservoirs, leading to higher atomic gas fractions. Following these results, Butler et al.(2017) studied a sample of dwarf irregular galaxies and found that these systems have relatively high speciﬁc angular momentum than expected from the known angular momentum-total mass scaling relation of spiral galaxies (e.g., Fall and Efstathiou, 1980; Romanowsky and Fall, 2012; Obreschkow and Glazebrook, 2014) and linked this deviation to the low baryonic mass fraction of dwarf galaxies. To expand on Obreschkow et al. (2016), there has been recent studies on the relation between angular momentum and the gas content of galaxies in different environments. For instance, Lutz et al.(2017, 2018) studied the gas content and speciﬁc angular momentum of a sample of HI rich galaxies and found that these galaxies have

240 5.5. DISCUSSION

Figure 5.9: Upper Panel: The stellar mass versus the gas fraction for all galaxies in the Dorado ﬁeld identiﬁed by ASKAP. The Dorado galaxies are numbered in accordance with Table 1. Overlaid in gray and blue circles are the ALFALFA and xGASS galaxy samples, respectively. The contours are at 1, 2 and 3 σ boundaries of the ALFALFA sample. Lower Panel: The average HI gas fractions for satellite galaxies that reside in different halo masses, these values are adopted from Brown 13.5 et al.(2017). For comparison, the halo mass of NGC 1566 is M halo 10 M , and the Dorado galaxies are overlaid in this panel as well. ∼

241 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP higher speciﬁc angular momentum in comparison with galaxies that lie within the same stellar mass range but have lower HI mass fractions. They argue that their sample is therefore more HI rich than other galaxies of the same stellar mass because of the high speciﬁc angular momentum of these galaxies which helps in sustaining more HI mass. Murugeshan et al.(2019) studied six

HI deficient galaxies in low density environment and found no tracers of harassments nor ram pressure stripping acting on these galaxies. Murugeshan et al. analysed these same galaxies in the q fatm plane and found that their HI contents were normal for their specific angular momentum. − Hence, they concluded that the environment is not responsible for their apparent HI deficiency, but instead they are a sample of low specific angular momentum galaxies. Most recently, Džudžar et al. (2019) studied a sample of Choir galaxy groups to find that their high specific angular momentum is the only reason behind their high gas mass fractions, and that the environmental effects are less prevalent in these choir groups.

The q fatm parameter-space can be used as a diagnostic tool to study galaxies in different − environments in order to disentangle the role of the internal physical properties (nature) and the environmental factors (nurture) in driving the transformation of galaxies in different environments. Here, we study the angular momentum of the galaxies in this group to explore the nature of the relationship between the total angular momentum and the gas mass fractions in these galaxies. For all the galaxies in our sample we estimate the atomic gas mass fraction using (Obreschkow et al., 2016):

1.35M f = HI , (5.4) atm M and the disc stability parameter (q):

jσ q = (5.5) GM where j is the speciﬁc angular momentum of the galaxies, G the Gravitational constant, M is the total mass and σ is the velocity dispersion of the HI gas. The total baryonic mass in the disc is deﬁned by M=M +1.35(MHI+MH2 ), where MH2 is the molecular gas mass and is approximated to ∗ be 0.04 of the total mass for galaxies with no molecular gas observations. The helium gas mass is accounted for in the previous equation by the factor 1.35. We adopt a constant value for the

1 velocity dispersion for all the galaxies in our sample σ = 10 km s− (Obreschkow et al., 2016). This adopted value is the average velocity dispersion in nearby disc galaxies found in different studies (Dickey et al., 1990; Leroy et al., 2008). We measure the speciﬁc angular momentum in the

242 5.5. DISCUSSION

ﬁve resolved galaxies in our sample using the following:

Σi(1.35MHI,i + M ,i)vrot,iri j = ∗ (5.6) Σi(1.35MHI,i + M ,i) ∗ where vrot is the rotation velocity at a radius ri estimated in the previous section. For the remaining unresolved galaxies, we use the approximation in Obreschkow et al.(2016) to determine their speciﬁc angular momentum:

j = rdiscvmax, (5.7)

where rdisc is the radius of the optical disc and vmax is the maximum rotation of the gas measured using the inclination corrected w50 values.

Figure 5.10 presents the location of all the Dorado ﬁeld galaxies in the q fatm parameter − space. The ﬁve resolved galaxies are shown by the coloured stars. We measure the q and fatm for these galaxies at three different radii: at a radius containing 50 per cent of the optical light (r50), at a radius containing 90 per cent of the optical light (r90) as well as at the total extent of the HI

2 disc (HI surface densities = 1 M pc− ). For the resolved galaxies, the gas mass fraction changes significantly with different radii, by more than an order magnitude in the case of NGC 1566. The angular momentum also follows a similar trend and increase monotonically with radius, especially for NGC 1566, NGC 1546 and IC 2032. In general, most of the galaxies in our sample agree within the scatter of the analytic correlation between the q-parameter and the gas mass fraction derived in Obreschkow et al.(2016). However, five galaxies in our sample clearly deviate from this correlation, namely, the HI sources 1, 11, 16, 18 and NGC 1543. These five galaxies have relatively high angular momentum but have lower gas mass fractions than predicted by the analytic relation (Obreschkow et al., 2016). This points to these galaxies being affected by interactions with the surrounding environment. Stevens et al.(2018) explored the connection between the angular momentum and gas fraction in the DARKSAGE semi-analytic model (Stevens et al., 2016; Croton et al., 2016) and the reasons leading to some galaxies deviating from the analytic prediction of

Obreschkow et al.(2016). Galaxies in DARKSAGE deviate from the q fatm parameter space for − two reasons, minor mergers and/or ram pressure stripping (Stevens et al., 2018). In our opinion, past tidal harassments could also contribute to the breaking of this connection. The HI sources 11 and 18 are in close physical proximity to the interacting triplet NGC 1533, IC 2038 and IC 2039 and might have interacted tidally with this triplet in the past which might have caused the low gas mass fraction seen in these two HI sources. The sources 1 and 16 are irregular dwarfs with

243 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

Figure 5.10: The q fatm parameter space for all the galaxies in the Dorado field. Overlaid are the HIPASS, THINGS− and LITTLE THINGS galaxies obtained from Obreschkow et al.(2016) as well as the HI deficient galaxies presented in Murugeshan et al.(2019). The black squares show the Virgo cluster VIVA sample of galaxies (Chung et al., 2009) obtained from Li et al. (2019, submitted). The grey shaded region shows the analytic correlation between the q-parameter and the gas mass fraction derived in Obreschkow et al.(2016). noticeably asymmetric HI and velocity field morphologies, which is an indication of either past interactions or ram pressure stripping in this group (Elagali et al., 2019). The last galaxy in the

Dorado ﬁeld that breaks the q fatm parameter space is NGC 1543 (see also Murugeshan et al., − 2019). This is a late-type barred galaxy with an HI ring surrounding the nucleus. Murugeshan et al. (2019) speculate that the NGC 1543 ring formed as a result of a collision interaction while on its ﬁrst passage through the group. This minor merger scenario can also explain the reason for the low gas mass fraction of NGC 1543.

5.6 Conclusion

In this work, we present the full field of view of ASKAP HI observations of the Dorado field obtained as part of the early science observations of WALLABY. Our main findings are summarised below:

• In the Dorado 30 deg2 ﬁeld, we detect a total of 24 HI sources in the velocity range between

700 1700 km s 1. Seven of these detections are new HI sources that have no previous − − known redshift in the literature. We detect no HI clouds or sources that have no optical

6 1 counterparts, down to the observations 5σ HI mass detection limit ( 4.0 10 M beam− ) ' ×

244 5.6. CONCLUSION

estimated at the adopted distance of the Dorado group.

• We detect HI bridges between two triplets NGC 1533/IC 2038(2039) and NGC 1602(1595)/ FS90077 as well as the galaxy pair IC 2058/PGC 75125, which indicates an ongoing tidal interactions between these systems. These galaxies are known to be interacting in the literature,

however ASKAP HI images show deeper and more extended emission associated with them.

For example, our HI image of the triplet NGC 1533/IC 2038 (2039) shows a more extended

bridge connecting this triplet and reveals an HI emission that extends for 200 associated with this system.

• We use 3DBAROLO to derive the rotation curves for five resolved galaxies ( 4 beams ≥ across the major axis) in the Dorado field, namely IC 2058, NGC 1566, IC 2032, NGC 1546 as well as NGC 1543, and estimate their corresponding dynamical masses. We explore non-circular or random kinematic structures associated with these galaxies and relate this with the interaction history and/or the internal properties of these galaxies. For instance, we find signs of kinematic asymmetries in the outer parts of the dwarf galaxy IC 2032.

The residual ﬂux and velocity maps show non-negligible amounts of gas (& 3σ) deviating from circular rotation, and the isovelocity contours of the receding and approaching sides

are different. Also, the HI morphology in the receding side of the galaxy is slightly more extended. Since the nearest system to this dwarf is the triplet NGC 1533/IC2038(2039), over

200 kpc physical separation, tidal interactions are less likely, so this dwarf may have passed through the centre of the group and experienced ram pressure winds. This stripping scenario also explains the ‘comet-like’ appearance with a ‘ﬁreball’ of star formation at one end and a fainter tail.

• In comparison with the ALFALFA sample of galaxies, the HI gas fractions and the star

formation rates of the detected HI sources are within the scatter except for a few outliers. This is reasonable since the Dorado group is unvirialised and only one third of the Dorado

group members show clear signs of HI disturbances.

• The speciﬁc angular momentum of the majority of the detected galaxies follows the analytic

relation in Obreschkow et al. (2016). The HI sources that break the relation, exhibit high

245 CHAPTER 5. WALLABY EARLY SCIENCE - VI. IMAGING THE NGC 1566 GALAXY GROUP WITH ASKAP

speciﬁc angular momentum but have low gas mass fractions, are galaxies affected by their environment through either through past tidal interactions or/and possible ram pressure stripping.

Acknowledgements

AE Acknowledges the funds he received from the International Centre of Radio Astronomy (ICRAR). This research was supported by the Australian Research Council Centre of Excellence for All-sky Astrophysics in 3 Dimensions (ASTRO 3D) through project number CE170100013. The

Australian SKA Pathﬁnder is part of the Australia Telescope National Facility which is managed by CSIRO. Operation of ASKAP is funded by the Australian Government with support from the

National Collaborative Research Infrastructure Strategy. ASKAP uses the resources of the Pawsey Supercomputing Centre. Establishment of ASKAP, the Murchison Radio-astronomy Observatory and the Pawsey Supercomputing Centre are initiatives of the Australian Government, with support from the Government of Western Australia and the Science and Industry Endowment Fund. We acknowledge the Wajarri Yamatji people as the traditional owners of the Observatory site. This publication makes use of data products from the Wide-ﬁeld Infrared Survey Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration. The VHS data was combined and analysed to measure source photometry using the VISTAview tool2.

2http://vistaview.icrar.org/

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Figure 5.11: The galaxies identified by SOFIA in the Dorado field, each row corresponds to the same galaxy. The galaxies are number in accordance with Table 5.1. The first column presents the HI column density distribution of the ASKAP observations overlaid on the VHS J-band; the n 20 2 contours are at 2 5σ level (5σ = 0.6 10 cm− ; n is a positive integer). The second and third columns are the velocity× field and velocity× dispersion, respectively, both clipped to show pixels of 20 2 column density values NHI 0.6 10 cm for each of the detected galaxies. & × −

264 5.A. GAS MORPHOLOGY AND KINEMATICS OF THE REMAINING GALAXIES

Figure 5.12: Same as Figure 5.11.

265 BIBLIOGRAPHY

266 Chapter 6

Conclusions and Summary

This PhD thesis presents detailed studies of the environmental processes that take place in galaxy groups, i.e., galaxy collision, tidal interactions as well as ram pressure stripping, and how each of these processes affect the atomic gas morphology and kinematics as well as the star formation rates in group galaxies. Understanding the physical processes operating in groups is of a fundamental importance for our understanding of galaxy evolution as whole, especially because most of the galaxies reside in groups and only few galaxy groups are studied in the literature. We use the state-of-the-art ASKAP radio telescope as well as the ATCA to map the atomic gas in two galaxy groups and explore similar groups in the EAGLE hydrodynamical simulations in order to confront the theory with observations.

In the first part of this thesis, we investigate the ramifications of galaxy-galaxy collisions and minor mergers, leading to the formation of the so-called ring galaxies, on the star formation and the neutral gas (atomic and molecular) of the participant galaxies. For the first time, we are able to offer a consistent physical interpretation of long standing questions about these galaxies, such as why do they appear to form stars so inefficiently and how their abundance evolve through cosmic time. Dif- ferent to other merger scenarios, we show that the drop-through collision leading to the formation of ring galaxies causes the quenching of the star formation in timescales less than 500 Myr and the transformation of gas rich discs to green valley galaxies. These results also motivated new searches for ring galaxies in the distant Universe and the discovery of the most distant ring galaxy to date, refer to Appendix A, at a lookback time of 10.8 Gyr (z = 2.19). Similar studies of distant ring galaxies can further advance our understanding of star formation and dynamics in young galaxies, as well as constrain the cosmic evolution of large spiral disks and galaxy groups.

267 CHAPTER 6. CONCLUSIONS AND SUMMARY

Besides galaxy collisions, we also explore the effects of less violent interactions, i.e., tidal interactions and ram pressure stripping, on group galaxies. We find evidence for moderate, but continuous, ram pressure stripping in the NGC 1566 galaxy group. This new discovery affirms the influence of the intergalactic medium on the outer HI discs of galaxies, which originally thought to be negligible in comparison with the influence of the intracluster medium on cluster galaxies. Also, several members of this group are undergoing tidal interactions with each other, and have tidal tails and/or bridges linking them. Exploring similar cases of ram pressure effects on group galaxies is now possible, thanks to the advancement of the ASKAP, MEERKAT and APERTIF instruments.

Comparing the NGC 1566 group with a large sample of HI sources in the nearby Universe, we ﬁnd that this group host several HI "poor" galaxies. These galaxies are either undergoing interactions or are companions of (or in close proximity to) a more massive galaxy, which hints at gas losses via tidal interactions. This agrees with previous studies which ﬁnd satellites that reside in

12 groups, with halo masses & 10 M , showing a noticeable decrease in both the gas content and the specific star formation rate at fixed stellar mass (Brown et al., 2017). Unlike the stacking results in Brown et al.(2017), we are able to explore the environmental and physical processes that lead to the gas deficiency in this group. Hence, the results presented in this thesis emphasise the importance of high resolution observations in order to fully understand the physical processes that transform group galaxies. Such high resolution HI imaging of groups for the entire sky will be available in few years thanks to the imaging capability of ASKAP as well as the other SKA precursors that have currently started operating. Below, we provide a summary of the important results of the preceding Chapters in this thesis and presents potential directions for future work.

6.1 The Collisional Pair NGC 922 & S2

In Chapter 2 (published in Elagali et al., 2018b), we present an HI study of the collisional galaxy pair NGC 922 and S2. We ﬁnd the HI morphology and kinematics of NGC 922 challenging to explain based on the simple drop-through collision proposed in Wong et al.(2006). Our new observations suggest that an additional tidal interactions must have occurred between this pair before or after the plunge of the intruder S2 through the disc of NGC 922 and hence provide further constraints for future simulations and modelling of NGC 922. This study shows that NGC 922 has a star formation rate that is high relative to galaxies of the same stellar mass. Since galaxy interactions and collisions are more common at higher redshifts, this suggests that collisional ring

268 6.2. COLLISIONAL RING GALAXIES IN EAGLE galaxies may account for a greater fraction of burst-dominated star forming galaxies than they do in the local Universe. Further, we also show that NGC 922 and other nearby collisional ring galaxies have higher HI gas fractions in comparison with galaxies that have similar stellar mass.

6.2 Collisional Ring Galaxies in EAGLE

In Chapter 3 (published in Elagali et al., 2018a), we use the EAGLE simulations to study the formation, evolution and characteristics of collisional ring galaxies through cosmic time. We predict the formation density rate of collisional ring galaxies up to redshift z = 2.5. This density rate is in broad agreement with the observations of ring galaxies (Lavery et al., 2004; Elmegreen and Elmegreen, 2006). This means that the numerical treatment of the ISM, star formation, and feedback model in the EAGLE simulations is adequate enough to reproduce a realistic ring-morphology population, which is especially outstanding because the calibration of the subgrid physics in these simulations did not include galaxy morphology.

We find ring galaxies in massive groups that are more concentrated than groups without ring galaxies at a fixed halo mass. Further, the majority of EAGLE ring galaxies are formed when a companion galaxy drops through the disc of a gas-rich galaxy (sometimes more than one drop-through). We find that the ring morphology induced by the collision remains visible in the simulation for around 700 Myr and after that the gas and star particles collapse back to the inner regions of ∼ the target galaxy. In the case of multiple encounters, the lifespan of the ring feature increases significantly and can reach up to 1400 Myr or by a factor of 2 in comparison with other rings ∼ ∼ which formed as a result of single encounters. Further, the gas kinematics of these simulated rings show signs of outwardly expanding velocity components in addition to the rotation velocity, in agreement with the predictions of the analytic caustic theory (Appleton and Struck-Marcell, 1996; Struck-Marcell and Lotan, 1990; Struck, 2010; Mapelli and Mayer, 2012) and the observations of ring galaxies (Conn et al., 2016; Fogarty et al., 2011; Higdon, 1996).

EAGLE ring galaxies have moderate star formation rates even though they have high HI gas fractions. These galaxies are inefﬁcient in forming stars because their ISM is characterised by much lower gas-phase pressure and gas-phase metallicity in comparison with star-forming galaxies that have similar stellar and gas masses. We show that the drop-through interaction is responsible for decreasing the ISM pressure internal to the ring as the ISM is swept up by the density wave outwards. On the other hand, at the radius of the ring, the star formation efﬁciency of ring galaxies

269 CHAPTER 6. CONCLUSIONS AND SUMMARY is indistinguishable from that of normal star-forming disc galaxies, and thus the main reason for their integrated lower efﬁciency is the different gas surface density proﬁles.

6.3 Asymmetries in NGC 1566

In Chapter 4 (published in Elagali et al., 2019), we present HI line observations of the spiral galaxy

NGC 1566 using 16 newly commissioned ASKAP antennas. We use ASKAPSOFT pipeline to ﬂag, calibrate and image these observations. Our HI line cube has a 3σ column density sensitivity of

19 2 1 NHI = 3.7 10 cm , measured over a 20 km s channel width, and an angular resolution × − − that equals 35 . The HI disc of NGC 1566 extends well beyond the observed optical disc ∼ 00 especially around the northern and western parts of NGC 1566. We derive the rotation curve of NGC 1566 and the dark matter content in this galaxy based on the Isothermal, NFW and Burkert dark matter halo proﬁles. However, we are unable to differentiate between these dark matter density proﬁles as the central region ( 2 kpc) of NGC 1566 is not resolved by our ASKAP observations. ∼

The HI morphology of NGC 1566 is asymmetric, with the south-eastern part of the HI disc sharply declining beyond 30 kpc from the centre, whereas the north-western edge is more ∼ extended and smoothly declines with radius, up to a radius of 50 kpc. We use this data to differen- ∼ tiate between the environmental processes that can cause asymmetries in galaxies, namely, ram pressure interactions with the IGM, galactic interactions and gas accretion from hosting/neighbouring

filaments. Based on the sensitivity limit of ASKAP observations, we detect no nearby companion galaxy or any HI clouds connected to, or in the nearby vicinity of, the HI disc of NGC 1566, which rules out the possibility of a past fly-by interaction or gas accretion. However, we find that ram pressure interactions with the IGM, assuming an IGM density of 5 10 5 cm 3, may be strong × − − enough to affect the HI disc of NGC 1566 and could be the reason for the asymmetries seen in the

HI morphology.

6.4 Interactions in the NGC 1566 group

In Chapter 5, we present a wide ﬁeld HI image of the NGC 1566 group (30 deg2) using ASKAP

20 2 7 1 observations. Down to the 5σ column density limit of 0.6 10 cm− ( 1.2 10 M beam− ), × ' × we detect a total of 24 HI sources in this ﬁeld, seven of which are new HI sources that have no

270 6.5. FUTURE PLANS

previous known redshift in the literature. Using the high sensitivity of the ASKAP observations, we are able to detect bridges in several interacting galaxies to a deeper level than known in previous studies, such as between the two triplets NGC 1533/IC 2038(2039) and NGC 1602(1595)/FS90077 as well as the galaxy pair IC 2058/PGC 75125. The speciﬁc angular momentum of the majority of the detected galaxies follows the analytic relation in Obreschkow et al. (2016). The HI sources that break the relation, which exhibit high speciﬁc angular momentum but have low gas mass fractions, are galaxies affected by their environment through either past tidal interactions or/and possible ram pressure stripping.

6.5 Future Plans

In this thesis, we discuss several questions regarding interactions in group environments and have opened many for future research. Below, we discuss some of the new avenues that will help improve our understanding of the topics discussed in the previous Chapters:

• Some observations indicate that ring galaxies are H2 deﬁcient even in the ring feature, where

the HI gas surface density is highest. To explain this puzzle, few observers hypothesise that the molecular hydrogen is destroyed by the continuous UV-ﬁeld from supernovae and OB stars in the ring (Higdon et al., 2015; Wong et al., 2017). This hypothesis is still an open question, especially because the subgrid physics modules in existing cosmological

simulations such as EAGLE and ILLUSTRIS-TNG do not have a detailed description of the cold phase of the ISM, and it is not possible to follow the formation and destruction of the molecular clouds. We can investigate this photodissociation process in smaller zoom-in

simulations such as the FIRE simulations (Hopkins et al., 2014, 2017). Alternatively, we

can create a zoom box around one of the high redshift EAGLE rings and attempt a more detailed simulations of the physics of the ISM and the different star formation prescriptions

in order to follow the formation and destruction of the H2 in ring galaxies. In addition to

the simulations, high resolution molecular hydrogen observations with ALMA will improve the current understanding of the cold gas phase of the ISM in ring galaxies and will be complementary to these new simulations.

• In Chapters 4 and 5, we investigate how the physical processes such as ram pressure stripping

and tidal harassments affect the HI gas in the NGC 1566 loose galaxy group. Thanks to the

271 CHAPTER 6. CONCLUSIONS AND SUMMARY

advent of ASKAP,HI observations of large areas on the sky are now possible and to a low

surface brightness sensitivity using much less observational time compared to the ASKAP early science commissioning phase. We intend to exploit these observations to probe different environments (groups and clusters) around galaxies and quantify the signiﬁcance/effects of gas accretion, galaxy-galaxy interactions and ram pressure stripping in large sample of galaxies. We also intend to use IFU spectroscopy and CO observations to explore the effects of the environment on the ionised and cold gas as well as to explore the star formation, the angular momentum, and misalignments between the ionised gas and stellar kinematics in galaxies. This will provide further insights to disentangle between the different environmental processes acting on group/cluster galaxies. For instance, ram-pressure removal of gas implies that the truncation of star formation is an outside-in process (Fossati et al., 2016; Kapferer et al., 2009) since the gas is preferentially removed in the outer parts of galaxies, which are less gravitationally bound (Elagali et al., 2019). On the other hand, the misalignments between the gas and stellar kinematics can be used to identify gas of external origin (Bryant et al., 2019), and the stellar and gas phase metallicity gradients are tracers of low metallicity external accretion and past merger history (Kobayashi, 2004). Hence,

combing HI observations and IFU spectroscopy will be key to disentangling the role of the nature, i.e. internal processes, and the nurture, i.e. the ensemble of environmental processes on the mass assembly and evolution of galaxies.

Another key component needed to answer these questions is the confrontation of theory and observations, given that the environment acts on galaxies via hydrodynamical and grav-

itational forces, hydrodynamical simulations (e.g., EAGLE, ILLUSTRIS-TNG) constitute a powerful tool to complete this study. Using simulations, we will have the leverage of tracing the evolution of galaxies and their hosts (groups) through cosmic time. This will allow us to answer questions such as: what is the timescale after which a galaxy is completely stripped of its gas since the time of in-fall in the group/cluster? How is this timescale dependant on the stellar mass and morphology of the stripped system and the halo mass of the host

group/cluster? We intend to also examine galaxy groups that are close to the cluster environments and those that are distant from peak density environments to answer the question: how the large-scale scale structure, e.g. within walls, ﬁlaments or voids can affect the growth

of galaxies via both tidal torques and/or gas accretion? Using simulations to study the different gas phases in galaxy groups/clusters will also allow us to test the physical treatment

272 6.5. FUTURE PLANS

of the gas in galaxy formation models. To sum up, we intend to explore the different gas

phases: atomic (via WALLABY & APERTIF), ionised (via MUSE, HECTOR & WEAVE) as well

as molecular hydrogen (via ALMA) and use this synergy to delve into the details of galaxy formation history and evolution. The measurables delivered by each facility along with the simulations will help us decode each galaxy’s genome and discover the strongest impacts through cosmic time on its formation and evolution.

• With the ASKAP observations presented in this thesis, we are unable to derive high-resolution (few hundred parsec scales) rotation curves for the observed galaxies and consequently fail to differentiate between different dark matter density proﬁles in their central regions. Such

high angular resolution (1000) observations will be achieved in the next coming years using baselines out to 6 km, which will be available for nearby selected galaxies. We plan to combine these data with other multiwavelength datasets, such as CO and Hα observations, to derive the rotation curves for nearby galaxies and study the central density proﬁle of dark matter halos to constrain the baryonic and dark matter mass distributions in nearby galaxies.

273 CHAPTER 6. CONCLUSIONS AND SUMMARY

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278 Appendix A

A Giant Galaxy in the Young Universe with a Striking Ring

A.1 Introduction

In this Appendix, I discuss part of the results and analysis of a journal article submitted to Na- ture astronomy (Yuan, Elagali et al.). I am the second author in this work in collaboration with Dr. Tiantian Yuan at Swinburne University and others. This paper reports the discovery of the most distant ring galaxy in the literature, at redshift z = 2.19 or at look-back time of 10.7 Gyr, using the the Cosmic Evolution Survey (COSMOS, Scoville et al., 2007) ﬁeld of the FourStar

Galaxy Evolution Survey (ZFOURGE, Straatman et al., 2016)1. My contribution to this paper is mainly in providing the theoretical background and the possible scenarios for the formation of this giant ring galaxy. The discovery of this ring galaxy (hereafter called R5519) was accidental, and came about while Dr. Yuan was visually searching for spiral galaxies in ZFOURGE images in the redshift range 1.8 z 2.5. R5519 is one of the largest galaxies among the 4000 . . ∼ galaxies inspected by Dr. Yuan, with a distinct ring morphology and a large diffuse disc. Figure A.1 shows a collage of images of this ring galaxy and its local environment made by Dr. Yuan using the Hubble Space Telescope (HST) and the ZFOURGE different wavebands. In the following sections, I brieﬂy discuss the results of this paper, with more focus on the theoretical aspects which was my contribution to this paper. I ﬁrst present the general properties of the highest redshift z 2.23 EAGLE collisional ring galaxy. Then, I compare between this simulated galaxy and ∼ R5519, but also investigate other possible scenarios for the formation of R5519 at such high redshift.

1http://z-fourge.strw.leidenuniv.nl

279 APPENDIX A. A GIANT GALAXY IN THE YOUNG UNIVERSE WITH A STRIKING RING A.2 The Observations of R5519

The ring galaxy R5519 was discovered during the systematic search for z & 2 spiral galaxies in the

Cosmic Evolution Survey (COSMOS) ﬁeld of the FourStar Galaxy Evolution survey (ZFOURGE,

Straatman et al., 2016). Dr. Yuan used ZFOURGE catalogue images to identify spiral structures in galaxies falling within the photometric redshift range of 1.8 . zp . 2.5. Because of the surface brightness dimming and smaller sizes of galaxies at z > 1, the visual spiral feature identification was restricted to galaxies with illuminated pixels larger than a radius of 0. 5 (> 4 kpc at z 2) in 00 ∼ the Hubble Space Telescope (HST) images. Dr. Yuan’s visual inspection simultaneously identified ring galaxies as well as other morphologically distinct objects such as mergers and gravitationally lensed galaxies. R5519 was flagged as one of the largest galaxies among the 4000 galaxies ∼ inspected, with a clear ring structure as well as a large diffuse disk (Fig. A.1).

The spectroscopic redshift of R5519 was confirmed to be at zs = 2.192 0.001 based on our ± Keck/MOSFIRE2 near-infrared (NIR) spectroscopy and Keck/OSIRIS3 adaptive-optics aided NIR integral field spectroscopy. A joint analysis of the MOSFIRE and OSIRIS spectroscopic data as well as the ZFOURGE Hα images, shows that the Hα kinematics of R5519 are consistent with a tilted rotating and expanding/contracting circular ring model. Taking the inclination angle (i = 29 5) ± and the position angle (PA = 28 10) from an ellipse fit to the ring structure as inputs to the ± kinematic model, the inferred rotational velocity at the fixed ring radius (Rring = 5.1 0.4 kpc) is ± V = 90 75 km/s and the radial expansion/contraction velocity is V = 226 90 km/s. The rot ± rad ± error bars in velocities represent uncertainties from observational measurements. The systematic errors caused by uncertain ranges of PA and i are of the same order of magnitude. This is the first time that an expansion/contraction velocity has been derived for a collisional ring galaxy (CRG) beyond the local universe.

R5519 resides in a small galaxy group environment, reminiscent of the loose groups in which local CRGs such as the Cartwheel galaxy are found (Higdon, 1995; Appleton and Struck-Marcell, 1996; Madore et al., 2009). A companion galaxy (ID 5593, hereafter G5593) is conﬁrmed at 30 ∼ kpc projected distance from R5519 with a 3D-HST survey (Momcheva et al., 2016) grism redshift

+0.005 of zgrism = 2.184 0.006. An additional group candidate (ID 5475, hereafter G5475) is found − ∼ 40 kpc away (Fig. A.1), with a photometric redshift of zp = 2.1 0.1 (Straatman et al., 2016). If ± +65 G5593 is the intruder of R5519, then the inferred timescale after collision is τc > 39 15 Myr; this − 2https://www2.keck.hawaii.edu/inst/mosfire/home.html 3https://www2.keck.hawaii.edu/inst/osiris/

280 A.3. THE MORPHOLOGY AND KINEMATICS OF THE Z = 2.23 EAGLE RING GALAXY is a lower limit due to unknown projection effects. R5519 has a total UV+IR star formation rate

1 (SFR) of 80.0 0.2 M yr− and stellar mass of log(M /M ) = 10.78 0.03. In comparison, ± ∗ ± 1 G5593 has a total SFR of 123 2 M yr− and log(M /M ) = 10.40 0.04. The furthest group ± ∗ ± candidate (G5475) within 50 kpc projected distance from the ring is a compact quiescent galaxy. The morphology of G5593 shows double nuclei and a tidal tail (Fig. A.1), suggestive of an ongoing merger of its own.

One of the most striking features of R5519 is the extended stellar light outside of the ring in multiple wavelengths (Fig. A.1). Comparing with other z 2 galaxies in the 3D-HST cata- ∼ logue (van der Wel et al., 2014), the R 4 (11.8 0.3 kpc from HST WFC3 images) of R5519 80 ± is 2.4σ larger than the median size (5.4 kpc) of all late-type galaxies (log(M /M ) > 9.5) and ∗ 1.5σ larger than the mean value (7.1 kpc) of the most massive (log(M /M ) 10.6) late-type ∗ ≥ galaxies at z 2. A morphological inspection on the other unusually large (1σ above the mean) ∼ and massive late-type galaxies in the COSMOS ﬁeld reveals that most of them are merging galaxies.

A.3 The Morphology and Kinematics of the z = 2.23 EAGLE Ring Galaxy

Elagali et al.(2018) present a noval identification method of ring galaxies in the EAGLE simulations suite mock optical images based on the radial flux distribution in these galaxies. The highest- redshift collisional ring galaxy in the EAGLE simulation identified by Elagali et al.(2018) is at a redshift of z = 2.23 and is visualised in Figure A.2. Only one ring galaxy was found in the entire (100 Mpc)3 volume at this redshift, which indicates that the formation of ring galaxies is highly hampered at higher redshifts. At lower redshifts (z < 2), we usually identify an order of 10 ring galaxies in each snapshot (Elagali et al., 2018). We later discuss the possible physical reasons for this difference in frequency. The three panels in Figure A.2 show the face-on u-band mock image

(left), the HI (middle) and H2 (left) column density maps of this galaxy. The HI and the H2 column densities were calculated using the prescriptions in Rahmati et al.(2013) for the neutral hydrogen fraction in gas particles and Gnedin and Kravtsov(2011) for the H I to H2 transition, as described in detail in Lagos et al.(2015). The colours in the column density images are coded according to the bar on-top of each panel and are in units of atoms per cm2. The u-band image is 60 kpc on a side while the HI and H2 column density images are 100 kpc on a side to better reﬂect the

4 R80 is the radius within which 80% of the total luminosity is included.

281 APPENDIX A. A GIANT GALAXY IN THE YOUNG UNIVERSE WITH A STRIKING RING

Figure A.1: A montage of images for the ring galaxy R5519 and the neighbouring environment adopted from the Nature paper of Yuan, Elagali et al. Left: RBG composite colour image centred on R5519, in which R represents the HST F606W band, G the F814W band and B the WFC3 band. Numbers in this image denote the ZFOURGE catalog ID. Middle: a second RBG composite colour image, R represents the HST F125W band, G the F160W band and B the Magellan Kstot band. The photometric redshifts of each object derived using ZFOURGE are labelled in white, whereas in yellow are the spectroscopic redshifts. Right: a third RBG composite colour image, in which R represents the FUV band, G the NUV band and B the optical image. The images are generated by stacking HST and ZFOURGE catalogue images that correspond to rest-frame FUV, NUV and optical wavelengths. more extended nature of the gas with respect to stars. The morphology of the atomic and molecular hydrogen distribution in this simulated galaxy is similar to the stellar distribution seen in the u-band image, except for the nucleus where there is a drop (“holes”) in the column density. This is similar to the morphology of observed local collisional ring galaxies, for example the atomic and molecular hydrogen distributions in the Cartwheel are mainly concentrated in the ring structure and the inner regions are dominated by lower column densities (Higdon, 1996; Higdon et al., 2012).

To analyse the kinematics of this z = 2.23 simulated collisional ring galaxy, we use four con- secutive EAGLE “snipshots”, which have better time resolution than the snapshots but a more limited set of output information. For our purpose, all the relevant information exists in the snipshots. The time resolution of the former is of the order 60 Myr (Crain et al., 2017). We define time t = 0 Myr as the time when the ring feature was identified in the snapshot at z = 2.23, this is an arbitrarily defined time and is not representative of the exact time when the collision occurred. In Figure A.3, we show four columns from left to right and denote the time of the four snipshots (relative to t = 0 defined above) on the top. The first row shows a 500 500 kpc view of the simulation Z Y axes × − and illustrates the large scale environment around the ring host. This view is centred exactly on the ring dynamical centre, the red-brown dots show the host galaxy’s stars, blue dots show the satellites’ stars, and the grey dots show the gas particles. The mass of the halo hosting these galaxies

12 is M200 = 1.2 10 M , which represents the mass enclosed within a radius that contains an ×

282 A.3. THE MORPHOLOGY AND KINEMATICS OF THE Z = 2.23 EAGLE RING GALAXY

Figure A.2: Visualisation of the z = 2.23 ring galaxy in the EAGLE RefL100N1504 simulation. The three panels corresponds to the face-on u-band image (left), the HI (middle) and H2 (left) column density maps of this galaxy. The colours in the column density images are coded according 2 to the colour bar at the top row and are in units of cm− . The u-band image is 60 kpc on a side while the HI and H2 column density maps are 100 kpc on a side. average density of 200 times the critical density or the mean matter density of the universe. This simulated collisional ring galaxy has a stellar mass of M = 1.4 1010 M and is the central ∗ × galaxy within this parent halo. There are three smaller satellite galaxies associated with the same

9 halo, their stellar masses are M & 10 M . The stellar mass ratio between the intruder satellite ∗ galaxy and the ring host is 0.38. The middle row shows a 50 50 kpc zoom-in view of the ring ∼ × host along the Z Y axes, the red-brown dots show the ring young stars with ages < 100 Myr, and − grey dots show the gas particles. There is a clear ring structure, made predominantly of young stars, in the three first snipshots, which disappears at t = 180 Myr. The bottom row shows the stars (solid line) and gas particles (dashed line) surface mass density (excluding those from the nucleus) as a function of radius in the ring host. The location of the ring feature is defined as the peak of the star particle surface mass density and the positions of these particles are shown as a function of radius in subsequent snipshots. We observe that the radius of the ring expands from 5, 6.5 to 15 kpc from t = 0, 60 to 120 Myr. We tag the gas and star particles that are identified to be part of the ring at t = 0 Myr and follow their positions in subsequent times at t = 60, 120, 180 Myr. The particles that make up each ring at these three time-frames are mostly composed of new particles that do not come from the ring formed at t = 0 Myr; this is consistent with the ring being a wave pattern.

In Figure A.4, we show the circular rotation and radial velocity of the gas particles at different radii in the z = 2.23 collisional ring galaxy for the four time sequence used in Figure A.3. We only consider the particles that lie within a radius of r = 30 kpc from the centre of the galaxy and assume a scale height of h = 5 kpc for the disc. In this case, the radial velocity is calculated for gas particles

283 APPENDIX A. A GIANT GALAXY IN THE YOUNG UNIVERSE WITH A STRIKING RING in the plane of the disc, where the plane is determined as that perpendicular to the total stellar spin vector of the galaxy. The radial velocity therefore describes the instantaneous radially outwards (positive) or inwards (negative) motion of the gas particles. Note that the instantaneous radial velocity does not reﬂect the motion of the density wave. The large variations in radial velocities at all radii from t = 0 Myr to t = 120 Myr is likely triggered by the interactions with satellite galaxy(ies) as shown in Figure A.3. The ring feature is disrupted at a later snipshot (t = 180 Myr) due to a second passage of a satellite galaxy (see Figure A.3). Hence, this ring feature has a lifespan of less than 180 Myr. This is relatively low compared to the lifespan of collisional ring galaxies in an idealised or isolated interaction simulations where the ring can expand for times > 500 Myr to its maximum radius and to very low surface brightness levels (Mapelli et al., 2008). In such isolated simulations, the intruder interacts once, not multiple times as in this EAGLE ring case, with the target galaxy and does not completely merger with the target. It is also interesting to highlight the very gas-rich nature of this system, as gas particles are seen across the image. This is also expected to contribute to the exact evolution of the ring galaxies, as continuing gas accretion may hamper the expansion of the density wave induced by the intruder.

A.4 Similarities between the z = 2.23 EAGLE Ring and the ZFOURGE Ring

One possible scenario for the formation of R5519 is via drop-through interactions with neighbouring satellite galaxy(ies). In this section, we discuss the possibility of this scenario using the z = 2.23

EAGLE ring galaxy as an example. However, we note that this EAGLE collisional ring galaxy is not meant to be used as a direct comparison or a replication of R5519. The total star formation rate of R5519, based on the combined rest-frame infrared luminosity and rest-frame UV luminosity, is

1 80.0 0.2 M yr− and possesses a stellar mass of log(M /M ) = 10.78 0.03. In comparison, ± ∗ ± the EAGLE ring galaxy is 0.23 times less massive than R5519, and has a star formation rate of ∼ 1 3.8 M yr− (more than 10 times lower than R5519). One of the most striking features of R5519 is its extended stellar light outside of the ring seen in multiple wavelengths (Figure A.1). The radial light distribution is steeper at a larger radius than the exponential proﬁle of a typical isolated disc galaxy, consistent with a post-collisional disc proﬁle of a collisional ring galaxy (Gerber et al., 1996).

Despite these differences, we see that the EAGLE ring’s dynamically active environment at z 2 causes stellar and gas particles to be tugged to larger radii, resulting in a somewhat large ∼

284 A.4. SIMILARITIES BETWEEN THE Z = 2.23 EAGLE RING AND THE ZFOURGE RING

Figure A.3: The time evolution of a collisional ring galaxy at z = 2.23 from the EAGLE cosmological hydrodynamical simulations (Elagali et al., 2018). The four columns from left to right denote four time snipshots (Crain et al., 2017) with respect to an arbitrarily defined ring formation time (t = 0 Myr). Snipshots are defined in the EAGLE simulations as having finer time resolution than snapshots. The top row shows a 500 500 kpc view, through the simulation Z Y axes, of the large scale environment in 2D centred on× the ring host; red-brown = host galaxy stars,− blue = satellite halo stars, and grey=gas particles. The middle row shows a 50 50kpc zoom-in view of the ring × host in 2D, red-brown = young stars with ages < 100 Myr, and grey=gas. The bottom row shows the star and gas particle surface mass density (excluding those from the nucleus) at each radius. The location of the ring feature is defined as the peak of the star particle surface mass density. The green and purple lines show the location of the ring’s star and gas particles identified at t = 0 Myr respectively, and the positions of these particles with radius in the subsequent snipshots. In a later time snipshot at t = 240 Myr (not shown here), the nucleus of the collisional ring galaxy is found to be broken into two parts after the major merger at t = 180 Myr.

285 APPENDIX A. A GIANT GALAXY IN THE YOUNG UNIVERSE WITH A STRIKING RING stellar disc that is not usually seen in nearby collisional ring galaxies. The simulated z = 2.23 collisional ring galaxy at t = 120 Myr looks morphologically similar to our observed R5519, lending an explanation for the diffuse large disc outside of the ring of R5519. The snipshot at t = 180 Myr (Figure A.3) may look like it is beginning to form a second ring as predicted by the analytic collisional ring galaxy theory (Struck, 2010; Smith et al., 2012). However, a later snipshot at t = 240 Myr (not shown in Figure A.3) reveals that the collisional ring galaxy is completed destroyed (including its nucleus) by the major merger at t = 180 Myr before a second ring can be deﬁnitively identiﬁed. This highlights the “hostile” environment typical of high redshift galaxies that can lead to the demise of collisional ring galaxies. R5519, on the other hand, resides in a similar galaxy group environment reminiscent of the loose groups in which local CRGs such as the Cartwheel galaxy are found (Higdon, 1995; Appleton and Struck-Marcell, 1996; Madore et al.,

2009). The closest companion galaxy to R5519 lies at 30 kpc projected distance. ∼

Further, the ground-based Hα narrow-band image from the ZFOURGE catalogue shows that the Hα kinematics of R5519 are consistent with a tilted rotating and expanding/contracting ring model. The inferred rotational velocity at the fixed ring radius (Rring = 5.0 0.4 kpc) is V , ± rot 5kpc = 90 75 km/s and the radial expansion/contraction velocity is V , = 226 90 km/s. The ± rad 5kpc ± error bars in these velocities represent the uncertainties from observational measurements such as the position angle and the inclination of this galaxy. This radial velocity component in the Hα is similar to theoretical models of collisional ring galaxies, including our paper on EAGLE ring galaxies (Elagali et al., 2018). These classical models predict that the ring is a radially expanding density wave. However, kinematic confirmation of the density wave is observationally challenging even for nearby galaxies. The current kinematic data alone is insufficient to rule out a regular merger. A collisional ring galaxy is a special type of a major merger with a drop-through collision.

Only a few percent ( 2%) of interacting galaxies are found to be collisional ring galaxies in the ∼ local universe (Madore et al., 2009). The observational difference between collisional ring galaxies and other common (regular) mergers is that collisional ring galaxies are geometrically simpler, characterised by a symmetric ring and a relatively well-preserved disc whereas regular mergers are morphologically messier and retain less information about the pre-disrupted disc (Appleton and Struck-Marcell, 1996). The prominent ring and the profile of the stellar light of R5519 is consistent with the definition of a collisional ring galaxy. Note that the chance of finding regular mergers projected as a symmetric ring is lower than finding a collisional ring galaxy (Romano et al., 2008). Moreover, in a regular merger or tidal interactions scenario, the star formation in the central regions of galaxies contributes significantly to the total SFR whereas the star formation of a collisional ring

286 A.4. SIMILARITIES BETWEEN THE Z = 2.23 EAGLE RING AND THE ZFOURGE RING

Figure A.4: The circular and radial velocity of the gas particles in the z = 2.23 EAGLE collisional ring galaxy. We show four time snipshots with respect to an arbitrarily deﬁned ring formation time (t = 0 Myr). For each time snipshot, the circular and radial velocity are shown in the upper and lower panels, respectively. This collisional ring galaxy is expanding from t = 0 Myr to 120 Myr, but the ring feature is disrupted due to a second passage of a satellite galaxy at time 180 Myr. During these four time snipshots the radius of this collisional ring galaxy ring (deﬁned as the peak of the 1D surface mass density of the stars; see Figure A.3) also changes from 5, 6.5, 15, to 7.5 kpc, respectively.

287 APPENDIX A. A GIANT GALAXY IN THE YOUNG UNIVERSE WITH A STRIKING RING galaxy usually occurs on the ring (Arabsalmani et al., 2019). We do not detect any signiﬁcant star formation activity in the central regions of the ring in R5519. We therefore think a collisional ring galaxy origin is a more reasonable explanation for R5519 than a regular merger.

A.5 An Alternative Secular Origin for the ZFOURGE Ring

A large fraction of ring galaxies in the local Universe have not formed through interactions and are the result of secular galaxy evolution, those are called the resonance ring galaxies (de Vaucouleurs, 1959; Herrera-Endoqui et al., 2015; Buta, 2017). These rings are preferentially found in barred galaxies, with the ring encircling the bar and forming the familiar θ-shape. These type of rings have formed by gas accumulation at Lindblad resonances under the continuous inﬂuence of gravity torques from the bars (see Figure A.5). Even though in the local Universe secular rings are more abundant than collisional ring galaxies (Theys and Spiegel, 1976; Buta and Combes, 1996), the

Hubble spiral sequence is not yet fully established at z > 2 (Conselice, 2014). Most galaxies at z > 2 are clumpy, irregular and have small/no bulge component (Bournaud and Elmegreen, 2009). If R5519 is a secular ring at z = 2.19, we ﬁnd two aspects that are challenging to explain in our current framework of disc galaxy formation and the origin of the Hubble sequence.

• R5519 have a large disc of diffuse light outside of the ring with a size of R80 = 11.6 kpc,

where R80 is the radius that contains 80 per cent of the galaxy’s total luminosity. The average

outer disc size of the most massive isolated disc galaxy at z 2 in the ZFOURGE sample is ∼ R80 = 6.4 kpc, a factor of 2 smaller than R5519. In fact, R5519’s extended disc is almost 1.5 times larger than the Milky Way at z = 0 (van der Kruit and Freeman, 2011). Such a large extended disc of a secular origin has not been reported in observations or simulations at

z > 2. The reason for the lack of Milky Way size large discs at z > 2 is that it takes time to build up angular momentum to sustain large discs. Theoretical work on the development of torques and angular momentum in the hierachical Universe show that the angular momentum of halos increase with cosmic time (Catelan and Theuns, 1996b,a), and at earlier times halos are still growing in angular momentum and are unable to sustain large discs (Zavala et al., 2016). Recent cosmological hydrodynamical simulations show that large discs do not form

until z 1 (Lagos et al., 2017). Moreover, the surface brightness distribution of the diffuse ≤ light of R5519 deviates from exponential disc, and is more consistent with a disc disrupted by a drop-through merger (Gerber et al., 1996; Romano et al., 2008).

• The existence of a resonance ring in R5519 implies the presence of a large bar within the nucleus of R5519. This disagrees with the current observations, because the fraction of spiral

288 A.5. AN ALTERNATIVE SECULAR ORIGIN FOR THE ZFOURGE RING

Figure A.5: Left: an example of a collsional ring galaxy in the local Universe (the Cartwheel) taken from the Hubble Space Telescope gallery. Right: a resonance ring galaxies NGC 5335 adapted from Buta(2017).

and barred galaxies decreases with redshift and becomes extremely difﬁcult to ﬁnd at z > 2 (Sheth et al., 2012; Law et al., 2012; Kraljic et al., 2012; Elmegreen and Elmegreen, 2014;

Conselice, 2014; Yuan et al., 2017). It is true that the lack of z > 2 bars in observations could be caused by the limited spatial resolution and the depth of high-z galaxy surveys. Bars are associated with the spiral discs of the Hubble sequence, however spiral arms are easier to observe than the bar structure. Many high-redshift observations have pointed out

that spiral galaxies themselves do not fully emerge until z < 2 (Elmegreen and Elmegreen, 2014; Conselice, 2014). For example, there are less than three galaxies observed with spi-

ral arms at z > 2 (Law et al., 2012; Yuan et al., 2017) and no conclusive barred spirals have been observationally conﬁrmed at z > 2 (Kraljic et al., 2012; Vincenzo et al., 2019). Cosmological hydrodynamical simulations show that the dynamically active environment at

z > 1 are hostile to the formation of bars (Sheth et al., 2012; Kraljic et al., 2012). Vincenzo et al.(2019), for example, investigate the formation of a barred spiral at z > 2 in zoom-in hydrodynamical simulations and find that these galaxies are predominantly in the field and have no companions. This finding is consistent with previous studies pointing out that bar formation at high redshift is difficult because any merger or violent gas accretion event can destroy the bar (Sheth et al., 2012; Kraljic et al., 2012). This picture is also consistent with the overall evolution of the Hubble sequence with respect to the evolution of the environment (Cen, 2014a,b). If we interpret R5519 as one of the outliers that formed a bar in the early Universe, then the active group environment it resides in is inconsistent with current theories of bar formation. A collisional ring origin, on the other hand, can naturally account for the

289 APPENDIX A. A GIANT GALAXY IN THE YOUNG UNIVERSE WITH A STRIKING RING

large diffuse disc and other observed properties of R5519. Note that non-circular orbits and projection effects can mimic the impression of an expanding/contracting component in the velocity ﬁeld, however the projected velocity of secular rings (Kormendy and Kennicutt, 2004) is usually much lower than the expansion/contracting velocity of collisional ring

galaxies. The large V , = 226 90 km/s we derive for R5519 is inconsistent with a exp 5kpc ± secular ring and more consistent with a collisional ring, e.g. V , . = 225 8 km/s for exp 5 8kpc ± the Apr 147 (Fogarty et al., 2011).

A.6 Conclusion

In this short Chapter, we discuss the discovery of the most distant ring galaxy, R5519, reported in the literature to-date and provide possible scenarios for the formation of this galaxy at the redshift of z = 2.19. We provide evidence in support of a collisional origin for R5519, and include a comparison with a simulated collisional ring galaxy at similar distance from the EAGLE simulation. Detailed studies of distant collisional ring galaxies, such as R5519, offer a pathway for refining our knowledge of star formation and dynamics in young galaxies, as well as constraining the cosmic evolution of large spiral discs and galaxy groups. This is because the conditions required to form ring galaxies are unique. For a galaxy to form and maintain the ring feature for timescales larger than 100 Myr, two additional conditions must be satisfied besides the drop-through collision. The first condition is related to the local environment in which the ring feature is formed. Local collisional ring galaxies are often found in loose groups instead of compact groups or galaxy clusters (Madore et al., 2009). Even though interactions rates are higher in dense environments (e.g, clusters), the high velocity dispersions between cluster members and fly-by interactions can destroy the ring features. Similarly, the galaxy environment at high redshift is very active because of the frequent interactions with satellites and high gas accretion rates (Cen, 2014a; Rodriguez-Gomez et al., 2015). Our z = 2.23 EAGLE ring is short-lived in this dynamically chaotic high-redshift environment. Finding and measuring the number density of objects like R5519 across cosmic time can provide an alternative constraint on the assembly and evolution of loose galaxy groups.

The second additional condition is related to the properties of the pre-collisional “ring feature” host. In the analytic caustic theory of collisional ring galaxy formation, the pre-collisional ring host is implicitly assumed to have a dynamically “thin" (low velocity dispersion) disc before the collision so that the expanding density wave can propagate through the disc medium. This assumption is

290 A.6. CONCLUSION valid for most spiral galaxies nearby but certainly not for high redshift galaxies which tend to have large velocity dispersions (Glazebrook, 2013). A vertically thick (high velocity dispersion) disc can dampen the propagation of density waves, similar to the suppression of density-wave spiral arms in high-redshift galaxies (Yuan et al., 2017). This implies that if the collision occurs between two clumpy, dispersion-dominated galaxies, it is hard for the ring feature to form and remain visible for as a long timescale as nearby collisional ring galaxies. At high redshifts (z 2), most galaxies are ≥ clumpy and have high velocity dispersions in the vertical disc (Glazebrook, 2013; Wisnioski et al., 2015). It is sensible to speculate that the lack of spiral galaxies and thin discs (Law et al., 2012;

Cen, 2014a; Yuan et al., 2017; Freeman and Bland-Hawthorn, 2002; Helmi et al., 2018) at z 2 ≥ will lead to a decreased chance of observing large collisional ring galaxies. Our predictions can be tested as more detailed observations and simulations of high-redshift collisional ring galaxies like R5519 are conducted.

An important current limitation for the theoretical exploration of high redshift ring galaxies is the limited cosmological volume of EAGLE and as a result the small number statistics. We note that large cosmological volumes have been produced with full hydrodynamics, but those usually have signiﬁcantly poorer resolution, preventing the study of the internal structure of galaxies.

In the future, simulations such as EAGLE-XL will give us access to 27 times more volume and comparable spatial resolution. This will be essential to understand the diversity of formation paths and properties (e.g. star formation rates, stellar masses, gas masses) of ring galaxies at the peak of star formation in the Universe.

291 APPENDIX A. A GIANT GALAXY IN THE YOUNG UNIVERSE WITH A STRIKING RING

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