The ram pressure stripped radio tails of galaxies in the Coma cluster
Hao Chen, Ming Sun, Masafumi Yagi, Hector Bravo-Alfaro, Elias Brinks, Jeffrey Kenney, Francoise Combes, Suresh Sivanandam, Pavel Jachym, Matteo Fossati, Giuseppe Gavazzi, Alessandro Boselli, Paul Nulsen, Craig Sarazin, Chong Ge, Michitoshi Yoshida, Elke Roediger
MMNRAS , 1–20 (2020) Preprint 25 June 2020 Compiled using MNRAS L A TEX style file v3.0
The ram pressure stripped radio tails of galaxies in theComa cluster
Hao Chen, , (cid:63) Ming Sun, † Masafumi Yagi, , Hector Bravo-Alfaro, Elias Brinks, Jeffrey Kenney, Francoise Combes, Suresh Sivanandam, , Pavel Jachym, Matteo Fossati, , Giuseppe Gavazzi, Alessandro Boselli, Paul Nulsen, , Craig Sarazin, Chong Ge, Michitoshi Yoshida, , and Elke Roediger, Department of Physics & Astronomy, University of Alabama in Huntsville, 301 Sparkman Drive, Huntsville, AL 35899, USA Department of Astronomy, University of Cape Town, Private Bag X3, Rondebosch 7701, South Africa Optical and Infrared Astronomy Division, National Astronomical Observatory of Japan, Mitaka, Tokyo, 181-8588, Japan Graduate School of Science and Engineering, Hosei University, 3-7-2, Kajinocho, Koganei, Tokyo, 184-8584, Japan Departamento de Astronom´ıa, Universidad de Guanajuato, Apdo. Postal 144, Guanajuato 36000, Mexico Centre for Astrophysics Research, University of Hertfordshire, College Lane, Hatfield AL10 9AB, UK Yale University Astronomy Department, P.O. Box 208101, New Haven, CT 06520-8101, USA Observatoire de Paris, LERMA, College de France, CNRS, PSL Univ, Sorbonne University, UPMC F-75014 Paris, France Department of Astronomy and Astrophysics, University of Toronto, 50 St. George St, Toronto, ON, Canada Dunlap Institute of Astronomy and Astrophysics, University of Toronto, 50 St. George St, Toronto, ON, Canada Astronomical Institute, Academy of Sciences of the Czech Republic, Boˇcn´ı II 1401, 141 31 Prague 4, Czech Republic Dipartimento di Fisica G. Occhialini, Universit¨a degli Studi di Milano Bicocca, Piazza della Scienza 3, I-20126 Milano, Italy Institute for Computational Cosmology and Center for Extragalactic Astronomy, Durham University, South Road, Durham DH1 3LE, UK Universit`a degli Studi di Milano-Bicocca, Piazza della Scienza 3, 20126 Milano, Italy Aix-Marseille Univ., CNRS, CNES, LAM, Marseille, France Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA ICRAR, University of Western Australia, 35 Stirling Hwy, Crawley, WA 6009, Australia Department of Astronomy, University of Virginia, Charlottesville, VA 22904, USA Subaru Telescope, National Astronomical Observatory of Japan, 650 North A’ohoku Place, Hilo, HI 96720 USA Hiroshima Astrophysical Science Center, Hiroshima University, Higashi-Hiroshima, Hiroshima 739-8526, Japan Milne Centre for Astrophysics, Department of Physics & Mathematics, University of Hull, Hull, HU6 7RX, UK
Accepted XXX. Received YYY; in original form ZZZ
ABSTRACT
Previous studies have revealed a population of galaxies in galaxy clusters with rampressure stripped (RPS) tails of gas and embedded young stars. We observed 1.4 GHzcontinuum and HI emission with the Very Large Array in its B–configuration in twofields of the Coma cluster to study the radio properties of RPS galaxies. The bestcontinuum sensitivities in the two fields are 6 and 8 µ Jy per 4 (cid:48)(cid:48) beam respectively,which are 4 and 3 times deeper than those previously published. Radio continuumtails are found in 10 (8 are new) out of 20 RPS galaxies, unambiguously revealing thepresence of relativistic electrons and magnetic fields in the stripped tails. Our resultsalso hint that the tail has a steeper spectrum than the galaxy. The 1.4 GHz continuumin the tails is enhanced relative to their H α emission by a factor of ∼ ∼ Key words: galaxies: clusters: individual (Coma) – galaxies: interactions – galaxies:ISM – radio continuum: galaxies (cid:63)
E-mail: [email protected] † E-mail: [email protected]
Ram pressure stripped (RPS) galaxies are characterized bygas being stripped from the affected galaxy by the intraclus- © a r X i v : . [ a s t r o - ph . GA ] J un H. Chen et al.
Right Ascension (J2000) D ec li n a t i o n ( J ) Figure 1.
XMM-Newton (cid:48) and 15.88 (cid:48) in radius) of the two VLA fields. ter medium (ICM, e.g., Gunn & Gott 1972; Nulsen 1982).Star formation (SF) can be triggered by ram pressure atthe early interaction stage by compression of interstellarmedium (ISM), as shown in observations and simulations(e.g., Koopmann & Kenney 2004; Crowl et al. 2006). Then,as the cold ISM is depleted, the galactic SF will be quenched(e.g., Quilis et al. 2000; Boselli et al. 2016b). Thus, rampressure stripping is an important process affecting galaxyevolution in rich environments like galaxy groups and clus-ters. The evolution of the stripped ISM is a significant areaof research. The mixing of the stripped cold ISM with thehot ICM will produce a multi-phase gas (e.g., Sun et al.2007; Ferland et al. 2009; J´achym et al. 2019). Some of the stripped ISM can turn into stars in the galactic halo and theintracluster space (e.g., Cortese et al. 2007; Sun et al. 2007;Owers et al. 2012; Ebeling et al. 2014; Cramer et al. 2019),especially in the high ICM-pressure environment (e.g., Sunet al. 2010). Thus, stripped tails emerge as ideal targets tostudy this multi-phase medium and SF conditions in an ex-treme environment.Ram pressure stripped tails are observed in X-rays (e.g.,Sun et al. 2006, 2010; Zhang et al. 2013), far-ultraviolet(FUV, e.g. Boissier et al. 2012), H α (e.g., Gavazzi et al.2001, 2017, 2018; Sun et al. 2007; Yagi et al. 2007, 2017;Yoshida et al. 2008, 2012; Smith et al. 2010; Yagi et al.2010, 2013; Fossati et al. 2012, 2016, 2018; Fumagalli et al. MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies (e.g., Sivanandam et al. 2010, 2014), CO (e.g., J´achymet al. 2013, 2014, 2017; Scott et al. 2013, 2015; Verdugo et al.2015; Moretti et al. 2018) and HI (e.g., Kenney et al. 2004,2014; Oosterloo & van Gorkom 2005; Chung et al. 2007,2009; Scott et al. 2010, 2012; Abramson & Kenney 2014;Ramatsoku et al. 2019; Serra et al. 2019; Deb et al. 2020).Extensive simulations (e.g., Quilis et al. 2000; Roediger &Br¨uggen 2008; Ruszkowski 2012) show that stripping hasa significant impact on galaxy evolution (e.g., disk trunca-tion, the formation of flocculent arms, the transformation ofdwarf galaxies). SF in the stripped tail has also been seenin simulations (e.g., Kapferer et al. 2009; Tonnesen & Bryan2010, 2012; Roediger et al. 2014).A complementary tool for studying stripped tails is theradio continuum emission. At 1.4 GHz, radio continuumemission is dominated by synchrotron radiation which isemitted by the relativistic electrons moving within a mag-netic field. Like the colder, denser (traced by HI) and hotter,more diffuse (traced by H α and X-ray) gas, the plasma con-taining relativistic electrons and magnetic fields is strippedby ram pressure (e.g., Gavazzi & Jaffe 1987). Murphy et al.(2009) identified radio-deficit regions along the outer edge of6 Virgo Cluster galaxies, revealing that relativistic electronsand magnetic fields on their leading edges had been removedby ram pressure. Furthermore, the stripped relativistic elec-trons can be re-accelerated (rejuvenated) in the tail (Pinzkeet al. 2013), either by turbulence and ICM shocks (Kang &Ryu 2011), or by new SNe. At the same time, local core-collapse supernovae can contribute to the relativistic elec-trons since H II regions (tracing massive stars) have beenfound in RPS tails (e.g. Sun et al. 2007; Yagi et al. 2010).However, fewer tails have been detected in radio continuumthan in H α and X-ray. For example, there are 17 strippedtails detected in H α in the Coma cluster (Gavazzi et al.2018), but only 2 of them have thus far been detected inradio continuum (Miller et al. 2009). Currently most of thestripped tails seen in late-type galaxies in radio continuumare short and are detected in nearby clusters, for exampleNGC 4522 (Vollmer et al. 2004), NGC 4402 (Crowl et al.2005) and others in the Virgo cluster. A few long RPS tails,such as CGCG 097-073, CGCG 097-079 and UGC 6697 havebeen reported in Abell 1367 (Gavazzi 1978; Gavazzi & Jaffe1985, 1987; Gavazzi et al. 1995).Why do RPS tails tend not to be detected in the radiocontinuum? How common are radio continuum tails behindRPS galaxies? How does the radio continuum emission intails correlate with emission in other bands? Are the ob-served radio continuum tails mainly related to star forma-tion in the tails or due to relativistic electrons which werestripped from the galaxy by ram pressure? To address thesequestions deep radio continuum data are needed, somethingwhich has become feasible with the new-generation wide-band correlators deployed on existing radio telescopes.As ram pressure stripping and SF activity in the tailsare believed to be more prominent in high-pressure environ-ments than in low-pressure environments (e.g., Sun et al.2010; Tonnesen et al. 2011; Poggianti et al. 2016, 2017), theComa cluster, as the most massive cluster at z < . , is anideal target for these studies. Coma has the richest opticaldata among nearby massive clusters, already with a sampleof over 20 late-type galaxies with one-sided SF or ionized gas tails (Smith et al. 2010; Yagi et al. 2010; Kenney et al. 2015;Gavazzi et al. 2018). There has been an increasing effort toobtain multi-wavelength observations of these galaxies. RPStails in bands other than H α have been detected. A spectac-ular example is D100, with a narrow tail observed in X-rayswith Chandra (Sanders et al. 2014) as well as CO detected atsub-mm wavelengths (J´achym et al. 2017), co-existing withthe narrow H α tail (Yagi et al. 2007).The deepest HI data on the Coma cluster to date arethose of Bravo-Alfaro et al. (2000) and Bravo-Alfaro et al.(2001), obtained with the VLA in its C–configuration. Theangular resolution of ∼ (cid:48)(cid:48) is not sufficient, though, for de-tailed study of the HI features in the galaxies and some ofthe narrow H α tails (e.g., D100’s with a width of ∼ (cid:48)(cid:48) ). Thedeepest radio 1.4 GHz continuum data on the Coma clusterwere presented in Miller et al. (2009), before the implemen-tation of the far more powerful WIDAR correlator. In thispaper, we present new HI (with higher spatial resolution)and 1.4 GHz continuum (with deeper sensitivity) data on 20RPS galaxies in the Coma cluster.We assume H = 70 km s − Mpc − , Ω m = . , and Ω Λ = . . At the redshift of the Coma cluster ( z = . ),D L =100.7 Mpc and (cid:48)(cid:48) = . kpc. To further study Coma galaxies at radio frequencies,we obtained 1.4 GHz continuum and HI data with theNRAO Karl G. Jansky Very Large Array (VLA) in its B-configuration in two fields centered at NGC 4848 and D100,respectively (Fig. 1), from June 1st to 11th, 2016 (Table 1,program code: SH0174, PI: Sun). The 1.4 GHz continuumdata were taken with two base-bands (A0/C0 and B0/D0)covering a frequency range from 0.9 GHz to 2.1 GHz; eachbase-band is constructed with 7 spectral windows of 128MHz, and each spectral window is divided into 64 channelsof 2 MHz. The HI spectral data were taken with two spec-tral windows (one for A0/C0 and the other for B0/D0) of64 MHz covering a velocity range of 1000 to 11200 km s − .Each spectral window is divided into 3584 channels of 17.8kHz (or 3.88 km s − velocity resolution).The VLA was in the B-configuration for all the observa-tions. 3C286 was observed to calibrate the flux density scaleand the bandpass. J1310+3220 was observed for the calibra-tion of antenna gains and phase. The data were calibratedand reduced with the CASA software; each field was cali-brated separately. Beyond the standard CASA pipeline, weremoved radio frequency interference (RFI) carefully withthe tfcrop and rflag mode of the flagdata task in CASA.About 40% of the data in the continuum band and 20%of the data in the HI band are identified as affected byRFI. Then, phase self-calibration was applied to the contin-uum data as there are strong sources ( −
40 mJy beam − )in the fields which left residuals after the standard calibra-tion. Complex gain calibration solutions were obtained with The National Radio Astronomy Observatory is a facility of theNational Science Foundation operated under cooperative agree-ment by Associated Universities, Inc.MNRAS000
40 mJy beam − )in the fields which left residuals after the standard calibra-tion. Complex gain calibration solutions were obtained with The National Radio Astronomy Observatory is a facility of theNational Science Foundation operated under cooperative agree-ment by Associated Universities, Inc.MNRAS000 , 1–20 (2020)
H. Chen et al.
Table 1.
VLA 1.4 GHz observations in B-configurationField R.A. (J2000) Decl. (J2000) Obs. date Total time (on-source) Continuum beam HI beam Central source1 12h58m02.65s +28 ◦ (cid:48) (cid:48)(cid:48) × . (cid:48)(cid:48) × . (cid:48)(cid:48) . (cid:48)(cid:48) × . (cid:48)(cid:48) NGC 48482 13h00m12.46s +27 ◦ (cid:48) (cid:48)(cid:48) × . (cid:48)(cid:48) × . (cid:48)(cid:48) . (cid:48)(cid:48) × . (cid:48)(cid:48) D100
Table 2.
Coma RPS galaxies studied in this paperNO. Galaxy R.A. (J2000) Decl. (J2000) Velocity rms 1.4 GHz Flux Density Field[h m s] [ ◦ (cid:48) (cid:48)(cid:48) ] [ km s − ] [ µ Jy beam − ] [ mJy ] . ± .
12 KUG 1255+283 (GMP 4555) 12 57 57.74 28 03 42.1 8136 8.04 . ± .
13 NGC 4858 (GMP 3816) 12 59 02.11 28 06 56.4 9416 9.85 . ± .
14 GMP 4570 12 57 56.81 27 59 30.6 4565 9.92 . ± .
15 GMP 4629 12 57 50.27 28 10 13.7 6918 5.73 . ± .
16 D100 (MRK 0060, GMP 2910) 13 00 09.15 27 51 59.4 5316 8.05 . ± .
27 KUG 1258+279A (GMP 2599) 13 00 33.70 27 38 15.6 7485 13.8 . ± .
28 MRK 0058 (GMP 3779) 12 59 05.30 27 38 39.6 5419 19.5 . ± .
29 IC 4040 (GMP 2559) 13 00 37.91 28 03 28.0 7675 10.7 . ± .
210 KUG 1257+278 (GMP 3271) 12 59 39.82 27 34 35.9 5011 17.7 . ± .
211 IC 3949 ∗ (GMP 3896) 12 58 55.89 27 49 59.9 7526 14.5 . ± .
212 GMP 3071 ∗
12 59 56.15 27 44 47.3 8920 9.04 . ± .
213 NGC 4853 ∗ (GMP 4156) 12 58 35.20 27 35 47.1 7688 40.8 . ± .
214 NGC 4911 ∗ (GMP 2374) 13 00 56.08 27 47 26.9 7985 10.4 . ± .
215 NGC 4921 ∗ (GMP 2059) 13 01 26.15 27 53 09.5 5470 14.6 . ± .
216 GMP 4060 ∗
12 58 42.60 27 45 38.0 8686 20.8 < .
217 NGC 4854 ∗ (GMP 4017) 12 58 47.44 27 40 29.3 8383 24.6 < ∗
13 00 01.08 28 04 56.2 7765 10.8 < ∗
13 00 08.07 27 46 24.0 8672 8.54 < ∗ (GMP 2640) 13 00 29.23 27 30 53.7 7395 23.6 < Note:
The properties of 20 RPS galaxies from Smith et al. (2010), Yagi et al. (2010) and Kenney et al. (2015) are listed. R.A., Decl.,and optical velocity are from LEDA (Makarov et al. 2014), except for GMP3016 whose velocity is from NED as there is no LEDAvalue. The local rms and continuum flux density (or 4 σ upper limit) measured by PyBDSF are also shown for each galaxy. The rmsvaries within the same field because of the primary beam correction. The synthesised beam sizes are shown in Table 1. Galaxies withan asterisk have no RPS tails detected in radio continuum. a 60s integration time sampling over 1-3 cycles of phase self-calibration until no further significant improvement was ob-tained. Amplitude self-calibration was tested but was notapplied as this brought no further improvements.To create the HI cube, continuum was carefully fit-ted with line free channels in the range of ± − centred on the systemic velocity of each RPS galaxy; thecontinuum was subtracted with the uvcontsub algorithm.Multi-scale, multi-frequency, multi-term synthesis with w-projection (which is a wide-field imaging technique) wereused in the tclean algorithm to clean the continuum and HIdata. The continuum map used a robust weighting of 0.5 toget a good compromise between optimized spatial resolutionand sensitivity. The HI cube used natural weighting to getthe best sensitivity. The overall radio continuum map of the two observed fieldsis shown in Appendix A. The PyBDSF source-detection package (Mohan & Rafferty 2015) was used to locate andmeasure the flux densities of the radio sources from our data.PyBDSF calculated the local rms for each pixel, evaluatedover a 5 (cid:48) box. Then, source peaks greater than 4 σ are iden-tified and all contiguous pixels higher than 2 σ are identifiedas belonging to one source. Finally, Gaussian fitting is usedto resolve the source position and flux density. In total, 1975and 1173 radio sources (64 of them are duplicates) are de-tected within the 10% response of field 1 and 2, respectively.Our source number density is 13.8 (for field 1) and 3.9 (forfield 2) times of that from Miller et al. (2009) in the cen-tral 6.4 (cid:48) (90% response radius of the VLA field at 1.4 GHz).To assess the reliability of our flux density measurement, wecompared our results with those from Miller et al. (2009)and the FIRST survey (White et al. 1997) (see Appendix Bfor details). Our flux density is consistent with both.With PyBDSF, we also derive the rms distribution foreach field. The rms in the center of fields 1 and 2 are 6 µ Jy and 8 µ Jy which are 4 and 3 times deeper than the Milleret al. (2009) data. Because of the primary beam correction,the rms increases from the center of the field to the outerregions. As NGC 4848 and D100 are at the center of theirfields, they have the deepest data in their respective fields.The full continuum source catalog is shown in Table B1. MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies Right Ascension (J2000) D ec li n a t i o n ( J ) Figure 2.
Left: VLA 1.4 GHz continuum images for RPS galaxies. Green contours show 2, 10 and 50 sigma flux density levels in theradio continuum. Radio emission in galactic disks is measured in symmetrical regions around galaxy centers outlined by solid circles, orellipses, or rectangles. Radio emission in one-sided asymmetric tails is also measured in regions shown by solid rectangles. Dashed blackellipses show D . Right: H α maps for the same galaxies from the Subaru data, except for KUG 1258+279A and NGC4921 (H α mapfrom GOLDMine http://goldmine.mib.infn.it/). The nuclei are marked with a white X. The dashed boxes are where H α emission andupper limits on the radio continuum emission are measured, the values of which are included in Fig. 6. The dashed line in the KUG1258+279A H α map points along the direction that Smith et al. (2010) identified as the ‘tail’.MNRAS000
Left: VLA 1.4 GHz continuum images for RPS galaxies. Green contours show 2, 10 and 50 sigma flux density levels in theradio continuum. Radio emission in galactic disks is measured in symmetrical regions around galaxy centers outlined by solid circles, orellipses, or rectangles. Radio emission in one-sided asymmetric tails is also measured in regions shown by solid rectangles. Dashed blackellipses show D . Right: H α maps for the same galaxies from the Subaru data, except for KUG 1258+279A and NGC4921 (H α mapfrom GOLDMine http://goldmine.mib.infn.it/). The nuclei are marked with a white X. The dashed boxes are where H α emission andupper limits on the radio continuum emission are measured, the values of which are included in Fig. 6. The dashed line in the KUG1258+279A H α map points along the direction that Smith et al. (2010) identified as the ‘tail’.MNRAS000 , 1–20 (2020) H. Chen et al.
Right Ascension (J2000) D ec li n a t i o n ( J ) Figure 2 (b).
MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies Right Ascension (J2000) D ec li n a t i o n ( J ) Figure 2 (c).
MNRAS000
MNRAS000 , 1–20 (2020)
H. Chen et al.
Right Ascension (J2000) D ec li n a t i o n ( J ) Figure 2 (d).
MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies Right Ascension (J2000) D ec li n a t i o n ( J ) Figure 2 (e).
MNRAS000
MNRAS000 , 1–20 (2020) H. Chen et al.
Table 3.
Radio continuum flux densities of RPS galaxies and their tailsNO. Region identification Component 1.4 GHz Flux Density Spectral Index Sky Area Tail Length[mJy] [arcsec ] [kpc]1 NGC 4848a nucleus . ± . − . ± . ± − ± ± − ± ± − ± ± − ± ± − ± ± − ± ± − ± ± ± ± ± − ± ± ± − ± ± − ± ± − ± ± ± − ± ± − ± ± ± Note:
The properties of radio continuum for each galaxy/tail region defined in the radio continuum maps of Fig. 2 are listed. Skyarea refers to the solid angle for each galaxy/tail region. The length of tails is estimated with a cut off at 2-sigma.
KUG1255+283 KUG1258+279A NGC4848 NGC4858 IC4040 −3.0−2.5−2.0−1.5−1.0−0.50.0 I nd e x nucleigalaxiestails Figure 3.
The in–band spectral index for the nucleus, galaxy andtail regions of five RPS galaxies with the spectral index in the tailconstrained.
Twenty RPS galaxies (Smith et al. 2010; Yagi et al. 2010;Kenney et al. 2015) are covered by our observations. Theirradio properties, measured from our data, are listed in Ta-ble 2. 15 of them are detected in the radio continuum, whilethe other five (GMP 4060, NGC 4854, GMP 3016, GMP2923 and KUG 1258+277) were not detected (Table 2).We detected significant extended radio emission co-incident with the H α or SF tails in 10 of the 20 RPSgalaxies (Fig. 2 & Table 3), or in 10 of 15 (66%) RPSgalaxies detected at 1.4 GHz, unambiguously revealing thewidespread occurrence of relativistic electrons and magneticfields in the stripped tails. Radio tails behind IC 4040 and KUG 1255+283 were reported in Miller et al. (2009), whilethe other 8 are new. The galaxy regions and tail regions arevisually defined based on the radio continuum maps (Fig. 2).Symmetrical disks are defined as galaxy regions from ra-dio continuum map by means of circles/ellipses/rectangles.The one-sided asymmetric structures extending beyond thegalaxy regions are defined as tails (solid rectangles). Bothgalaxy regions and tail regions were set to match approxi-mately the 2-sigma contour of the radio emission. The 1.4GHz continuum flux densities of galaxies and tails are listedin Table 3.The radio continuum tail is always spatially coincidentwith the H α or SF tail but usually shorter than either (ex-cept for KUG 1257+278), at least at the sensitivity level ofthe current radio data. All tails, except for KUG 1257+278,extend beyond D (diameter of a galaxy at an isophotallevel of 25 mag arcsec − in the B -band, taken from LEDA(Makarov et al. 2014).)The in-band spectral index was also derived, taking ad-vantage of the wide bandwidth covered by the correlator. Be-cause of the weak signal in general, we divided the 1.4 GHzcontinuum into two parts, a low frequency (approximately0.9 — 1.5 GHz) and a high frequency one (approximately1.5 — 2.1 GHz). Then flux density maps centered at 1.25GHz and 1.75 GHz were created, and the power law indexfor the galaxy and tail was derived using the convention of S ∼ ν α . We only measure the spectral index of a region whenit is detected above 5 σ in both bands. The results are shownin Table 3 and Fig. 3. Our results hint that the tail has asteeper spectrum than the galaxy, which is consistent withaging of electrons and a lack of fresh injection of relativisticelectrons in the tails. But the limited depth and frequencycoverage of our data do not allow us to reach a definite con-clusion. MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies Table 4.
HI deficiency for late-type galaxies in the Coma clusterNO. Galaxy Type D M HInormal rms HI m B i M HI Def HI Def
HI BA00 [arcsec] [ M (cid:12) ] [ µ Jy beam − ] [mag] [degree] [ M (cid:12) ] < . > . -3 NGC 4858 (GMP 3816) Sbc 30.8 14.6 362 15.42 35.6 < . > . > . < . > . -5 GMP 4629 N/A 17.3 4.2 207 17.13 37.8 < . > . -6 D100 (MRK 0060, GMP 2910) I 23.9 9.2 246 15.83 45.0 < . > . -7 KUG 1258+279A (GMP 2599) Sb 33.7 16.7 440 15.32 44.6 < . > . < . > . < . > . -11 IC 3949 (GMP 3896) S 57.3 46.0 533 14.91 90.0 < . > . -12 GMP 3071 S0/a 26.8 6.4 290 16.94 90.0 < . > . -13 NGC 4853 (GMP 4156) S0p 41.5 15.2 1649 14.21 40.3 < . > − . -14 NGC 4911 (GMP 2374) Sb 68.9 69.7 326 13.44 34.7 37.34 0.27 0.5815 NGC 4921 (GMP 2059) Sb 119.7 210.4 504 13.34 24.8 < . > . < . > − . -17 NGC 4854 (GMP 4017) S0 57.3 29.0 1019 14.77 53.8 < . > − . -18 GMP 3016 N/A N/A - 372 N/A N/A - - -19 GMP 2923 E 18.5 3.0 267 17.32 90.0 < . > . -20 KUG 1258+277 (GMP 2640) S0p 29.4 7.6 965 15.62 73.4 < . > − . > . Note:
The expected HI mass ( M HInormal ) is derived from table IV of Haynes & Giovanelli (1984) with D from LEDA and galaxytype from Dressler (1980) (except for GMP 4060 and GMP 2923 from LEDA). For the galaxies without morphology type information(N/A), M HInormal is derived from the average relation of M HI and D for all the galaxies. The rms HI values are calculated for the HIcubes for a channel width of 21.5 km s − . The beam sizes of the HI cubes are shown in Table 1. The values for m B and inclination ( i )are also from LEDA. N/A means not available. Def HI BA00 is based on Def HI from Bravo-Alfaro et al. (2000). Table 5.
HI flux of RPS galaxies Galaxy HI Flux
SoFiA
HI Flux
D25
HI Flux
BA00 [ Jy km s − ] [ Jy km s − ] [ Jy km s − ] NGC 4848 . ± .
005 0 . ± .
132 0 . ± . NGC 4911 . ± .
004 1 . ± .
32 0 . ± . IC 4040 . ± .
005 0 . ± .
11 0 . ± . Note:
HI Flux
SoFiA is the HI flux calculated in the 2 σ
3D mask created by SoFiA. HI Flux
D25 is calculated in the D region definedin the optical (black ellipse in Fig. 4). HI Flux BA00 is measured by Bravo-Alfaro et al. (2000).
The HI data have a spatial resolution of ∼ (cid:48)(cid:48) using naturalweighting to match the narrow width of the gaseous strippedtails of RPS galaxies. The sensitivities of the HI spectra for20 RPS galaxies are listed in Table 4. Although our highspatial resolution HI data are less sensitive than the lowresolution data from Bravo-Alfaro et al. (2000) (observed inC configuration with the VLA) for sources with radii largerthan 30 (cid:48)(cid:48) , the better spatial resolution can resolve the peakHI structure of galaxies and give better HI upper limits forsmall galaxies and narrow tails, such as D100’s tail which isdiscussed in section 4.2.2.HI emission is detected in the integrated intensity (mo-ment 0) maps for NGC 4848, NGC 4911 and IC 4040 asshown in Fig. 4. The HI integrated intensity maps are madefrom the HI cubes with 2 σ
3D masks created by SoFiA(Serra et al. 2015). The D regions are also indicated inFig. 4. The gap in the D region for NGC 4848 is the resultof a mask which was applied to remove an image side-lobeartifact. The positions of HI emission peaks in the integrated intensity maps are consistent with those from Bravo-Alfaroet al. (2000). Corresponding HI spectra for the D regionsare shown in Fig. 5 and the HI results are listed in Ta-ble 5. The HI flux within the SoFiA 3D mask (HI Flux SoFiA )is lower than that within D (HI Flux D25 ), revealing thatthere is some weak HI emission within D that was notpicked up by SoFiA, especially for NGC 4911. Increasingthe radii of the areas of integration to twice D does notlead to a further increase in HI flux, suggesting that we havemeasured the total flux within D . Our HI Flux D25 is largerthan HI Flux
BA00 ; these may be because Bravo-Alfaro et al.(2000) ignored the diffuse HI emission within D where theS/N ratio is lower than 3.Six galaxies (NGC 4848, MRK 0058, KUG 1258+279A,IC 4040, NGC 4911 and NGC 4921) in our sample are de-tected in HI by Bravo-Alfaro et al. (2000), three of which(MRK 0058, KUG 1258+279A and NGC 4921) are not de-tected in this work. The HI in these three galaxies is prob-ably of an extended nature and lacks bright clumps. Also,these galaxies fall at about the 50% response radius of field2 so the sensitivity of the data may not be good enough. MNRAS000
BA00 ; these may be because Bravo-Alfaro et al.(2000) ignored the diffuse HI emission within D where theS/N ratio is lower than 3.Six galaxies (NGC 4848, MRK 0058, KUG 1258+279A,IC 4040, NGC 4911 and NGC 4921) in our sample are de-tected in HI by Bravo-Alfaro et al. (2000), three of which(MRK 0058, KUG 1258+279A and NGC 4921) are not de-tected in this work. The HI in these three galaxies is prob-ably of an extended nature and lacks bright clumps. Also,these galaxies fall at about the 50% response radius of field2 so the sensitivity of the data may not be good enough. MNRAS000 , 1–20 (2020) H. Chen et al.
Figure 4.
Integrated intensity (moment 0) map of HI within the2 σ
3D mask created by SoFiA. The beam size is listed in Table 1.Black ellipses show the D regions. The nuclei of galaxies areshown as black X. Figure 5.
HI spectra for the D regions of three galaxies withHI detected. The HI spectra are fitted with a Gaussian functionshown as the red line. The observed 1.4 GHz continuum should be dominated bysynchrotron radiation. Typically for spiral galaxies, the ra-dio continuum emission is a tracer for star formation as rela-tivistic electrons are accelerated by core-collapse supernovae(e.g., Condon et al. 1991; Bell 2003; Murphy et al. 2008).As H II regions (tracing massive stars) have been found instripped tails (Gavazzi et al. 1995; Sun et al. 2007; Crameret al. 2019), the local supernovae could contribute to thesynchrotron radiation in tails. On the other hand, relativis-tic electrons can be stripped by ram pressure as in radiohead-tail galaxies (e.g., Gavazzi & Jaffe 1987; Murphy et al.2009). In the following we investigate if the observed radiocontinuum emission is mainly related to star formation inthe tail or due to relativistic electrons which were strippedfrom the galaxy by ram pressure. MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies α and IR We compared the 1.4 GHz continuum and H α surface bright-ness (SB) in tails (blue) and galaxies (red) in Fig. 6, to studythe origin of the radio tails. Each number indicates a galaxyas listed in Table 2. Blue arrows show 3 sigma 1.4 GHz SBupper limits for the dashed box regions in Fig. 2. The 1.4GHz SB is the ratio between 1.4 GHz flux density and skyarea (listed in Table 3) for each tail/galaxy region defined inFig. 2. H α SB values for the same tail/galaxy regions are es-timated from the
Subaru data as in Yagi et al. (2017), exceptthat of KUG 1258+279A which is from GOLDMine. Broad-band over–subtraction and [NII], [SII], and [OI] contami-nation are corrected. The intrinsic extinction is unknown.About 1 mag of intrinsic extinction was measured for twoisolated HII regions in the Virgo cluster (Gerhard et al. 2002;Cortese et al. 2004). We simply adopt this value on our H α SB. If the relativistic electrons emitting radio in strippedtails come from local supernovae in the tails, it is expectedthat the same 1.4 GHz — H α relation in galaxies will befollowed. However, all the RPS galaxies, except IC 4040 (9in Fig. 6), show a higher 1.4 GHz to H α ratio in the tailsthan in the galaxies. The mean ratio in the tails ( . × mJy erg − s cm , blue dashed line) is 6.8 times that in galax-ies ( . × mJy erg − s cm , red dashed line). Moreover,H α emission in tails only provides an upper limit on the SFRas the stripping and heating (caused by the RPS shock) ofionized gas can account for much of the diffuse H α emis-sion (e.g., Sun et al. 2010; Fossati et al. 2016; Cramer et al.2019). So the radio continuum in the Coma tails is generallytoo strong compared to the value expected from the SFR.Our results suggest that the radio-emitting plasma in tailsis not formed in situ, but stripped from the galaxy by rampressure. While the above analysis is based on rectangularregions roughly matching the radio emission, we also exam-ined the change on the results by using tail regions coincidingwith a constant surface brightness limit of 20 µ Jy beam − (beam size . (cid:48)(cid:48) × . (cid:48)(cid:48) ). Both radio surface brightness andthe H α surface brightness have been re-measured with thisnew set of regions. The averaged 1.4 GHz to H α ratio in thenew set of the tail regions is now 46% higher than the bluedashed line in Fig. 6, and the scatter still remains about thesame. For simplicity and the sizable scatter here, we stick tothe results from the simple rectangular regions as shown inFig. 2.There is also evidence that the ram pressure may alsoenhance the radio emission of the main bodies of the galax-ies. With the 70, 100 and 160 µ m flux densities for these RPSgalaxies from the Herschel point source catalog, the radio-IRrelations of these RPS galaxies are compared with those forstar forming galaxies (Fig. 7). Because the main bodies aremuch brighter than the tails, the radio and IR emissions ofthe RPS galaxies are dominated by their main bodies. TheTIR luminosity of the RPS galaxies in the Coma cluster isderived from the 70, 100 and 160 µ m flux densities with thecalibration coefficients listed in Table 3 of Galametz et al.(2013). The corresponding far-infrared (FIR, − µ m)fraction was given by the 70 to 160 µ m flux density ratiowith the Equation 3 of Murphy et al. (2008). The IR to 1.4GHz ratio is generally expressed by the parameter q (Bell −18 −17 −16 −15 H α Surface Brightness [erg s −1 cm −2 arcsec −2 ]0.0010.010 . GH z S u rf ace B r i gh t n e ss [ m J y a r c s ec − ] H α Surface Brightness [10 erg s −1 kpc −2 ] . GH z S u rf ace B r i gh t n e ss [ W H z − kp c − ]
12 345 6 78 9
Coma RPS tailsComa RPS galaxies
Figure 6.
The 1.4 GHz surface brightness versus H α surfacebrightness for RPS galaxies and their tails in regions defined inFig. 2. Each number indicates a galaxy as listed in Table 2, andtails are in blue and galaxies are in red. Blue arrows are 1.4 GHzupper limits for the dashed box regions in Fig. 2 where an H α tail is detected but a radio extension is absent from the currentdata. Bottom right arrows show the extinction correction of onemagnitude applied on H α . The mean ratios of 1.4 GHz to H α forRPS tails and galaxies in the Coma cluster are shown as blue andred dashed lines respectively. q IR = log ( IR3 . × W m − ) − log ( S . W m − Hz − ) , (1)The mean values of q TIR and q FIR in the RPS galaxies of theComa cluster ( q TIR = . ± . and q FIR = . ± . , the reddashed lines in Fig. 7) are systematically lower than that instar forming galaxies ( q TIR = . ± . and q FIR = . ± . ,the black solid lines in Fig. 7, Bell 2003; Yun et al. 2001).On average, the 1.4 GHz continuum in the RPS galaxiesin the Coma cluster is enhanced relative to their TIR (bya factor of 1.9) and FIR (by a factor of 2.0) compared tonormal star-forming galaxies. Previous works also reportedthe radio excess in the Coma cluster galaxies (e.g., Gavazzi& Jaffe 1986; Gavazzi et al. 1991) and in the Virgo clustergalaxies (e.g., Murphy et al. 2009; Vollmer et al. 2013). Themean q TIR and q FIR in the Virgo cluster galaxies are . ± . ± . (the blue dashed lines in Fig. 7, Mur-phy et al. 2009). Several mechanisms have been proposed toexplain the enhanced radio continuum in RPS galaxies, in-cluding the enhancement of magnetic field strength causedby ISM shear motions (Murphy et al. 2009), the compressionof ISM/magnetic field by the ICM ram and/or thermal pres-sure (Murphy et al. 2009), and turbulence and ICM shockscaused by ram pressure (Kang & Ryu 2011). The generalradio enhancement in the Coma tails may be related to tur-bulence in the tail and wake of the galaxy, a further investi-gation of which would require more theoretical or simulationwork which is beyond the scope of the current paper. MNRAS , 1–20 (2020) H. Chen et al. TIR luminosity [L O • ]10 . GH z l u m i no s it y [ W H z - ]
12 34 6 7 8 9
10 11 13
SF galaxies (Bell 2003)Virgo RPS galaxies (Murphy+2009)Coma RPS galaxies (This paper) FIR luminosity [L O • ]10 . GH z l u m i no s it y [ W H z - ]
12 34 67 8 9
10 11 13
SF galaxies (Yun+2001)Virgo RPS galaxies (Murphy+2009)Coma RPS galaxies (This paper)
Figure 7.
The 1.4 GHz luminosity versus total infrared (TIR,top plot) and far-infrared (FIR, bottom plot) luminosity for theRPS galaxies. The galaxies in the Coma cluster are in red andthose in the Virgo cluster (Murphy et al. 2009) are in blue. Eachnumber indicates a galaxy as listed in Table 2. The mean ratiosof 1.4 GHz to TIR (FIR) for RPS galaxies in the Coma clusterand in the Virgo cluster are shown as red and blue dashed linesrespectively. The relations between the 1.4 GHz continuum andTIR (Bell 2003) or FIR (Yun et al. 2001) for star forming galaxiesare also shown as black solid lines.
If the 1.4 GHz emission in the stripped tail originates fromthe relativistic electrons stripped from the galaxy by rampressure, the relativistic electrons will cool/age along thetail. The age of the relativistic electrons in the tail, t , isroughly d / v , where d is the distance from the galaxy and v is the mean velocity of electrons stripped from the galaxy.The 1.4 GHz emission can only be detected before the rela-tivistic electrons have cooled, or roughly t < t syn (1.4 GHz),which is equivalent to d < t syn ( . ) × v , where t syn (1.4GHz) is the synchrotron cooling time at 1.4 GHz. With thisassumption, if the tail length detected at 1.4 GHz does notexceed t syn ( . ) × v , it will support the assumption thatthe bulk of the radio continuum in tails is stripped from thegalaxy.Using equation 9 of Feretti & Giovannini (2008), thesynchrotron cooling time (in Myr) can be estimated as: t syn ( ν ) = B . B + B [( + z ) ν ] − . (2) where the magnetic field B is in µ G, the frequency is in GHzand B CMB = . ( + z ) µ G is the magnetic field of theCosmic Microwave Background. Faraday rotation measurestudies indicate that magnetic field strengths in clusters areon the order of a few µ G, with strengths up to tens of µ Gin cluster cores (e.g., Perley & Taylor 1991; Taylor & Per-ley 1993; Feretti et al. 1995, 1999; Taylor et al. 2002, 2006,2007; Bonafede et al. 2010). However, if the magnetic fieldin the tail has its origin in the galaxy, it would be strongerthan the magnetic field in the ICM. Assuming a magneticfield strength of 10 µ G, t syn will be ∼ Myr at 1.4 GHzfor the Coma cluster (z=0.0231). The mean velocity of therelativistic electrons stripped from the galaxy is about 500km s − in simulations (Tonnesen & Bryan 2010). Assumingan isotropic distribution of tail directions, the velocity in theplane of the sky is ∼ km s − ( × π / − ). Then,the maximum expected length of the radio continuum tailat 1.4 GHz in our observations would be 15.2 kpc, which isindeed larger than the observed length of the radio contin-uum tails (2.8–12.6 kpc) in the Coma cluster. The observedsteepening of radio spectra in tails (especially for NGC 4858)also supports synchrotron cooling in the tail. The large un-certainty in the assumed relative velocity of the galaxy withrespect to the ICM and the value for the magnetic field leadto a large range for the maximum length limitation of the ra-dio continuum tails and the above should thus be consideredto be an order of magnitude calculation based on plausiblevalues. There are 10 galaxies without radio continuum tail detec-tions. For 5 of them (GMP 4060, NGC 4854, GMP 3016,GMP 2923 and KUG 1258+277), no radio continuum is de-tected at all from the galaxy (Table 2). There is no deepH α data for KUG 1258+277. The H α tails of the other 4(GMP 4060, NGC 4854, GMP 3016 and GMP 2923, cor-responding to Fig. 4 l, k, e and d in Yagi et al. 2010) areweak and detached from the galaxies. At the same time,GMP 4060, GMP 3016, and GMP 2923 did not show de-tectable H α emission in the galaxies. So perhaps they are ata late evolutionary stage of stripping which might accountfor the weak (undetected) radio emission.There are 5 galaxies (NGC 4911, NGC 4921, NGC 4853,IC 3949 and GMP 3071) with radio continuum detectionsin the galaxies, but radio continuum tails are not detected.NGC 4911 and NGC 4921 are both face-on galaxies. TheirH α extensions are short, which is consistent with the lackof long radio tails there. Whereas NGC 4911 does not havea significant radio tail, its radio continuum within D isasymmetric, with an extension towards the short H α tail.NGC 4853, IC 3949 and GMP 3071 do have significant H α tails, but no radio continuum tails are detected. Deeper datamay be required to probe the radio continuum emission inthe tails of these galaxies. The RPS galaxies should be deficient in HI as it is stripped(Quilis et al. 2000; Boselli et al. 2016a). To quantify the
MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies missing atomic gas in the RPS galaxies in Coma, the HIdeficiency parameter (Def HI , Haynes & Giovanelli 1984), orthe log of the ratio of the normal (i.e., expected) HI massand the observed HI mass, is studied with our HI results.The whole HI flux within D is used to derive the Def HI forNGC 4848, NGC 4911 and IC 4040. For other galaxies, onlya lower limit within D can be placed on Def HI . The upperlimit on the HI flux is S HI ul = × rms × (cid:112) W × W × (cid:113) A galaxy / A beam , (3)where rms is the local noise in the HI data, in units of Jybeam − channel − , W is the expected velocity width of thegalaxy in units of km s − , W is the channel width of our datacube (21.5 km s − ), A galaxy is the D area of the galaxy, and A beam is the area of the beam size for our HI data cube.The expected velocity width of the galaxy is given by W = ( V M sin i ) + W t (Bottinelli et al. 1983), which is derivedfrom the maximum rotational velocity ( V M ) by correctingthe projection of inclination ( i ) and profile broadening ( W t )caused by random/turbulent motions. V M is obtained from m B using the Tully-Fisher relation for individual galaxies(Fig. 13 of Meyer et al. 2016). The values of m B and i arefrom LEDA and listed in Table 4. The value of W t used hereis 5 km s − from Verheijen & Sancisi (2001). The HI mass isgiven by Meyer et al. (2017): M HI = . × ( + z ) × D × S HI , (4)where M HI is the HI mass in M (cid:12) , S HI is the HI flux in Jy km s − , D is the luminosity distance of the galaxy in Mpc(100.7 Mpc used for all the Coma galaxies), and z is the red-shift of the galaxy (0.0231 used for all the Coma galaxies).The normal or expected HI mass ( M HInormal ) is derived basedon table IV of Haynes & Giovanelli (1984), on D (LEDAdata base) and the galaxy type (Dressler 1980, LEDA database). All the results are listed in Table 4.As the HI disks of galaxies within ∼ D (e.g., Chung et al. 2009), our HI deficiencies shouldtherefore be accurate. Def HI for IC 4040 and the lower lim-its on Def HI for NGC 4858, MRK 0058, NGC 4921, KUG1258+279A and KUG 1258+277 are consistent with Bravo-Alfaro et al. (2000) (Def HI BA00 ). However, in this work wefound NGC 4848 and NGC 4911 to be less HI-deficient thanBravo-Alfaro et al. (2000). This is because the HI flux islarger than that in Bravo-Alfaro et al. (2000) which is alsodiscussed in section 3.3. Most (15 of 20) RPS galaxies inour sample are proved to be HI deficient with our detectionand lower limit estimate. Having said that, we should keepin mind that M HI could be an underestimate as any missingflux has not been considered here. The mixing of the stripped cold ISM with the hot ICM willproduce multi-phase gas (e.g., Sun et al. 2007; Ferland et al.2009; J´achym et al. 2019). It is interesting to study howmulti-phase gases co-exist and evolve in the RPS tails. D100is the only galaxy in our sample with CO detections in theRPS tails (J´achym et al. 2017). With our deep and high spa-tial resolved HI observations, the state of co-exist molecularand atomic gas in the RPS tails of D100 could be discussed. We were able to derive the molecular gas mass, M H ,in the narrow tail (regions D100 T1 to D100 T4 in J´achymet al. 2017) of D100 from CO single dish observations andGalactic standard CO– H relation. The lower limit of M H for the whole tail could be got by the sum of D100 T1 toD100 T4 as the tail is not fully covered by the CO obser-vations. We constrain the HI mass limit in the narrow H α tail of D100 ( ∼ (cid:48)(cid:48) width) with the new data ( ∼ (cid:48)(cid:48) spatialresolution vs. ∼ (cid:48)(cid:48) spatial resolution for Bravo-Alfaro et al.(2000)). The HI mass upper limit is derived using equation2 and 3 for the whole tail (4 (cid:48)(cid:48) × (cid:48)(cid:48) ) assuming that theHI velocity width is the same as CO (the velocity coverageof CO in D100 T1 to D100 T4 is about 200 km s − ). TheHI upper limit is also derived for each part of the H α tailcovered by CO observations (D100 T1 to D100 T4). Resultsare listed in Table 6. Integrated over the entire tail there ismore than . × M (cid:12) of H (assuming a Galactic CO-to-H2 conversion factor), and less than . × M (cid:12) of HI.The molecular gas fraction in (parts of) D100’s stripped tailis surprisingly high, which is normally only observed in theinner disk of galaxies. This is consistent with the results in(parts of) the RPS tail of ESO137-001 (J´achym et al. 2014).In galaxies, the molecular to atomic gas ratio ( R H ) cor-relates with the interstellar pressure in the galactic disk(Blitz & Rosolowsky 2004, 2006). The ICM thermal pres-sure ( P / k B ) at the projected position of D100 (0.17 Mpcaway from the cluster centroid) is . × K cm − , whichis derived from model C of the Planck Collaboration et al.(2013). Although the environments are completely different,the high molecular to atomic gas ratio ( R H ) and high ICMthermal pressure in (parts of) D100’s stripped tail agree wellwith the typical relation between R H and pressure in thegalactic disk (Fig. 9 in Krumholz et al. 2009). The high ICMthermal pressure offers an explanation for the molecular gasdominated stripped tail (J´achym et al. 2014). On the otherhand, formation of molecular gas and ionization of atomicgas in the stripped tail also increase R H . We obtained deep VLA 1.4 GHz data in two fields of theComa cluster and focused on 20 RPS galaxies by studyingtheir radio continuum and HI properties. Our main resultsare summarized as follows:(1) Radio continuum tails are found in 10 of 20 rampressure stripped (RPS) galaxies in Coma, revealing thewidespread occurrence of relativistic electrons and magneticfields in the RPS tails. 8 of the RPS tails are new detections.(2) The wide-band 1.4 GHz data allow spectral indicesto be derived which provide a hint that the radio spectraof the tails are steeper than those of the galaxies, as ex-pected from synchrotron cooling without injection of freshrelativistic electrons in the tails.(3) The 1.4 GHz continuum of the stripped tails is en-hanced relative to their SFR traced by H α emission by afactor of ∼ MNRAS000
MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies missing atomic gas in the RPS galaxies in Coma, the HIdeficiency parameter (Def HI , Haynes & Giovanelli 1984), orthe log of the ratio of the normal (i.e., expected) HI massand the observed HI mass, is studied with our HI results.The whole HI flux within D is used to derive the Def HI forNGC 4848, NGC 4911 and IC 4040. For other galaxies, onlya lower limit within D can be placed on Def HI . The upperlimit on the HI flux is S HI ul = × rms × (cid:112) W × W × (cid:113) A galaxy / A beam , (3)where rms is the local noise in the HI data, in units of Jybeam − channel − , W is the expected velocity width of thegalaxy in units of km s − , W is the channel width of our datacube (21.5 km s − ), A galaxy is the D area of the galaxy, and A beam is the area of the beam size for our HI data cube.The expected velocity width of the galaxy is given by W = ( V M sin i ) + W t (Bottinelli et al. 1983), which is derivedfrom the maximum rotational velocity ( V M ) by correctingthe projection of inclination ( i ) and profile broadening ( W t )caused by random/turbulent motions. V M is obtained from m B using the Tully-Fisher relation for individual galaxies(Fig. 13 of Meyer et al. 2016). The values of m B and i arefrom LEDA and listed in Table 4. The value of W t used hereis 5 km s − from Verheijen & Sancisi (2001). The HI mass isgiven by Meyer et al. (2017): M HI = . × ( + z ) × D × S HI , (4)where M HI is the HI mass in M (cid:12) , S HI is the HI flux in Jy km s − , D is the luminosity distance of the galaxy in Mpc(100.7 Mpc used for all the Coma galaxies), and z is the red-shift of the galaxy (0.0231 used for all the Coma galaxies).The normal or expected HI mass ( M HInormal ) is derived basedon table IV of Haynes & Giovanelli (1984), on D (LEDAdata base) and the galaxy type (Dressler 1980, LEDA database). All the results are listed in Table 4.As the HI disks of galaxies within ∼ D (e.g., Chung et al. 2009), our HI deficiencies shouldtherefore be accurate. Def HI for IC 4040 and the lower lim-its on Def HI for NGC 4858, MRK 0058, NGC 4921, KUG1258+279A and KUG 1258+277 are consistent with Bravo-Alfaro et al. (2000) (Def HI BA00 ). However, in this work wefound NGC 4848 and NGC 4911 to be less HI-deficient thanBravo-Alfaro et al. (2000). This is because the HI flux islarger than that in Bravo-Alfaro et al. (2000) which is alsodiscussed in section 3.3. Most (15 of 20) RPS galaxies inour sample are proved to be HI deficient with our detectionand lower limit estimate. Having said that, we should keepin mind that M HI could be an underestimate as any missingflux has not been considered here. The mixing of the stripped cold ISM with the hot ICM willproduce multi-phase gas (e.g., Sun et al. 2007; Ferland et al.2009; J´achym et al. 2019). It is interesting to study howmulti-phase gases co-exist and evolve in the RPS tails. D100is the only galaxy in our sample with CO detections in theRPS tails (J´achym et al. 2017). With our deep and high spa-tial resolved HI observations, the state of co-exist molecularand atomic gas in the RPS tails of D100 could be discussed. We were able to derive the molecular gas mass, M H ,in the narrow tail (regions D100 T1 to D100 T4 in J´achymet al. 2017) of D100 from CO single dish observations andGalactic standard CO– H relation. The lower limit of M H for the whole tail could be got by the sum of D100 T1 toD100 T4 as the tail is not fully covered by the CO obser-vations. We constrain the HI mass limit in the narrow H α tail of D100 ( ∼ (cid:48)(cid:48) width) with the new data ( ∼ (cid:48)(cid:48) spatialresolution vs. ∼ (cid:48)(cid:48) spatial resolution for Bravo-Alfaro et al.(2000)). The HI mass upper limit is derived using equation2 and 3 for the whole tail (4 (cid:48)(cid:48) × (cid:48)(cid:48) ) assuming that theHI velocity width is the same as CO (the velocity coverageof CO in D100 T1 to D100 T4 is about 200 km s − ). TheHI upper limit is also derived for each part of the H α tailcovered by CO observations (D100 T1 to D100 T4). Resultsare listed in Table 6. Integrated over the entire tail there ismore than . × M (cid:12) of H (assuming a Galactic CO-to-H2 conversion factor), and less than . × M (cid:12) of HI.The molecular gas fraction in (parts of) D100’s stripped tailis surprisingly high, which is normally only observed in theinner disk of galaxies. This is consistent with the results in(parts of) the RPS tail of ESO137-001 (J´achym et al. 2014).In galaxies, the molecular to atomic gas ratio ( R H ) cor-relates with the interstellar pressure in the galactic disk(Blitz & Rosolowsky 2004, 2006). The ICM thermal pres-sure ( P / k B ) at the projected position of D100 (0.17 Mpcaway from the cluster centroid) is . × K cm − , whichis derived from model C of the Planck Collaboration et al.(2013). Although the environments are completely different,the high molecular to atomic gas ratio ( R H ) and high ICMthermal pressure in (parts of) D100’s stripped tail agree wellwith the typical relation between R H and pressure in thegalactic disk (Fig. 9 in Krumholz et al. 2009). The high ICMthermal pressure offers an explanation for the molecular gasdominated stripped tail (J´achym et al. 2014). On the otherhand, formation of molecular gas and ionization of atomicgas in the stripped tail also increase R H . We obtained deep VLA 1.4 GHz data in two fields of theComa cluster and focused on 20 RPS galaxies by studyingtheir radio continuum and HI properties. Our main resultsare summarized as follows:(1) Radio continuum tails are found in 10 of 20 rampressure stripped (RPS) galaxies in Coma, revealing thewidespread occurrence of relativistic electrons and magneticfields in the RPS tails. 8 of the RPS tails are new detections.(2) The wide-band 1.4 GHz data allow spectral indicesto be derived which provide a hint that the radio spectraof the tails are steeper than those of the galaxies, as ex-pected from synchrotron cooling without injection of freshrelativistic electrons in the tails.(3) The 1.4 GHz continuum of the stripped tails is en-hanced relative to their SFR traced by H α emission by afactor of ∼ MNRAS000 , 1–20 (2020) H. Chen et al.
Table 6.
HI in the stripped tail of D100Tail Region Area HI FWHM CO M HI M H R H2 P/k B [km s − ] [ M (cid:12) ] [ M (cid:12) ] [10 K cm − ]D100 T1 (cid:48)(cid:48) × . (cid:48)(cid:48) < .
95 0 . ± . > (cid:48)(cid:48) × . (cid:48)(cid:48) < ± > (cid:48)(cid:48) × . (cid:48)(cid:48) < ± > (cid:48)(cid:48) × . (cid:48)(cid:48) < ± > (cid:48)(cid:48) × . (cid:48)(cid:48) < > > Note:
The HI mass upper limit in the stripped tail of D100 is derived from the data presented here. The size of the HI tail is set bythe H α emission ( (cid:48)(cid:48) × . (cid:48)(cid:48) ); the velocity width is assumed to be the same as CO (FWHM CO ). D100 M H2 results adjusted to thecosmology of this paper are from CO 1–0 (J´achym et al. 2017). The exception is the tail region of T1 where we used CO 2–1 becausethe beam of CO 1–0 covered part of disk. P/K B is the ICM thermal pressure at the projected position of D100. by a factor of ∼
2, compared with normal SF galaxies. Rampressure interaction may enhance radio emission in the mainbodies of the galaxies.(5) The length of the radio continuum tails is consistentwith the distance being limited by the synchrotron coolingtime and stripping velocity.(6) HI detections in three RPS galaxies (NGC 4848,NGC 4911 and IC 4040) and HI upper limits for the othergalaxies are presented as well. Most (15 of 20) RPS galaxiesare consistent with being deficient in HI.(7) The H /HI mass ratio for the H α tail of D100 ex-ceeds 1.6 — the cold gas in D100’s stripped tail is dominatedby molecular gas. This is likely a consequence of the highICM pressure in the stripped tail.Our continuum data imply that there are magneticfields in the stripped tails, which could explain the colli-mated and bifurcated structure of the stripped gas. SF maybe less efficient in a stripped tail due to suppression by mag-netic fields and turbulence. We suggest that the observedradio continuum emission is due to ram pressure strippingof relativistic electrons from the host galaxy, and not somuch to any local SF. This is consistent with our findingof a steeper radio spectrum in the stripped tail compared tothat of the host galaxy, the overly luminous radio continuumflux density compared to the H α , and the shorter tail lengthin radio (limited by the synchrotron cooling time) than inH α . Our results also suggest that cluster late-type galaxiescan inject relativistic electrons into the ICM, which couldpotentially feed radio relics in clusters (Ge et al. 2019). ACKNOWLEDGEMENTS
We acknowledge useful discussions with the NRAO helpdesk staff, Xing Lu, Ruoyu Liu, Neal Miller and StephenWalker. Support for this work was provided by NSF grant1714764 and NASA/EPSCoR grant NNX15AK29A. We alsoacknowledge the support from the National Aeronauticsand Space Administration through
Chandra
Award Num-ber GO6-17111X and GO6-17127X issued by the
Chandra
X-ray Center, which is operated by the Smithsonian Astro-physical Observatory for and on behalf of the National Aero-nautics Space Administration under contract NAS8-03060.HC’s work has been supported by the South African De-partment of Science and Innovation and the National Re-search Foundation through a Fellowship within the SARAOResearch Chair held by RC Kraan-Korteweg. MF has re- ceived funding from the European Research Council (ERC)under the European Union’s Horizon 2020 research and in-novation programme (grant agreement No 757535). This re-search has made use of data and/or software provided by theHigh Energy Astrophysics Science Archive Research Center(HEASARC), which is a service of the Astrophysics ScienceDivision at NASA/GSFC and the High Energy AstrophysicsDivision of the Smithsonian Astrophysical Observatory. Thisresearch has made use of the NASA/IPAC Infrared ScienceArchive, which is funded by the National Aeronautics andSpace Administration and operated by the California Insti-tute of Technology. Based in part on data collected at SubaruTelescope, which is operated by the National AstronomicalObservatory of Japan.
DATA AVAILABILITY
Most of the data underlying this article are available in thearticle and in its online supplementary material. All the datawill be shared on reasonable request to the correspondingauthors.
REFERENCES
Abramson A., Kenney J. D. P., 2014, AJ, 147, 63Bell E. F., 2003, ApJ, 586, 794Bellhouse C., et al., 2017, ApJ, 844, 49Blitz L., Rosolowsky E., 2004, ApJ, 612, L29Blitz L., Rosolowsky E., 2006, ApJ, 650, 933Boissier S., et al., 2012, A&A, 545, A142Bonafede A., Feretti L., Murgia M., Govoni F., Giovannini G.,Dallacasa D., Dolag K., Taylor G. B., 2010, A&A, 513, A30Boselli A., et al., 2016a, A&A, 587, A68Boselli A., et al., 2016b, A&A, 596, A11Boselli A., et al., 2018, A&A, 615, A114Bottinelli L., Gouguenheim L., Paturel G., de Vaucouleurs G.,1983, A&A, 118, 4Bravo-Alfaro H., Cayatte V., van Gorkom J. H., Balkowski C.,2000, AJ, 119, 580Bravo-Alfaro H., Cayatte V., van Gorkom J. H., Balkowski C.,2001, A&A, 379, 347Butler A., et al., 2018, A&A, 620, A16Chung A., van Gorkom J. H., Kenney J. D. P., Vollmer B., 2007,ApJ, 659, L115Chung A., van Gorkom J. H., Kenney J. D. P., Crowl H., VollmerB., 2009, AJ, 138, 1741Condon J. J., Anderson M. L., Helou G., 1991, ApJ, 376, 95MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies Condon J. J., Cotton W. D., Greisen E. W., Yin Q. F., PerleyR. A., Taylor G. B., Broderick J. J., 1998, AJ, 115, 1693Cortese L., Gavazzi G., Boselli A., Iglesias-Paramo J., 2004, A&A,416, 119Cortese L., et al., 2007, MNRAS, 376, 157Cramer W. J., Kenney J. D. P., Sun M., Crowl H., Yagi M.,J´achym P., Roediger E., Waldron W., 2019, ApJ, 870, 63Crowl H. H., Kenney J. D. P., van Gorkom J. H., Vollmer B.,2005, AJ, 130, 65Crowl H. H., Kenney J. D., van Gorkom J. H., Chung A., RoseJ. A., 2006, in American Astronomical Society Meeting Ab-stracts. p. 211.11Deb T., et al., 2020, arXiv e-prints, p. arXiv:2004.04754Dressler A., 1980, ApJS, 42, 565Ebeling H., Stephenson L. N., Edge A. C., 2014, ApJ, 781, L40Feretti L., Giovannini G., 2008, in Plionis M., L´opez-Cruz O.,Hughes D., eds, Lecture Notes in Physics, Berlin SpringerVerlag Vol. 740, A Pan-Chromatic View of Clusters ofGalaxies and the Large-Scale Structure. p. 24 ( arXiv:astro-ph/0703494 ), doi:10.1007/978-1-4020-6941-3 5Feretti L., Dallacasa D., Giovannini G., Tagliani A., 1995, A&A,302, 680Feretti L., Dallacasa D., Govoni F., Giovannini G., Taylor G. B.,Klein U., 1999, A&A, 344, 472Ferland G. J., Fabian A. C., Hatch N. A., Johnstone R. M., PorterR. L., van Hoof P. A. M., Williams R. J. R., 2009, MNRAS,392, 1475Fossati M., Gavazzi G., Boselli A., Fumagalli M., 2012, A&A, 544,A128Fossati M., Fumagalli M., Boselli A., Gavazzi G., Sun M., WilmanD. J., 2016, MNRAS, 455, 2028Fossati M., et al., 2018, A&A, 614, A57Fumagalli M., Fossati M., Hau G. K. T., Gavazzi G., Bower R.,Sun M., Boselli A., 2014, MNRAS, 445, 4335Galametz M., et al., 2013, MNRAS, 431, 1956Gavazzi G., 1978, A&A, 69, 355Gavazzi G., Jaffe W., 1985, ApJ, 294, L89Gavazzi G., Jaffe W., 1986, ApJ, 310, 53Gavazzi G., Jaffe W., 1987, A&A, 186, L1Gavazzi G., Boselli A., Kennicutt R., 1991, AJ, 101, 1207Gavazzi G., Contursi A., Carrasco L., Boselli A., Kennicutt R.,Scodeggio M., Jaffe W., 1995, A&A, 304, 325Gavazzi G., Boselli A., Mayer L., Iglesias-Paramo J., V´ılchezJ. M., Carrasco L., 2001, ApJ, 563, L23Gavazzi G., Consolandi G., Yagi M., Yoshida M., 2017, A&A,606, A131Gavazzi G., Consolandi G., Gutierrez M. L., Boselli A., YoshidaM., 2018, A&A, 618, A130Ge C., et al., 2019, MNRAS, 486, L36Gerhard O., Arnaboldi M., Freeman K. C., Okamura S., 2002,ApJ, 580, L121Gunn J. E., Gott III J. R., 1972, ApJ, 176, 1Haynes M. P., Giovanelli R., 1984, AJ, 89, 758Helou G., Soifer B. T., Rowan-Robinson M., 1985, ApJ, 298, L7J´achym P., Kenney J. D. P., Rˇzuiˇcka A., Sun M., Combes F.,Palouˇs J., 2013, A&A, 556, A99J´achym P., Combes F., Cortese L., Sun M., Kenney J. D. P., 2014,ApJ, 792, 11J´achym P., et al., 2017, ApJ, 839, 114J´achym P., et al., 2019, ApJ, 883, 145Kang H., Ryu D., 2011, ApJ, 734, 18Kapferer W., Sluka C., Schindler S., Ferrari C., Ziegler B., 2009,A&A, 499, 87Kenney J. D. P., van Gorkom J. H., Vollmer B., 2004, AJ, 127,3361Kenney J. D. P., Geha M., J´achym P., Crowl H. H., Dague W.,Chung A., van Gorkom J., Vollmer B., 2014, ApJ, 780, 119 Kenney J. D. P., Abramson A., Bravo-Alfaro H., 2015, AJ, 150,59Kim K.-T., 1994, A&AS, 105, 403Koda J., Yagi M., Yamanoi H., Komiyama Y., 2015, ApJ, 807,L2Koopmann R. A., Kenney J. D. P., 2004, ApJ, 613, 866Krumholz M. R., McKee C. F., Tumlinson J., 2009, ApJ, 693, 216Makarov D., Prugniel P., Terekhova N., Courtois H., Vauglin I.,2014, A&A, 570, A13Meyer S. A., Meyer M., Obreschkow D., Staveley-Smith L., 2016,MNRAS, 455, 3136Meyer M., Robotham A., Obreschkow D., Westmeier T., DuffyA. R., Staveley-Smith L., 2017, Publ. Astron. Soc. Australia,34, 52Miller N. A., Hornschemeier A. E., Mobasher B., 2009, AJ, 137,4436Mohan N., Rafferty D., 2015, PyBDSF: Python Blob Detec-tion and Source Finder, Astrophysics Source Code Library(ascl:1502.007)Moretti A., et al., 2018, MNRAS, 480, 2508Murphy E. J., Helou G., Kenney J. D. P., Armus L., Braun R.,2008, ApJ, 678, 828Murphy E. J., Kenney J. D. P., Helou G., Chung A., Howell J. H.,2009, ApJ, 694, 1435Nulsen P. E. J., 1982, MNRAS, 198, 1007Oosterloo T., van Gorkom J., 2005, A&A, 437, L19Owers M. S., Couch W. J., Nulsen P. E. J., Randall S. W., 2012,ApJ, 750, L23Perley R. A., Taylor G. B., 1991, AJ, 101, 1623Pinzke A., Oh S. P., Pfrommer C., 2013, MNRAS, 435, 1061Planck Collaboration et al., 2013, A&A, 554, A140Poggianti B. M., et al., 2016, AJ, 151, 78Poggianti B. M., et al., 2017, ApJ, 844, 48Quilis V., Moore B., Bower R., 2000, Science, 288, 1617Ramatsoku M., et al., 2019, MNRAS, 487, 4580Roediger E., Br¨uggen M., 2008, MNRAS, 388, 465Roediger E., Bruggen M., Owers M. S., Ebeling H., Sun M., 2014,MNRAS, 443, L114Ruszkowski M., 2012, The Role of Magnetic Fields and Micro-physics in Ram Pressure Stripping, NASA ATP ProposalSanders J. S., Fabian A. C., Sun M., Churazov E., Simionescu A.,Walker S. A., Werner N., 2014, MNRAS, 439, 1182Scott T. C., et al., 2010, MNRAS, 403, 1175Scott T. C., Cortese L., Brinks E., Bravo-Alfaro H., Auld R.,Minchin R., 2012, MNRAS, 419, L19Scott T. C., Usero A., Brinks E., Boselli A., Cortese L., Bravo-Alfaro H., 2013, MNRAS, 429, 221Scott T. C., Usero A., Brinks E., Bravo-Alfaro H., Cortese L.,Boselli A., Argudo-Fern´andez M., 2015, MNRAS, 453, 328Serra P., et al., 2015, MNRAS, 448, 1922Serra P., et al., 2019, A&A, 628, A122Sivanandam S., Rieke M. J., Rieke G. H., 2010, ApJ, 717, 147Sivanandam S., Rieke M. J., Rieke G. H., 2014, ApJ, 796, 89Smith R. J., et al., 2010, MNRAS, 408, 1417Smolˇci´c V., et al., 2017, A&A, 602, A1Snowden S. L., Mushotzky R. F., Kuntz K. D., Davis D. S., 2008,A&A, 478, 615Struble M. F., 2018, MNRAS, 473, 4686Sun M., Jones C., Forman W., Nulsen P. E. J., Donahue M., VoitG. M., 2006, ApJ, 637, L81Sun M., Donahue M., Voit G. M., 2007, ApJ, 671, 190Sun M., Donahue M., Roediger E., Nulsen P. E. J., Voit G. M.,Sarazin C., Forman W., Jones C., 2010, ApJ, 708, 946Taylor G. B., Perley R. A., 1993, ApJ, 416, 554Taylor G. B., Fabian A. C., Allen S. W., 2002, MNRAS, 334, 769Taylor G. B., Gugliucci N. E., Fabian A. C., Sanders J. S., GentileG., Allen S. W., 2006, MNRAS, 368, 1500MNRAS000
Abramson A., Kenney J. D. P., 2014, AJ, 147, 63Bell E. F., 2003, ApJ, 586, 794Bellhouse C., et al., 2017, ApJ, 844, 49Blitz L., Rosolowsky E., 2004, ApJ, 612, L29Blitz L., Rosolowsky E., 2006, ApJ, 650, 933Boissier S., et al., 2012, A&A, 545, A142Bonafede A., Feretti L., Murgia M., Govoni F., Giovannini G.,Dallacasa D., Dolag K., Taylor G. B., 2010, A&A, 513, A30Boselli A., et al., 2016a, A&A, 587, A68Boselli A., et al., 2016b, A&A, 596, A11Boselli A., et al., 2018, A&A, 615, A114Bottinelli L., Gouguenheim L., Paturel G., de Vaucouleurs G.,1983, A&A, 118, 4Bravo-Alfaro H., Cayatte V., van Gorkom J. H., Balkowski C.,2000, AJ, 119, 580Bravo-Alfaro H., Cayatte V., van Gorkom J. H., Balkowski C.,2001, A&A, 379, 347Butler A., et al., 2018, A&A, 620, A16Chung A., van Gorkom J. H., Kenney J. D. P., Vollmer B., 2007,ApJ, 659, L115Chung A., van Gorkom J. H., Kenney J. D. P., Crowl H., VollmerB., 2009, AJ, 138, 1741Condon J. J., Anderson M. L., Helou G., 1991, ApJ, 376, 95MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies Condon J. J., Cotton W. D., Greisen E. W., Yin Q. F., PerleyR. A., Taylor G. B., Broderick J. J., 1998, AJ, 115, 1693Cortese L., Gavazzi G., Boselli A., Iglesias-Paramo J., 2004, A&A,416, 119Cortese L., et al., 2007, MNRAS, 376, 157Cramer W. J., Kenney J. D. P., Sun M., Crowl H., Yagi M.,J´achym P., Roediger E., Waldron W., 2019, ApJ, 870, 63Crowl H. H., Kenney J. D. P., van Gorkom J. H., Vollmer B.,2005, AJ, 130, 65Crowl H. H., Kenney J. D., van Gorkom J. H., Chung A., RoseJ. A., 2006, in American Astronomical Society Meeting Ab-stracts. p. 211.11Deb T., et al., 2020, arXiv e-prints, p. arXiv:2004.04754Dressler A., 1980, ApJS, 42, 565Ebeling H., Stephenson L. N., Edge A. C., 2014, ApJ, 781, L40Feretti L., Giovannini G., 2008, in Plionis M., L´opez-Cruz O.,Hughes D., eds, Lecture Notes in Physics, Berlin SpringerVerlag Vol. 740, A Pan-Chromatic View of Clusters ofGalaxies and the Large-Scale Structure. p. 24 ( arXiv:astro-ph/0703494 ), doi:10.1007/978-1-4020-6941-3 5Feretti L., Dallacasa D., Giovannini G., Tagliani A., 1995, A&A,302, 680Feretti L., Dallacasa D., Govoni F., Giovannini G., Taylor G. B.,Klein U., 1999, A&A, 344, 472Ferland G. J., Fabian A. C., Hatch N. A., Johnstone R. M., PorterR. L., van Hoof P. A. M., Williams R. J. R., 2009, MNRAS,392, 1475Fossati M., Gavazzi G., Boselli A., Fumagalli M., 2012, A&A, 544,A128Fossati M., Fumagalli M., Boselli A., Gavazzi G., Sun M., WilmanD. J., 2016, MNRAS, 455, 2028Fossati M., et al., 2018, A&A, 614, A57Fumagalli M., Fossati M., Hau G. K. T., Gavazzi G., Bower R.,Sun M., Boselli A., 2014, MNRAS, 445, 4335Galametz M., et al., 2013, MNRAS, 431, 1956Gavazzi G., 1978, A&A, 69, 355Gavazzi G., Jaffe W., 1985, ApJ, 294, L89Gavazzi G., Jaffe W., 1986, ApJ, 310, 53Gavazzi G., Jaffe W., 1987, A&A, 186, L1Gavazzi G., Boselli A., Kennicutt R., 1991, AJ, 101, 1207Gavazzi G., Contursi A., Carrasco L., Boselli A., Kennicutt R.,Scodeggio M., Jaffe W., 1995, A&A, 304, 325Gavazzi G., Boselli A., Mayer L., Iglesias-Paramo J., V´ılchezJ. M., Carrasco L., 2001, ApJ, 563, L23Gavazzi G., Consolandi G., Yagi M., Yoshida M., 2017, A&A,606, A131Gavazzi G., Consolandi G., Gutierrez M. L., Boselli A., YoshidaM., 2018, A&A, 618, A130Ge C., et al., 2019, MNRAS, 486, L36Gerhard O., Arnaboldi M., Freeman K. C., Okamura S., 2002,ApJ, 580, L121Gunn J. E., Gott III J. R., 1972, ApJ, 176, 1Haynes M. P., Giovanelli R., 1984, AJ, 89, 758Helou G., Soifer B. T., Rowan-Robinson M., 1985, ApJ, 298, L7J´achym P., Kenney J. D. P., Rˇzuiˇcka A., Sun M., Combes F.,Palouˇs J., 2013, A&A, 556, A99J´achym P., Combes F., Cortese L., Sun M., Kenney J. D. P., 2014,ApJ, 792, 11J´achym P., et al., 2017, ApJ, 839, 114J´achym P., et al., 2019, ApJ, 883, 145Kang H., Ryu D., 2011, ApJ, 734, 18Kapferer W., Sluka C., Schindler S., Ferrari C., Ziegler B., 2009,A&A, 499, 87Kenney J. D. P., van Gorkom J. H., Vollmer B., 2004, AJ, 127,3361Kenney J. D. P., Geha M., J´achym P., Crowl H. H., Dague W.,Chung A., van Gorkom J., Vollmer B., 2014, ApJ, 780, 119 Kenney J. D. P., Abramson A., Bravo-Alfaro H., 2015, AJ, 150,59Kim K.-T., 1994, A&AS, 105, 403Koda J., Yagi M., Yamanoi H., Komiyama Y., 2015, ApJ, 807,L2Koopmann R. A., Kenney J. D. P., 2004, ApJ, 613, 866Krumholz M. R., McKee C. F., Tumlinson J., 2009, ApJ, 693, 216Makarov D., Prugniel P., Terekhova N., Courtois H., Vauglin I.,2014, A&A, 570, A13Meyer S. A., Meyer M., Obreschkow D., Staveley-Smith L., 2016,MNRAS, 455, 3136Meyer M., Robotham A., Obreschkow D., Westmeier T., DuffyA. R., Staveley-Smith L., 2017, Publ. Astron. Soc. Australia,34, 52Miller N. A., Hornschemeier A. E., Mobasher B., 2009, AJ, 137,4436Mohan N., Rafferty D., 2015, PyBDSF: Python Blob Detec-tion and Source Finder, Astrophysics Source Code Library(ascl:1502.007)Moretti A., et al., 2018, MNRAS, 480, 2508Murphy E. J., Helou G., Kenney J. D. P., Armus L., Braun R.,2008, ApJ, 678, 828Murphy E. J., Kenney J. D. P., Helou G., Chung A., Howell J. H.,2009, ApJ, 694, 1435Nulsen P. E. J., 1982, MNRAS, 198, 1007Oosterloo T., van Gorkom J., 2005, A&A, 437, L19Owers M. S., Couch W. J., Nulsen P. E. J., Randall S. W., 2012,ApJ, 750, L23Perley R. A., Taylor G. B., 1991, AJ, 101, 1623Pinzke A., Oh S. P., Pfrommer C., 2013, MNRAS, 435, 1061Planck Collaboration et al., 2013, A&A, 554, A140Poggianti B. M., et al., 2016, AJ, 151, 78Poggianti B. M., et al., 2017, ApJ, 844, 48Quilis V., Moore B., Bower R., 2000, Science, 288, 1617Ramatsoku M., et al., 2019, MNRAS, 487, 4580Roediger E., Br¨uggen M., 2008, MNRAS, 388, 465Roediger E., Bruggen M., Owers M. S., Ebeling H., Sun M., 2014,MNRAS, 443, L114Ruszkowski M., 2012, The Role of Magnetic Fields and Micro-physics in Ram Pressure Stripping, NASA ATP ProposalSanders J. S., Fabian A. C., Sun M., Churazov E., Simionescu A.,Walker S. A., Werner N., 2014, MNRAS, 439, 1182Scott T. C., et al., 2010, MNRAS, 403, 1175Scott T. C., Cortese L., Brinks E., Bravo-Alfaro H., Auld R.,Minchin R., 2012, MNRAS, 419, L19Scott T. C., Usero A., Brinks E., Boselli A., Cortese L., Bravo-Alfaro H., 2013, MNRAS, 429, 221Scott T. C., Usero A., Brinks E., Bravo-Alfaro H., Cortese L.,Boselli A., Argudo-Fern´andez M., 2015, MNRAS, 453, 328Serra P., et al., 2015, MNRAS, 448, 1922Serra P., et al., 2019, A&A, 628, A122Sivanandam S., Rieke M. J., Rieke G. H., 2010, ApJ, 717, 147Sivanandam S., Rieke M. J., Rieke G. H., 2014, ApJ, 796, 89Smith R. J., et al., 2010, MNRAS, 408, 1417Smolˇci´c V., et al., 2017, A&A, 602, A1Snowden S. L., Mushotzky R. F., Kuntz K. D., Davis D. S., 2008,A&A, 478, 615Struble M. F., 2018, MNRAS, 473, 4686Sun M., Jones C., Forman W., Nulsen P. E. J., Donahue M., VoitG. M., 2006, ApJ, 637, L81Sun M., Donahue M., Voit G. M., 2007, ApJ, 671, 190Sun M., Donahue M., Roediger E., Nulsen P. E. J., Voit G. M.,Sarazin C., Forman W., Jones C., 2010, ApJ, 708, 946Taylor G. B., Perley R. A., 1993, ApJ, 416, 554Taylor G. B., Fabian A. C., Allen S. W., 2002, MNRAS, 334, 769Taylor G. B., Gugliucci N. E., Fabian A. C., Sanders J. S., GentileG., Allen S. W., 2006, MNRAS, 368, 1500MNRAS000 , 1–20 (2020) H. Chen et al.
Taylor G. B., Fabian A. C., Gentile G., Allen S. W., CrawfordC., Sanders J. S., 2007, MNRAS, 382, 67Tonnesen S., Bryan G. L., 2010, ApJ, 709, 1203Tonnesen S., Bryan G. L., 2012, MNRAS, 422, 1609Tonnesen S., Bryan G. L., Chen R., 2011, ApJ, 731, 98Valentijn E. A., Perola G. C., Jaffe W. J., 1977, A&AS, 28, 333Verdugo C., Combes F., Dasyra K., Salom´e P., Braine J., 2015,A&A, 582, A6Verheijen M. A. W., Sancisi R., 2001, A&A, 370, 765Vollmer B., Beck R., Kenney J. D. P., van Gorkom J. H., 2004,AJ, 127, 3375Vollmer B., Soida M., Beck R., Chung A., Urbanik M., Chy˙zyK. T., Otmianowska-Mazur K., Kenney J. D. P., 2013, A&A,553, A116White R. L., Becker R. H., Helfand D. J., Gregg M. D., 1997,ApJ, 475, 479Wold I. G. B., Owen F. N., Wang W.-H., Barger A. J., KeenanR. C., 2012, ApJS, 202, 2Yagi M., Komiyama Y., Yoshida M., Furusawa H., Kashikawa N.,Koyama Y., Okamura S., 2007, ApJ, 660, 1209Yagi M., et al., 2010, AJ, 140, 1814Yagi M., Gu L., Fujita Y., Nakazawa K., Akahori T., Hattori T.,Yoshida M., Makishima K., 2013, ApJ, 778, 91Yagi M., Koda J., Komiyama Y., Yamanoi H., 2016, ApJS, 225,11Yagi M., Yoshida M., Gavazzi G., Komiyama Y., Kashikawa N.,Okamura S., 2017, ApJ, 839, 65Yoshida M., et al., 2008, ApJ, 688, 918Yoshida M., Yagi M., Komiyama Y., Furusawa H., Kashikawa N.,Hattori T., Okamura S., 2012, ApJ, 749, 43Yun M. S., Reddy N. A., Condon J. J., 2001, ApJ, 554, 803Zhang B., et al., 2013, ApJ, 777, 122
APPENDIX A: RADIO CONTINUUM MAP
A mosaic image of our radio continuum observations isshown in Fig. A1. It is the combination of the 30% responseregions of the two fields. Magnified images of two sourcesare shown at the right to show them in detail. The narrowtail near the end of NGC 4869 is shown better in our imagethan in Miller et al. (2009). The source in the lower rightpanel (the small dashed-line box on the left) has an interest-ing complex structure which is not covered by Miller et al.(2009). This source is cataloged as 1254+28W07 (Valentijnet al. 1977; Kim 1994) and NVSS J125720+282727 (Con-don et al. 1998), but the complex structure is not resolvedbecause of the low spatial resolution ( > (cid:48)(cid:48) ). APPENDIX B: FLUX DENSITYMEASUREMENTS
The full continuum source catalog is shown in Table B1.As Fig. B1 shows, our deep observations reveal many morecontinuum sources than Miller et al. (2009). To assessthe reliability of the flux density measurement of ourwork, we compared our results with Miller et al. (2009)and the FIRST survey (White et al. 1997). Because thespatial resolution is about 5 (cid:48)(cid:48) , we cross match Miller’s andFIRST sources with ours for match radii of 5 (cid:48)(cid:48) . Sources areexcluded when more than one match is found within 5 (cid:48)(cid:48) .The comparisons are shown in Fig. B2. Our flux densitymeasurements are consistent with those from Miller et al.(2009) and FIRST. There is still some scatter, which could be caused by: 1) different algorithms of source detection(PyBDSF in our work and SAD in Miller et al. (2009) andFIRST); and/or 2) varying radio sources (e.g., AGN). Wealso applied the PyBDSF algorithms to the data of Milleret al. (2009) to compare the flux density measurements.By applying PyBDSF on both our data and the Milleret al. (2009) data, the consistency on the flux densitymeasurements is improved, with the root-mean-squarederror of the fits decreasing from 3.0 mJy to 1.3 mJy. Linearfits of comparison show that: S Miller + [ mJy ] = ( . ± . ) × S Ours [ mJy ] S FIRST [ mJy ] = ( . ± . ) × S Ours [ mJy ] S Miller + [ mJy ] = ( . ± . ) × S PyBDSF Miller [ mJy ] Power-law fits confirm the linear relations betweenour flux density and the Miller et al. (2009) flux density (anindex of . ± . ), and the FIRST flux density (an indexof . ± . ). APPENDIX C: 1.4 GHZ CONTINUUM SOURCECOUNTS
Our new data generated the largest 1.4 GHz source catalogin the Coma cluster. We compare source counts in our re-gions with those from previous work in Fig. C1. The sourcecount is defined as S . dN / dS / A . S is the source flux density, N is the source number and A is the observed sky area insteradians. Our source count is restricted to the 1975 and1173 sources identified by PyBDSF in two fields. The skyarea, A , is different for each bin of flux density, S . For fluxdensity S , we used the sky area for which the rms was lessthan / S , because only sources with a peak flux densitygreater than 4 σ are identified by PyBDSF. In both fields,the rms increases with the distance from the center becauseof the primary beam correction. So the region which is sensi-tive enough to detect faint sources is smaller than the regionto detect bright sources. The uncertainties in source countswere derived from Poisson statistics. Our results are consis-tent with other studies between 0.1 mJy and 110 mJy. Forlower flux density, the source counts are lower than Smolˇci´cet al. (2017). However, we need to keep in mind that thecompleteness and bias corrections used by Smolˇci´c et al.(2017) have not been applied to our data. The correctioncould be as large as a factor of 3 for a source with a fluxdensity of 5 σ (Table 4 and Fig. 18 of Smolˇci´c et al. 2017).At the same time, some inconsistency could be caused by dif-ferent software packages (e.g. PyBDSF in our work, BLOB-CAT in Smolˇci´c et al. (2017); Butler et al. (2018) and SADin Miller et al. (2009); Wold et al. (2012)) and thresholds(source peaks greater than S / N > or S / N > ) used toidentify sources in these studies. MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies Coma
GMP 4629GMP 4570GMP 2923GMP 3071IC 4040NGC 4921 GMP 3016 IC 3949NGC 4911 KUG1257+278KUG1258+279A MRK 0058NGC 4858 KUG 1255+283NGC 4848D100
NGC4869 Cont
Figure A1.
Left: Radio continuum image of two fields in the Coma cluster (30% response shown here). Galaxies shown in Fig. 1 and alsocovered by 30% response are marked as solid squares. Right: zoom-in radio continuum images of two interesting sources: the narrow-angletailed galaxy NGC 4869, and an interesting source at 12:57:20.54 +28:27:31.28 (J2000.0) marked as a dashed square.
Table B1.
The full continuum source catalogID (1) R.A. (J2000) (2) Decl. (J2000) (3) S int (4) ∆ S int (5) field (6)[h m s] [ ◦ (cid:48) (cid:48)(cid:48) ] [ µ Jy ] [ µ Jy ] Note:
Column descriptions: (1) source ID; (2) and (3) right ascension and declination (J2000) of radio source peak; (4) and (5) integralflux density and its associated error at 1.4 GHz; (6) The field (shown in Fig. 1 and Table 1) in which source was found. For the 64sources covered by both fields, their results from the two fields are listed in different rows respectively and the same source ID wasset for the same source.(This table is available in its entirety in machine-readable form online. A portion is shown here for guidance regarding its form andcontent.)
APPENDIX D: RADIO EMISSION FROMULTRA-DIFFUSE GALAXIES IN THE COMACLUSTER
We also examined the radio emission from ultra-diffusegalaxies (UDGs) in the Coma cluster. With the radio sourcesdetected in Miller et al. (2009), Struble (2018) found no evi-dence for radio emission from UDGs. We match the 3084radio sources detected in our deep observations and theUDGs in the Coma cluster (Koda et al. 2015; Yagi et al.2016). There are 449 UDGs within our fields (field 1 + 2, out to the 10% primary beam response level) that cover4152.9 arcmin . We find that there is 1 match within a 2 (cid:48)(cid:48) radius, 6 matches within a 5 (cid:48)(cid:48) radius and 26 matches withina 10 (cid:48)(cid:48) radius. Following Struble (2018), the expected num-ber of random matches within an offset of r is ∼ r < (cid:48)(cid:48) , ∼ r < (cid:48)(cid:48) , ∼ r < (cid:48)(cid:48) . Thus, even withour much deeper radio data, there is no evidence of radioemission from UDGs in the Coma cluster. MNRAS000
We also examined the radio emission from ultra-diffusegalaxies (UDGs) in the Coma cluster. With the radio sourcesdetected in Miller et al. (2009), Struble (2018) found no evi-dence for radio emission from UDGs. We match the 3084radio sources detected in our deep observations and theUDGs in the Coma cluster (Koda et al. 2015; Yagi et al.2016). There are 449 UDGs within our fields (field 1 + 2, out to the 10% primary beam response level) that cover4152.9 arcmin . We find that there is 1 match within a 2 (cid:48)(cid:48) radius, 6 matches within a 5 (cid:48)(cid:48) radius and 26 matches withina 10 (cid:48)(cid:48) radius. Following Struble (2018), the expected num-ber of random matches within an offset of r is ∼ r < (cid:48)(cid:48) , ∼ r < (cid:48)(cid:48) , ∼ r < (cid:48)(cid:48) . Thus, even withour much deeper radio data, there is no evidence of radioemission from UDGs in the Coma cluster. MNRAS000 , 1–20 (2020) H. Chen et al.
Right Ascension (J2000) D ec li n a t i o n ( J ) Figure B1.
Comparison of sources detected with PyBDSF between this work and Miller et al. (2009). Red apertures show the radiocontinuum sources defined with PyBDSF. Blue circles show the sources detected by Miller et al. (2009) (the radii are fixed at 20 (cid:48)(cid:48) ).This paper has been typeset from a TEX/L A TEX file prepared bythe author. MNRAS , 1–20 (2020) he ram pressure stripped radio tails of galaxies M ill e r + . GH z f l ux d e n s it y [ m J y ] Sources in NGC4848 fieldSources in D100 field F I R S T . GH z f l ux d e n s it y [ m J y ] Sources in NGC4848 fieldSources in D100 field M ill e r + . GH z f l ux d e n s it y [ m J y ] Figure B2.
Top two plots show the flux density comparison be-tween this work and Miller et al. (2009) and the FIRST survey(White et al. 1997). The bottom plot shows the comparison be-tween Miller et al. (2009) and PyBDSF results on their data.Dotted lines show an uncertainty of 0.3 index, while dashed andsolid lines show the linear and power-law fitting results. [mJy]0.11.010.0100.01000.0 S . d N / d S [ J y . s r − ] Coma NGC4848 field, JVLA 1.4GHz (this paper)Coma D100 field, JVLA 1.4GHz (this paper)Coma, VLA 1.4GHz (Miller et al. 2009)A370 VLA 1.4GHz (Wold et al. 2012)A2390 VLA 1.4GHz (Wold et al. 2012)COSMOS, JVLA 3GHz−>1.4GHz (Smolcic al. 2018)XXL South field, ATCA 2.1GHz−>1.4GHz (Bulter al. 2018)
Figure C1.
Comparison of source counts from our work and pre-vious results. For comparison, 3 GHz and 2.1 GHz flux densitiesfrom Smolˇci´c et al. (2017) and Butler et al. (2018) are convertedto 1.4 GHz flux densities assuming a spectral index of -0.7 and-0.75, respectively.MNRAS000