Molecular line emission in NGC 4945, imaged with ALMA

Molecular line emission in NGC 4945, imaged with ALMA

C. Henkel Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany Astronomy Department, Faculty of Science, King Abdulaziz University, P.O. Box 80203, Jeddah 21589, Saudi Arabia Xinjiang Astronomical Observatory, Chinese Academy of Sciences, 830011 Urumqi, China Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    S. Mühle Argelander Institut für Astronomie, Universität Bonn, Auf dem Hügel 71, 53121 Bonn, Germany Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    G. Bendo Jodrell Bank Centre for Astrophysics, University of Manchester, Oxford Road, Manchester M13 9PL, UK UK ALMA Regional Centre Node, University of Manchester, Oxford Road, Manchester M13 9PL, UK Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    G. I. G. Józsa Argelander Institut für Astronomie, Universität Bonn, Auf dem Hügel 71, 53121 Bonn, Germany SARAO, SKA South Africa, The Park, Park Road, Pinelands 7405, South Africa Rhodes University, RARG, RATT, P.O.Box 94, Grahamstown 6140, South Africa Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    Y. Gong Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany Purple Mountain Observatory & Key Laboratory for Radio Astronomy, Chinese Academy of Sciences, 2 West Beijing Road, 210008 Nanjing, PR China Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    S. Viti Department of Physics and Astronomy, UCL, Gower St., London, WC1E 6BT, UK Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    S. Aalto Dept. of Earth and Space Sciences, Chalmers University of Technology, Onsala Observatory, 43992 Onsala, Sweden Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    F. Combes LERMA, Observatoire de Paris, College de France, CNRS, PSL Univ., UPMC, Sorbonne Univ., Paris, France Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    S. García-Burillo Observatorio de Madrid, OAN-IGN, Alfonso XII, 3, 28014-Madrid, Spain Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    L. K. Hunt INAF-Osservatorio Astrofisico di Arcetri, Largo E. Fermi, 5, 50125, Firenze, Italy Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    J. Mangum National Radio Astronomy Observatory, 520 Edgemont Road, Charlottesville, VA 22903, USA Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    S. Martín European Southern Observatory, Alonso de Córdova 3107, Vitacura Casilla 763 0355, Santiago, Chile Joint ALMA Observatory, Alonso de Córdova 3107, Vitacura Casilla 763 0355, Santiago, Chile Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    S. Muller Dept. of Earth and Space Sciences, Chalmers University of Technology, Onsala Observatory, 43992 Onsala, Sweden Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    J. Ott National Radio Astronomy Observatory, P.O. Box O, 1003, Lopezville Road, Socorro, NM 87801-0387. USA Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    P. van der Werf National Radio Astronomy Observatory, P.O. Box O, 1003, Lopezville Road, Socorro, NM 87801-0387. USA Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    A. A. Malawi Astronomy Department, Faculty of Science, King Abdulaziz University, P.O. Box 80203, Jeddah 21589, Saudi Arabia Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    H. Ismail Astronomy Department, Faculty of Science, King Abdulaziz University, P.O. Box 80203, Jeddah 21589, Saudi Arabia Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    E. Alkhuja Astronomy Department, Faculty of Science, King Abdulaziz University, P.O. Box 80203, Jeddah 21589, Saudi Arabia Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    H. M. Asiri Astronomy Department, Faculty of Science, King Abdulaziz University, P.O. Box 80203, Jeddah 21589, Saudi Arabia Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    R. Aladro Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    F. Alves Argelander Institut für Astronomie, Universität Bonn, Auf dem Hügel 71, 53121 Bonn, Germany Max-Planck-Institut für Extraterrestrische Physik, Giessenbachstraße 1, 85748 Garching, Germany    Y. Ao National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, 181-8588, Tokyo, Japan    W. A. Baan Netherlands Institue for Radioastronomy ASTRON, 7991 PD Dwingeloo, The Netherlands    F. Costagliola Dept. of Earth and Space Sciences, Chalmers University of Technology, Onsala Observatory, 43992 Onsala, Sweden Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    G. Fuller Jodrell Bank Centre for Astrophysics, University of Manchester, Oxford Road, Manchester M13 9PL, UK Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    J. Greene Dept. of Astrophysical Sciences, Princeton University, Princeton, NJ 08544, USA    C. M. V. Impellizzeri European Southern Observatory, Alonso de Córdova 3107, Vitacura Casilla 763 0355, Santiago, Chile Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    F. Kamali Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    R. S. Klessen Universität Heidelberg, Zentrum für Astronomie, Inst. für Theoretische Astrophysik, Albert-Ueberle-Str. 2 and Interdisziplinäres Zentrum für Wissenschaftliches Rechnen, Im Neuenheimer Feld 205, 69120 Heidelberg, Germany    R. Mauersberger Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    X. D. Tang Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany Xinjiang Astronomical Observatory, Chinese Academy of Sciences, 830011 Urumqi, China Key Laboratory of Radio Astronomy, Chinese Academy of Sciences, 830011 Urumqi, China    K. Tristram European Southern Observatory, Alonso de Córdova 3107, Vitacura Casilla 763 0355, Santiago, Chile Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    M. Wang Purple Mountain Observatory & Key Laboratory for Radio Astronomy, Chinese Academy of Sciences, 2 West Beijing Road, 210008 Nanjing, PR China Leiden Observatory, Leiden University, P.O. Box 9513, 2300 RA Leiden, The Netherlands    J. S. Zhang Center for Astrophysics, Guangzhou University, Guangzhou 510006, China chenkel@mpifr-bonn.mpg.de
Received date ; accepted date
Key Words.:
galaxies: starburst – galaxies: structure – galaxies: ISM – Nuclear reactions, nucleosynthesis, abundances – galaxies: individual objects: NGC 4945 – Radio lines: ISM
offprints: C. Henkel,

NGC 4945 is one of the nearest (3.8 Mpc; 1 19 pc) starburst galaxies. To investigate structure, dynamics, and composition of its dense nuclear gas, ALMA band 3 (3–4 mm) observations were carried out with 2 resolution. Measured were three HCN and two HCO isotopologues, CS, CH, SiO, HCO, and CHCH. Spectral line imaging demonstrates the presence of a rotating nuclear disk of projected size 102 reaching out to a galactocentric radius of 100 pc with position angle PA = 452, inclination = 752 and an unresolved bright central core of size 2. The continuum source, representing mostly free-free radiation from star forming regions, is more compact than the nuclear disk by a linear factor of two but shows the same position angle and is centered 039 014 northeast of the nuclear accretion disk defined by HO maser emission. Near the systemic velocity but outside the nuclear disk, both HCN =10 and CS =21 delineate molecular arms of length 15 (285 pc) on opposite sides of the dynamical center. These are connected by a (deprojected) 0.6 kpc sized molecular bridge, likely a dense gaseous bar seen almost ends-on, shifting gas from the front and back side into the nuclear disk. Modeling this nuclear disk located farther inside (100 pc) with tilted rings provides a good fit by inferring a coplanar outflow reaching a characteristic deprojectd velocity of 50 km s. All our molecular lines, with the notable exception of CHCH, show significant absorption near the systemic velocity (571 km s), within a range of 500 – 660 km s. Apparently, only molecular transitions with low critical H-density ( 10 cm) do not show absorption. The velocity field of the nuclear disk, derived from CHCH, provides evidence for rigid rotation in the inner few arcseconds and a dynamical mass of = (2.10.2) 10 M inside a galactocentric radius of 245 (45 pc), with a significantly flattened rotation curve farther out. Velocity integrated line intensity maps with most pronounced absorption show molecular peak positions up to 15 (30 pc) southwest of the continuum peak, presumably due to absorption, which appears to be most severe slightly northeast of the nuclear maser disk. A nitrogen isotope ratio of N/N 200–450 is estimated. This range of values is much higher then previously reported on a tentative basis. Therefore, with N being less abundant than expected, the question for strong N enrichment by massive star ejecta in starbursts still remains to be settled.

1 Introduction

Observing the optically highly obscured central regions of galaxies hosting an active supermassive nuclear engine is fundamental for our understanding of galaxy evolution. The transfer of mass into sub-parsec scale accretion disks, the feeding of the nuclear engine and its feedback, potentially affecting the large scale appearance of the parent galaxy, are basic phenomena to be studied. The nearby, almost edge-on HO megamaser galaxy NGC 4945 is such an active galaxy. Its central region is known to show a rich molecular spectrum hosting not only a nuclear starburst but also an active galactic nucleus (AGN; see, e.g., Marconi et al. 2000; Yaqoob 2012). Past molecular single-dish studies include Henkel et al. (1990, 1994), Dahlem et al. (1993), Mauersberger et al. (1996), Curran et al. (2001), Wang et al. (2004), Hitschfeld et al. (2008), Pérez-Beaupuits et al. (2011), and Monje et al. (2014). As an outstanding galaxy in the Cen A/M 83 group at a distance of only 3.8 Mpc (e.g., Karachentsev et al. 2007; Mould & Sakai 2008; 1 corresponds to 19 pc), NGC 4945 hosts one of the three brightest IRAS (Infrared Astronomical Satellite) point sources beyond the Magellanic Clouds (Wang et al. 2004).

Although this Seyfert 2 galaxy has an infrared luminosity a few times that of the Milky Way, it nevertheless suffers from a lack of interferometric studies of molecular high density tracers, a consequence of its southern location. Greenhill et al. (1997) used the 22 GHz HO maser lines to map the circumnuclear accretion disk at a galactocentric radius of 0.45 pc (0025), which shows a similar position angle as the large scale disk, revealing a binding mass of order = (1–2) 10 M. More recently, 183 and 321 GHz HO and even Class I methanol maser emission has also been detected (Hagiwara et al. 2016; Humphreys et al. 2016; Pesce et al. 2016; McCarthy et al. 2017). Cunningham & Whiteoak (2005) mapped the 3 mm lines of HCN (partially), HNC, and HCO with an angular resolution of 56 35 and found an elongated morphology, with a position angle consistent with the inner HO and the outer large scale disk. If this gas were to trace a rotating disk, its radius would be 57 pc, with a rotational velocity of 135 km s, enclosing a mass of 2.510 M. Also based on ATCA (Australia Telescope Compact Array) data, more recently Green et al. (2016) published HCN, HCO, and HNC spectra, maps of integrated intensity, and the velocity field with 7 resolution.

A basic concern for the present study is the morphology of the dense molecular gas. In the following, we try to improve our understanding of the interplay between the active galactic nucleus, the surrounding dense gas environment, and its connection to clouds at larger galactocentric radii. After providing a comprehensive view on measured spatial and kinematical components we will look (1) for the presence of a molecular hole or the opposite, a molecular hotspot, at the very center, (2) for the presence of a rotating molecular disk, (3) for the existence of a molecular bar and (4) for the development of spiral patterns. Radial motions, frequently observed in other galaxies (e.g., Turner 1985; Feruglio et al. 2010; Bolatto et al. 2013; García-Burillo et al. 2014) are another phenomenon to be adressed.

Figure 1: An image of the 3 mm continuum emission from the central region of NGC 4945 (1 coresponds to a projected linear scale of 19 pc). The map is based on the line-free sky frequencies of 88.27–88.38 GHz and 88.54–88.70 GHz in the HCN window, 85.60–86.01 GHz and 86.28–86.42 GHz in the HCN window, 97.35–97.72 GHz and 98.01–98.24 GHz in the CS window, and 100.15–100.46 GHz in the continuum window (see Sect. 2). The beam size is 255 134 (0.1 Jy beam is equivalent to 6.8 K) with a position angle of 44.2. Contour levels are 0.5, 1, 2, 5, 10, 20, 40, 60, and 80% of the peak flux density of 96 mJy beam. The restored beam is shown in the lower left corner. In view of the limited extent of the emission, no primary beam correction has been applied.
Line SEST ALMA SEST ALMA Comments to the ALMA profiles
(GHz) (mJy) Jy km s
HCO =10 85.16222 —/— 20/20 2.70.4 weak, with absorption
c-CH 2–1 85.33889 180/130 130/130 606 21.40.7 strong, absorption NE of nucleus
CHCH 5–4 85.45727 —/— 40/40 7.50.5 strong, only emission
c-CH 4–4 85.65642 —/— weak, blended, see Sect. 4.3.1
H42 85.68840 —/— 40/40 12.90.8 strong, only emission
HCN =10 86.05497 20? 20 10? 1.10.2 weak
HCN =10 86.33992 65/75 50/60 212 9.40.8 strong, with absorption
HCO 10 86.67076 —/— 30/30 weak, blended
HCO =10 86.75429 —/— 40/40 strong, with absorption, blended
SiO =21 86.84696 —/— 40/40 strong, with absorption, blended
HCN =10 88.63116 1300/1100 1300/1000 4959 249.33.0 strong, with absorption
CS =21 97.98095 880/770 650/500 37520 112.41.2 strong, with absorption

a) For the integrated flux density and other information on the H42 line, see Bendo et al. (2016).
b) Approximate flux densities for the blue- and redshifted line peaks, obtained from the 15-m Swedish-ESO Submillimeter Telescope (SEST; Wang et al. 2004) and ALMA. Uncertainties in the ALMA data are 5–10 mJy from the noise level and 10% of the flux density from the calibration, while differences in the chosen area (10 arcsec) do not lead to significant changes. For the SEST, rms uncertainties are of order 40 mJy (c-CH), 12 mJy (HCN/HCN), 150 mJy (HCN), and 100 mJy (CS). As indicated by a question mark, the SEST detected HCN only tentatively. HCO provides a quartet of hyperfine components among which the strongest feature is seen by ALMA.
c) From Gaussian fits (see Wang et al. 2004 for the SEST data).
d) For the weak lines, it is not clear whether absorption is not observed because of a too low signal-to-noise ratio or whether it is truly absent. In the case of HCN, with two more abundant isotopologues showing weak absorption in a specific velocity interval and with HCN not being seen within exactly this velocity range, absorption is likely present.

Table 1: Detected spectral features toward NGC 4945

While trying to bridge the gap between the well known large scale structure (e.g., Dahlem et al. 1993; Ott et al. 2001) and the sub-parsec maser disk, we will also compare observed morphologies with that of the simultaneously measured radio continuum, mainly representing free-free emission from star forming regions (Bendo et al. 2016). Furthermore, we will analyze our images in the light of data taken at other wavelengths and briefly compare encountered structures with those from other (carefully selected) prominent nearby starburst galaxies. A superficial first look into chemical properties is also included. Finally, a relatively simple kinematical model of a particularly well ordered component of the dense molecular gas is presented.

Initially, the origin of the rare nitrogen isotope N was the major motivation to carry out the present study. Traditionally being assigned to low mass stars including novae, more recent models of stellar nucleosynthesis, describing the evolution of massive rapidly rotating stars, also permit the production of significant amounts of N through mixing of protons into helium burning shells (e.g., Woosley & Weaver 1995; Timmes et al. 1995). Direct confirmation of this scenario can be obtained by measuring a low (100) N/N ratio in a well developed starburst galaxy, where enrichment of the interstellar medium (ISM) by ejecta from massive stars must have occurred. Chin et al. (1999) tentatively reported such a ratio for NGC 4945, based on observations with the 15-m SEST (Swedish-ESO Submillimeter Telescope). Making use of the higher sensitivity and resolution provided by ALMA (Atacama Large Millimeter/Submillimeter Array) this result can now be checked.

The focus of this article is the molecular gas. Describing the observations in Sect. 2, overall distributions of molecular line emission (and absorption), including the continuum, are presented in Sect. 3. Sect. 4 provides the analysis, while main results are summarized in Sect. 5. The H42 radio recombination line, which also lies within the covered frequency range, has already been discussed by Bendo et al. (2016) in an accompanying paper.

Figure 2: Distribution of the main spatial HCN = 10 features with integrated HCN intensity as background. Assigned numbers are also presented in Col. 2 of Table 2, together with additional information. This includes the offset from position 9, the peak of the 3 mm continuum emission (see Fig. 1), which is responsible for the absorption features seen in most spectral lines. For the three quadratic areas surrounded by white dashed lines, see Sect. 3.3. 1 corresponds to a projected linear scale of 19 pc. The beam size (271 156, position angle 43) is given in the lower left.

2 Observations and data reduction

Band 3 (3–4 mm) images of NGC 4945 (project code: 2012.1.00912.S) were acquired in Cycle 1 with ALMA in a compact 32 antennae configuration on 2013 January 24. The minimal and maximal baseline lengths were 20 m and 420 m. Among the 32 antennas only one antenna, DV11, had to be flagged (bandpass). Four spectral windows were placed in Band 3, two in the lower sideband (LSB) and two in the upper sideband (USB) with the centers of the sidebands being separated by 12 GHz. The data were obtained using spectral averaging from

  • (1) an HCN (hereafter HCN) window centered at a sky frequency of 88.462948 GHz (rest frequency: 88.6316 GHz) with a bandwidth of 468.8 MHz (1588 km s) and 480 channels with a spacing of 3.3 km s,

  • (2) an HCN (hereafter HCN) window (this also includes HCN (HCN) and other lines) centered at 85.891221 GHz (86.05497 GHz) with a bandwidth of 1875 MHz (6545 km s) and 1920 channels with a spacing of 3.4 km s,

  • (3) a CS (hereafter CS) window centered at 97.794508 GHz (97.980950 GHz) with a bandwidth of 937.5 GHz (2874 km s) and 960 channels with a spacing of 3.0 km s, and

  • (4) a continuum window (see also Sect. 3.1) at 99.809726 GHz (100 GHz) with a bandwidth of 1875 MHz. Here the spacing of each of the 128 channels is 44.0 km s.

The systemic velocity is only known to an accuracy of about 25 km s (e.g., Chou et al. 2007; Bendo et al. 2016). Here we adopt a barycentric velocity of = c = 571 km s (optical convention; see, e.g., Dahlem et al. 1993). The differences between radio and optical velocities, = 1.1 km s, and between the barycentric and Local Standard of Rest reference systems, = 4.6 km s, are too small to play an important role in the following analysis.

Amplitude calibration was obtained by observing the quasar J1107–448. A flux density of 1.4 Jy has been taken from the ALMA calibrator data base. Bandpass and phase calibrators were J1427–4206 and J1248–4559, respectively. While in the following relative, not absolute intensities are mainly addressed, we nevertheless estimate an absolute flux accuracy of 10% (e.g., Fomalont et al. 2014; Wilson et al. 2014).

The phase tracking center was = 13 05 2728 and = –49 28 044 ( = 13 02 3207 and = –49 12 010; compare with Greenhill et al. 1997). The full width to half power (FWHP) field of view is 60. To investigate basic mophological properties we used “Briggs” weighting with a robustness parameter of 0.0 (approximately half way between natural and uniform weighting) for a reasonable trade off between sensitivity and resolution. Full synthesized half power beam sizes are then 26 15, with position angles of 42. In those cases, where we intended to include emission from spatial scales as large as possible, i.e. for moment 1 and 2 as well as position-velocity plots (Sects. 4.6.1 and 4.6.2), Briggs weighting with a robustness parameter of +2.0, to come close to natural weighting, has been applied. This leads to almost circular full half power beam sizes of 21 to 23. According to the “ALMA Cycle 1 Technical Handbook” (their Fig. 6.1), the largest accessible angular scales are, at a 50% intensity level, close to 20.

The observations lasted almost an hour, with an on-source integration time of about 30 min. Five full cycles between phase calibrator and NGC 4945 were performed, integrating on the latter source for about 6 min per cycle. Sensitivities in the HCN- and HCN-related windows are 0.6 mJy beam in 25 km s wide channels and 0.7 mJy beam in 30 km s wide channels of the CS window. For the continuum emission near 100 GHz an rms noise level of 0.13 mJy beam has been derived in a 310 MHz wide window (100.15–100.46 GHz, see also Fig. 1). Two versions of image cubes were created. One included the continuum emission. The other was continuum subtracted based on a linear fit to the line-free regions given in the caption to Fig. 1. The data were reduced, calibrated and imaged with the Common Astronomy Software Application (CASA)111http://casa.nrao.edu, version 4.0 (see McMullin et al. 2007).

Velocity Source Comments
(km s) (arcsec)
350 1 –3 –7 Peak 1, becoming detectable at 315 km s; there strongest feature
2 –1 –1 Peak 2, at 350 km s strongest feature, slightly SW from the center
3 –8 –4 Peak 3, weak, southern extension at 360 km s
390 1,2,3 –1 –1 Peak 2 emission extends to the SW, Peaks 1 and 3 join the forming NE-SW ridge
420 2 –1 –1 Peak 2 dominant, the molecular ridge expands toward the SW
460 2 –1 –2 Peak 2 dominant, position angle of ridge changes from 40 to 60
4 –6 –5 Appearance of a secondary peak
480 2 –1 –2 Peak 2 (inner SW ridge) gets weaker, appearance of weak extended arms
500 2 –1 –2 Peak 2 not dominant, appearance of an extended arm in the NW
515 –1 –2 Absorption, surrounded by emission in the SE, NW, and SW
5 +6 –8 Weak extended arms in the NW; given offset marks a hotspot in the SE
530–570 The morphology of the core emission resembles a at position angle 40
575 6 –2 –3 Absorption SW of the dynamical center, appearance of an extended arm in the SE
600 6 –2 –3 Absorption SW of the dynamical center, widespread weak line emission
7 +2 +2 A second region exhibiting absorption, located in the NE
8 –1 +2 Prominent (at this velocity) but not a particularly strong emission peak
5 +6 +8 Prominent (at this velocity) but not a particularly strong emission peak
620–665 9 0 0 Absorption dominates, morphology similar to that of the continuum emission
(see Fig. 1)
Widespread weak line emission, strongest on the NW side of the absorbing region,
from there extensions to the NE; also extended emission SE of the center
5 +6 +8 The SE hotspot (see 515 km s) fades
650–675 Absorption fades
665 10 +1 0 Appearance of an emission component slightly NE of the center
665–700 10 +1 0 While the position of the line peak is kept, the emission spreads toward the NE
The extended arms disappear
710 11 +3 +2 At this position emission becomes most intense, at higher velocities
line intensities start to decrease, disappearance of the extended arms
in the NW and SE
745 0 0 SW edge of the ridge of line emission
12 +4 +3 NE edge of the ridge of line emission
770 9 0 0 Weak line emission peak
13 +5 +9 Weak line emission peak

a) The table presents the positions of the main HCN = 10 hotspots. The numbers of the molecular hotspots given in Col. 2 can be visualized in Fig. 2 (white crosses). In the table they follow, from top to bottom, the barycentic velocity axis of the data cube. Therefore the table is best read from top to bottom. Position offsets given in Cols. 3 and 4 are relative to our 3 mm continuum peak at = 13 05 275, = –49 28 05 (Sect. 3.1). 1 corresponds to 19 pc (Sect. 1).

Table 2: HCN =0 line features toward NGC 4945 (see also Fig. 2)

3 Results

3.1 Continuum emission and the position of the HO maser disk

Figure 1 shows the 3 mm nuclear continuum emission of NGC 4945 obtained from the line free parts of the four measured bands (see also Sect. 2). Most of it (84% 10%) originates from free-free emission (see Bendo et al. 2016 for a detailed analysis, which also suggests that missing flux is not a problem here). This is the only detected source of continuum emission inside the area covered by the 1 sized primary beam. The integrated continuum flux density is 3101 mJy with a peak of 88.60.3 mJy beam. Note that the errors given are those of the fit and do not include the estimated calibration uncertainty (10%, see Sect. 2). The beam deconvolved size (full width to half maximum) is well described by an ellipse, centered at = 13 05 274930 and = –49 28 05229. Because of the high signal-to-noise ratio of several 100 (rms noise; 0.13 mJy beam, see Sect. 2) the nominal position error is only 0003. In view of changing weather conditions, the angle between phase calibrator and the target (5.4) and the level of errors in the antenna position measurements, we cautiously estimate that the absolute position uncertainty (compare, e.g., with Chou et al. 2007 and Lenc & Tingay 2009) is of order 01 (see also the ALMA knowledgebase in https://help.almascience.org/snoopi/index.php?, which provides (see the “Recent Articles” column, question 8), an absolute rms position uncertainty of 004).

Figure 3: HCN (HCN) =10 channel maps with the crosses indicating (as in case of Figs. 7, 9, 27, 28, and 29) the position of the radio continuum peak. The spacing between tick-marks along the y-axis corresponds to 10. Central barycentric velocities in km s are given in the upper left corners of each map. The systemic velocity is near 571 km s (Sect. 2). The wedge provides intensities in units of Jy beam. Note the dark areas between 525 and 650 km s, where the line is seen in absorption. The large scale rotation of the highly inclined galaxy is such that emission at low velocities is mainly originating from the southwest of the center, while high velocity emission arises from the northeast. 0.1 Jy correspond to 4.0 K main beam brightness temperature. 1 represents a projected linear scale of 19 pc. The beam size is shown in the lower right corner of the figure.

Using the Very Long Baseline Array (VLBA) Greenhill et al. (1997) determined the position of the circumnuclear 22 GHz HO maser disk, obtaining = 13 02 3228 and = –49 12 019 for the core of the curved line with a radius of 0025, formed by the maser hotspots. In the J2000 system this accretion disk is then centered at = 13 05 2748 and = –49 28 056 and coincides within the limits of our resolution with the dynamical center (see Sect. 4.6.1). In the following the position of the maser disk is therefore taken as synonymous for the dynamical center. The offset to our continuum peak is () = (–013, –037), with estimated uncertainties of 01 in either measurement. This 2.8 discrepancy is right at the limit of our sensitivity, but supporting evidence for its reality is provided below, in Sects. 4.1.2 and 4.6.1.

Beam deconvolved angular dimensions of the continuum source in NGC 4945 are 583001 132001 (projected 110 pc 25 pc) along the major and minor axes with a position angle of 41.40.1 (compare with Bendo et al. 2016 and, for 13 cm and 1.3 mm, with Lenc & Tingay 2009 and Chou et al. 2007, respectively).

3.2 A brief summary of spectral line properties

In this Sect. 3, we provide information on spectral properties, referring exclusively to the ALMA data presented here. The basis for detected spatial features is provided by Fig. 2. Additional information related to this image and its positions marked by white crosses is displayed in Table 2, which also provides relevant radial velocities. For those mainly interested in the more general discussion of these data, also including results from other studies, we refer to Sect. 4.

As expected, no spectral lines were found in the continuum spectral window and there are also no detected features close to those of HCN and CS (cf. Sect. 2). Table 1 summarizes the spectral features encountered in the three measured frequency bands containing molecular lines and coarsely characterizes their strength and line shape (i.e., exclusively emission or also absorption). The next subsections discuss estimates of the recovered flux for individual spectral features and provide a detailed description of the HCN =10 morphology. Finally, a short characterization of other lines is also presented. As a general rule, molecular lines from the nuclear region of NGC 4945 tend to show relatively strong emission at blueshifted velocities in the (barycentric) velocity interval 400–475 km s southwest of the dynamical center, less emission near the systemic velocity ( 571 km s) and a second spectral peak of line emission near = 710 km s northeast of the dynamical center. Also worth noting in this context is that the front side of the galactic plane is crossing the minor axis slightly southeast of the nucleus (e.g., Marconi et al. 2000).

Figure 4: Maps of primary beam corrected velocity integrated intensity (moment 0), obtained with a robustness parameter of 0.0 (Sect. 2). Data reduced with a parameter of +2.0 (close to natural weighting) do not reveal any additional features. 1 corresponds to a projected linear scale of 19 pc. Upper left: HCN = 10 emission with a restored beam of 271 156 and position angle PA = 43 (see the red spot in the lower left corner of the image). Contour levels: 5, 10, 20, 40, 60, and 80% of the integrated peak flux density of 24.0 Jy km s beam. Upper right: HCN =10 with a restored beam of 272 157, PA = 42. Center, left: HCN 10 (only 400 km s) with 272 157, PA = 42. Center right: CS = 21. Contours: 5, 10, 20, 40, 60, and 80% of the integrated peak flux density of 14.1 Jy km s beam. Restored beam: 257 138, PA = 41. Lower left: CH 21 with 272 157, PA = 42.. Lower right: CHCH 54. Restored beam: 272 157, PA = 42. To convert Jy km s beam into K km s beam, multiply by 37 (HCN), 39 (HCN, HCN, CH, and CHCH), or 36 (CS).

3.2.1 Recovered flux density

A basic problem, when analyzing interferometric data is the amount of missing flux. Since our largest accessible angular scale is with 20 at a 50% intensity level (Sect. 2) rather large, the amount of missing flux from the nuclear region should be either negligible or small (see, e.g., the single-dish maps of Dahlem et al. 1993, Henkel et al. 1994, and Mauersberger et al. 1996). Table 1 (Cols. 3 and 4) provides peak flux densities derived from the 15-m SEST and those obtained by our measurements, for the blue- and red-shifted line peaks. It is apparent that at the peak velocities most of the flux has been collected by the ALMA observations. The only discrepancy, which may be significant, is that of CS, where the ALMA data suggest peak flux densities 30% lower than those obtained with the SEST. However, even here the discrepancy does not reach the 3 level.

Cols. 5 and 6 of Table 1 display our integrated intensities from 2-dimensional Gaussian fits. A comparison with SEST data obtained with a 1 arcmin beam size indicates that we have collected 30, 36, 44, and 50% of the total SEST emission for CS, CH, HCN, and HCN. Uncertainties in these percentages range from 15% to 35% of the given values, depending on the strength of the respective line.

While our data apparently represent most of the flux at velocities near the line peaks, only 30–50% of the total velocity integrated flux density has been recovered. This can be explained by the fact that the area covered by our ALMA data encompasses only 10% of the SEST beam. Emission, mainly arising at near systemic velocities from the outer parts of the galaxy, is not covered by our data. As a consequence, we will assume in the following, that our calibration is correct within the previously estimated 10% limits of uncertainty and that missing flux is not relevant within the central 10 (190 pc) discussed below. Line ratios, discussed in Sect. 4.3, may be even less affected as long as spatial distributions are similar.

Figure 5: Upper panel, black: HCN =10 spectrum with a channel spacing of 3.3 km s from the central 0.25 arcmin of NGC 4945. The covered area is limited by = 13 05 26 and 13 05 29 and by = –49 28 50 and –49 28 20. The ordinate displays unweighted average values from pixels of the entire area (moment –1) as a function of barycentric velocity in units of mJy beam. For spatially integrated flux densities, multiply by 20020. Lower panel, black: A CS =21 spectrum with a channel spacing of 3.0 km s from the same region. Ordinate and abscissa: same units as for HCN. For spatially integrated flux densities, multiply by 24024. Applied pixel size for both spectra: 11. Overlayed spectra (red) with lower velocity resolution and larger beam size (1), matched to the amplitude of our spectra, are taken from Wang et al. (2004).

3.2.2 The HCN =10 line

As already mentioned, HCN =10 is the strongest spectral feature observed by us. Therefore it is used here as a reference for all other molecular lines. Table 2 and Fig. 2 display some basic information on detected spectral components. Fig. 3 provides channel maps with a spacing of 25 km s, Fig. 4 (upper left) shows a map of integrated HCN =10 emission, and Fig. 5 displays line profiles.

The spectrum (Fig. 5) has been obtained from an area of size 30 30 (570 pc 570 pc). In comparison with the SEST spectra (Henkel et al. 1990, 1994; Curran et al. 2001; Wang et al. 2004), the spectral intensity dip near = 630 km s is more pronounced. This is readily explained by the large SEST beam, 1, and the confinement of the HCN absorption (see below) to a small region and does not indicate substantial amounts of missing flux near the systemic velocity in our data. The absorption component, discussed below, was already identified by Cunningham & Whiteoak (2005) and by Green et al. (2016), using the Australia Telescope Compact Array (ATCA) with beam sizes of 56 35 and 7 7, respectively.

Most of the emission follows a narrow ridge, hereafter termed the nuclear disk, near position angle PA +45 east of north, with line peaks detected at low velocities slightly to the southwest and at high velocities slightly to the northeast of the center. At the outermost velocities, 350 km s (spatial Peak 2 in Fig. 2 and Table 2) and 770 km s (spatial Peak 11 in Fig. 2 and Table 2), these peaks show an offset of 35 in right ascension and 35 in declination. The offset corresponds to 100 pc.

At more intermediate blue- and redshifted velocities with respect to the systemic one, the nearly edge-on viewed (Sect 4.5 and 4.6.2) nuclear disk becomes visible, directed from the center toward the southwest with its approaching and toward the northeast with its receding part. Combined, the morphology of the ridge emission resembles to some degree the mm-wave continuum distribution displayed in Fig. 1, but the line emission is with 10 (190 pc) more extended.

As already mentioned, the trough in the spectrum around = 630 km s (Figs. 3 and 5) is caused by absorption. Absorption is seen toward the center and slightly northeast over the widest velocity range, 510 – 565 km s and 575 – 660 km s (Fig. 3). Toward the southwest, absorption is also seen, but only between 600 and 660 km s. Between 620 and 660 km s, both absorption components are spatially connected and reveal together a morphology that closely resembles that of the mm-wave continuum (Fig. 1). Overall, most of the absorption is seen at velocities redshifted with respect to the systemic one (see Figs. 3 and 6 and the discussion in Sect. 4.7). Absolute HCN = 10 absorption line intensities reach 50% of the peak continuum flux density (Sects. 3.1 and 3.3).

No absorption is seen beyond the extent of the central continuum source. However, emission is also observed well outside the ridge (Table 2 and Figs. 2 and 3). To these spectral features belongs Peak 1 at 400 km s, while Peak 3, also seen at low velocities, delineates the southwestern edge of the nuclear disk well displaced from the very center (for details, see Table 2). Already slightly below = 500 km s, weakly emitting clouds appear to form a secondary central ridge, approximately perpendicular to the position angle of the main body of the galaxy. At 500 – 700 km s, arms can be seen at the ends of this secondary ridge, 10 (projected 190 pc) from the center, directed toward the northeast from its northwestern edge, and toward the southwest from its southeastern edge (Fig. 3 and, for further discussion, Sect. 4.7).

Figure 6: Lower panel: HCN = 10 spectra with a channel spacing of 3.3 km s. Black: spectrum from the northeastern region (peak near = 710 km s) showing absorption. Red: spectrum from the southwestern region (peak near = 430 km s) showing absorption. The separation between the two regions is the line perpendicular to the nuclear disk crossing the peak of the radio continuum emission (see Figs. 2 and 3). Upper panel: The same for HCN = 10, but with a channel width of 24 km s. The ordinate provides average flux densities per beam over the covered area. The galaxy’s systemic velocity is close to = 571 km s

.

3.2.3 Rare HCN isotopologues

The HCN = 10 line (see Fig. 4, upper right, and Fig. 7) shows, with smaller amplitudes and lower signal-to-noise ratios, a similar behavior as HCN = 10. Similar to HCN, absorption is apparent in the channels between = 525 and 650 km s. Again, we find more absorption at velocities slightly redshifted relative to the systemic one, but the clear asymmetry apparent in the HCN = 10 line is here only seen in case of the redshifted northeastern part of the nuclear disk (Fig. 6 and Sect. 4.7 for further discussion). As for the main isotopologue, the absorption is confined to the extent of the nuclear continuum source, while the surrounding gas exhibits emission.

Fig. 4 and, in the Appendix, Fig. 27 show that we have also detected the N bearing isotopologue. While signal-to-noise ratios are low, HCN is detected in the low velocity channels and then again at 710 km s, displaced by a few arcseconds to the northeast, following the overall kinematics of the central region (Sect. 3.2.1). Therefore, HCN is seen in those parts of the spectrum where the more abundant isotopologues show relatively strong emission. In the velocity range, where HCN and HCN are affected by weak absorption, no HCN signal is seen, consistent with the overall weakness of HCN and HCN in this velocity interval. Fig. 8 also shows this effect. While the line is seen in the “total” and central spectrum and tentatively also in the southwestern one, the northeastern spectrum from the receding part of the galaxy does not show any significant peak. The spectra of the more abundant HCN isotopologues are contaminated by absorption mostly in this part of the spectrum and it is very likely that this also holds in the case of HCN.

Figure 7: HCN (HCN) = 10 channel maps. For the units of velocity, the ordinate, and the wedge as well as for the position of the crosses, see Fig. 3. As in case of the main isotopic species, absorption is present near the systemic velocities between 525 and 650 km s. 0.01 Jy correspond to 0.42 K. The beam size is shown in the lower right corner of the figure.
Figure 8: Spectra from the broad HCN window with seven contiguous channels smoothed, yielding a channel spacing of 23 km s. The molecular species related to detected spectral features are given in the two top panels and are described in more detail in Table 1. For H42, see Bendo et al. (2016). Because almost 2 GHz are covered, leading to a complex correlation between velocity and frequency, amplitudes are presented as a function of observed frequency. The lowest panel presents a spectrum from the central 0.25 arcmin, similar to the spectra shown in Fig. 5. The three upper panels show spectra from smaller areas, visualized by dashed white lines in Fig. 2, and include (from top to bottom) regions to the northeast of the center, close to the nucleus, and to the southwest of the center. The four numbers in parenthesis on the right hand side of each of these designations provide the eastern and western edge (in time seconds of Right Ascension) and the southern and northern edge (in arcseconds of Declination) of the chosen area (compare with the moment 0 maps of Figs. 2 and 4). The ordinate displays unweighted average flux densities per beam size for all pixels in the given areas, so that the values for the largest region (lowest panel), encompassing most areas with weak (or even absent) emission, are smallest. Multiply by 404 (upper three panels) and 25025 to obtain spatially integrated flux densities. Applied pixel size: 11.
Figure 9: CHCH 5 4 channel maps. For the units of velocity, ordinate, and the wedge as well as the position of the crosses, see Fig. 3. 0.01 Jy correspond to 0.42 K. The beam size is shown in the lower right corner of the figure.

3.2.4 The CS = 21 line

After HCN = 10, CS = 21 is the strongest molecular transition in our spectral windows. With respect to the velocity integrated HCN line, the integrated CS emission (Figs. 4, 5 and, for the channel maps, Fig. 28 in the Appendix) shows overall a similar spatial distribution (Fig. 4). However, being weaker than the HCN line and with all lines being measured with approximately the same sensitivity (Sect. 2), the emission outside the nuclear ridge is near the detection threshold. Therefore, the observed ridge appears to be narrower. A 2-dimensional Gaussian fit reveals for CS = 21 a beam deconvolved FWHP size of 88  07 19  01 at a PA = +432  11. Obviously, the emission from the nuclear ridge or disk is, as in the case of HCN = 10, more extended than that of the continuum (Sect. 3.1).

Because of the morphological similarity between the 3 mm radio continuum and the integrated CS = 21 line according to Figs. 1 and 4, we might expect that CS is even more affected by the continuum radiation than HCN and its weaker isotopologues. However, we find the opposite. While CS line intensities near the systemic velocity are also weaker than in the line wings near = 430 and 710 km s (Figs. 5 and 28), the dip near the systemic velocity only amounts to about one third of the peak flux density near 430 km s (and not to more than two thirds), indicating a much smaller influence of the underlying nuclear 3 mm radio continuum emission. Inspection of individual channels confirms this view.

3.3 Other lines

Cyclic CH (c-CH) 21 is our third strongest line. Its emission is displayed in Figs. 4, 8, and 29. No strong emission is seen outside the nuclear disk in the moment 0 map. However, in the channel maps of the Appendix, the outer arms are still weakly present. With respect to absorption, the line shows properties, which are intermediate between those of HCN =10 and CS =21. c-CH 21 does not exhibit absorption over the entire extent of the continuum source. We find instead weak absorption northeast of the center at those velocities, where HCN shows more widespread absorption. A comparison of the overall spectra (lower panels of Figs. 5 and 8) and individual channels also show that the deficit in emission (with respect to the line wings) at near systemic velocities (520 – 620 km s) is larger than that of CS but smaller than that of HCN.

There are still a few weaker molecular spectral features to be mentioned (see Fig. 8, where the top three panels (from the areas delineated by dashed white lines in Fig. 2) demonstrate the shift in radial velocity, when moving along the plane of the galaxy from the southwest to the northeast with rising velocities). As is apparent from the channel maps (Fig. 9), the CHCH 54 transition shows no indication for absorption. Unlike CS = 21, which still shows a dip in the overall emission spectrum near the systemic velocity (Fig. 5), the CHCH profile is flat-topped and a detailed inspection of individual channels shows no pixel with negative flux densities.

While HCO is too weak for a detailed analysis, both SiO =21 and HCO =10 (hereafter HCO) show strong absorption. Because of heavy line blending and because this will not play a major role in the discussion (Sect. 4), the following is not displayed in the form of figures but is only briefly noted: Selecting the most extreme pixels, for SiO and HCO isotopologues absolute flux density levels for absorption in 23 km s channels reach values similar to the emission peaks. For SiO, these levels are 9 mJy beam; for HCO we find 14 mJy beam. The latter is corroborated, although with lower signal-to-noise ratios, by the HCO = 10 line near 85 GHz (hereafter HCO; see Fig. 8), where the values become about 5 mJy beam. This is more extreme than in the cases of HCN and HCN = 10, where we find (again for the most extreme pixels) absorption levels down to –48 and –9 mJy beam, but peak emission at 235 and 14 mJy beam in 3.3 and 23 km s wide channels, respectively. For completeness, the corresponding values for CS = 21 (3.0 km s channels), c-CH 21 and CHCH 54 (23 km s channels) are 95 versus –2 mJy beam, 30 versus –2 mJy beam, and 9 mJy beam versus no pixel with significant negative flux density.

4 Discussion

4.1 Spatial morphology

4.1.1 The nuclear core

The molecular emission observed by us near the center originates primarily from a highly inclined rotating nuclear disk (Sects. 3.2 and 3.3). If it were a ring as suggested by Bergmann et al. (1992) or Wang et al. (2004), we should be able to define an inner radius. Like Bendo et al. (2016) for the H42 recombination line, however, we see no evidence for a central hole in our data, suggesting that any such hole in the disk has a size smaller than the synthesized beam, which measures 2.6” (50 pc) along the major axis of the disk. This is consistent with Chou et al. (2007), who could also not identify such a central void in their CO and CO = 21 maps. Note, however, that their data were obtained with a synthesized beam size of 51 28. Our measurements have higher resolution and may trace higher density gas.

Instead of a void we find a central region with strongly enhanced integrated emission. This region is not larger than our beam size and may still be unresolved. It encompasses only 40 pc, which agrees with the less tight upper size limit obtained by Chou et al. (2007), 55 pc, and is seen in all our maps of integrated intensity (Fig. 4). The cold X-ray reflector (200 pc 100 pc, with 50% of the X-ray emission arising from the innermost 30 pc) reported by Marinucci et al. (2012) agrees well with the sizes of our dense and dusty nuclear disk and the unresolved central core. The feature with a steep velocity gradient identified in CO = 21 by Lin et al. (2011; their fig. 4) may also arise from this core. Only the c-CH 21 line (Fig. 4) shows a peak which is more extended than our beam size along the minor axis of the galaxy, possibly because it is tracing more diffuse gas (e.g., Thaddeus et al. 1985) in spite of its rather high critical density (see Sect. 4.2 for details) or because it is enhanced by photon dominated regions (for the chemistry, see, e.g., Aladro et al. 2011) associated with the starburst ring proposed by Marconi et al. (2000; see below). The central 40 pc appear to host dense gas with particularly high column densities. While there is clearly molecular gas associated with the starburst ring of radius 25 (50 pc), proposed by Marconi et al. (2000) on the basis of near infrared Pa data, this putative ring is not associated with peaks of molecular emission (with c-CH as the possible exception; see Fig. 4). In this context we should note that the proposed Pa ring is actually a sequence of line peaks along the main axis of the galaxy with similar extent as the nuclear disk observed by us (see Marconi et al.’s 2000 fig. 2, upper and middle left panels).

For a better understanding of the central region, the radius of the sphere of influence of the nuclear engine is also of interest. With

(e.g., Barth 2003), the mass of the nuclear engine = (1–2) 10 M (Greenhill et al. 1997), distance = 3.8 Mpc (e.g., Karachentsev et al. 2007; Mould & Sakai 2008), and adopting = 120 km s as the stellar velocity dispersion (Oliva et al. 1995), we obtain 0016–0032 or 0.3–0.6 pc. Obviously, this scale is far smaller than our resolution and resembles that of the size of the HO megamaser disk (Greenhill et al. 1997).

Figure 10: Angular offsets (along the main axis with position angle 225 toward the southwest) of molecular versus continuum centroid position (obtained from two-dimensional Gaussian fits) as a function of depth of the spectral dip near the systemic velocity in percent of the line peaks at 430 km s. The percentages given refer to a covered area of 0.25 arcmin as described in Fig. 5. Molecular species are indicated. For details, also including line assignments, formal uncertainties and rest frequencies, see Sect. 4.1.2 and Table 1).

As already mentioned, the centrally peaked molecular line intensity distributions (Fig. 4) contradict the previously proposed molecular ring structure originally discussed by Bergman et al. (1992; see also fig. 6 in Wang et al. 2004) on the basis of single-dish CO/CO line intensity ratios as a function of velocity. The same also holds for the molecular ring with an inner radius of 200 pc, reported on the basis of single-dish CO =10 and 21 observations by Dahlem et al. (1993). What we clearly see in our denser gas tracers are peaks in molecular line emission near = 430 and 710 km s, potentially forming a circumnuclear ring at radius 5 (95 pc) which encloses the nuclear disk (see also Sect. 4.7). The nuclear disk itself extending from this radius toward the unresolved nuclear core provides a rich molecular environment, giving rise to strong emission from a multitude of molecular high density tracers.

4.1.2 Peak positions

Table 3 provides the coordinates of some relevant peak positions, i.e. those from the continuum and from the redshifted, systemic, and blueshifted emission of the CHCH line, which appears not to be affected by absorption. Table 4 presents the central positions of our unblended lines, obtained by two-dimensional Gaussian fits using the CASA-viewer. The scatter in these positions encompasses 007 in right ascension and 04 in declination and the weakest line, HCN, is not an outlier. We do not see large differences between the peak positions of the different lines. However, there appears to be a small offset relative to the peak of the continuum emission. For the seven transitions displayed in Table 4, we obtain an average offset of () = (–060 049, –060 033). Excluding HCN, average and sample standard deviations change insignificantly to (–068 049, –068 –032). While the standard deviations of the mean seem to imply that the offset is not significant, we note that all seven molecular peaks lie slightly to the south-west of the maximum of the continuum emission, apparently following the main axis of the projected disk. The molecular peak closest to the position of most intense continuum emission (02, toward the south; for formal position uncertainties see Table 4) is that of CHCH, showing no absorption (Sect. 3.2.5). Bendo et al. (2016) obtain 02 for their H42 line, which is also not showing absorption.

Feature , Accuracy
(2000)
665 km s 13 05 27.60/–49 28 03.8 02
3 mm continuum 13 05 27.49/–49 28 05.3 01
570 km s 13 05 27.47/–49 28 06.2 02
475 km s 13 05 27.18/–49 28 08.4 02

Approximate peak positions of the 3 mm continuum emission (see Sect. 3.1) and the three locations of maximal line width of the CHCH 54 transition, where the first column provides the associated barycentric radial velocity. The positions are close to the midline of the inclined nuclear disk of NGC 4945.

Table 3: Relevant CHCH and continuum coordinates

Keeping in mind that all data were obtained simultaneously so that relative positions should be very accurate (possibly more precise than the 01 mentioned in Sect. 3.1 for absolute positions), we may thus ask, whether there is a connection between the offset and the degree of absorption a line is showing. This correlation is displayed in Fig. 10. Since we do not know the depth of the absorption trough in the HCN line (footnote d of Table 1 and Sect. 3.2.3), this spectral feature has not been included. Errors in the depth of the dip near the systemic velocity may amount to 8% while systematic position errors are difficult to quantify. We therefore refrain from showing a fit to the data but conclude that there appears to be a change in central position with increasing absorption depth. As already mentioned, this drift points to the southwest. In view of this trend, shown by Fig. 10, the effect may be attributed to absorption predominantly occurring northeast of the dynamical center, shifting the fitted center of the molecular emission to positions southwest of the 3 mm continuum peak. This supports, in addition to the discussion of coordinates in Sect. 3.1 (and the moment 2 map of CHCH 54 provided and discussed in Sect. 4.6.1), our notion that the continuum peak is slightly displaced from the dynamical center of the galaxy.

Line Beam PA Source size PA Inclination Integrated
undeconvolved peak flux
() (deg) () (deg) (Jy km s) (Jy km s
beam)
HCO 2.721.57 42 13 05 27.44 –49 28 06.1 7.01.4 44 78 2.7 0.73
(1.30.3) (03) (03) (0.4) (0.09)
c-CH 2.721.57 42 13 05 27.44 –49 28 05.8 10.52.7 45 75 21.4 2.73
(0.40.1) (01) (01) (0.7) (0.08)
CHCH 2.721.57 42 13 05 27.49 –49 28 05.4 10.02.1 43 78 7.5 1.17
(0.80.1) (01) (02) (0.5) (0.07)
HCN 2.721.57 42 13 05 27.48 –49 28 05.5 6.01.3 31 77 1.1 0.36
(1.50.5) (05) (04) (0.2) (0.05)
HCN 2.721.57 42 13 05 27.35 –49 28 06.0 9.95.4 49 57 9.4 0.70
(0.80.4) (05) (02) (0.8) (0.05)
HCN 2.711.56 43 13 05 27.38 –49 28 06.3 11.45.9 47 59 249.3 14.78
(0.20.1) (01) (01) (3.0) (0.17)
CS 2.571.38 41 13 05 27.44 –49 28 05.7 10.33.3 47 71 112.4 10.47
(0.30.1) (01) (01) (1.2) (0.10)
HNC 5.593.54 15 13 05 27.48 –49 28 05.5 8.54.2 36.80
(0.50.5)
HCO 5.593.54 15 13 05 27.39 –49 28 06.9 14.44.8 33.00
(0.50.5)
3 mm 2.551.34 44 13 05 27.49 –49 28 05.3 5.81.3 41 76.9
01 (0.1) (0.1)
HO VLBA 13 05 27.48 –49 28 05.6 45

a) HCO, HCO, and SiO (see Table 1 and Fig. 8) are severely blended and therefore excluded from the list. Our remaining seven molecular lines are presented on top, following increasing frequency. Note that for most transitions, integrated intensities (; Col. 9) provide the difference between emission and absorption. Only CHCH appears to be entirely free of absorption. Formal accuracies of the peak positions in Cols. 4 and 5 are 04, 01, 02, 04, 03, 05, and 05 from top to bottom for our ALMA data with 2 resolution (see Col. 2). Numbers in parentheses indicate standard deviations of the 2-dimensional Gaussian fits to the original (not beam deconvolved data) and do not include the 10% calibration uncertainty given in Sect. 2. For comparison, we have also included parameters from the HNC and HCO data of Cunningham & Whiteoak (2005), from our 3 mm continuum data (see Sect. 3.1), and from the 22 GHz ( 1.3 cm) HO megamaser disk discovered by Greenhill et al. (1997). The uncertainty in the Cunningham & Whiteoak (2005) positions is 04; that of the megamaser disk is 01. The uv-coverage of the Cunningham & Whiteoak data is limited, leading for example to a position angle of 29 for the continuum. Therefore their PA values are not included and we also do not derive rotating disk inclinations from their data.

Table 4: Parameters of velocity integrated molecular lines

It is remarkable that HCN and HCN, which should show (aside from potential saturation effects in the main isotopologue) similar spatial distributions, are characterized by similar offsets. HCO, which is even more affected by absorption, has its peak at a lower offset with respect to the continuum. Whether this is due to a lower signal-to-noise ratio (the line is weak, see Table 1 and Fig. 8) or the consequence of a different spatial distribution remains an open question.

Table 4 also provides results from the 3 mm lines of HNC and HCO =10, taken from Cunningham & Whiteoak (2005). While the sensitivity and uv-coverage of these data do not match those of our ALMA images, we nevertheless note that their HCO profile, showing strong absorption, appears to peak southwest of the continuum peak, while HNC, less affected by absorption, appears to peak closer to the continuum position.

4.1.3 Beyond the nuclear disk

Off the central ridge (i.e. the nuclear disk), our map of integrated HCN = 10 emission (Figs. 2 and 4) shows a distinct secondary peak at = 13 05 281, = –49 28 13, (,) = (+6,–8) toward the southeast with respect to the continuum peak. The ratio of flux densities between the secondary and primary peaks is 0.150.03. Near this offset position, at (+3,–6), lies Knot B, following the nomenclature of Marconi et al. (2000, their fig. 2). Knot B is a region of enhanced near infrared Pa emission, related to star formation, and is located above NGC 4945’s galactic plane, toward the observer. Here we speculate that both objects are connected; Knot B represents merely the star forming front side of a giant molecular complex, while our data (not affected by obscuration) are revealing the entire cloud. In this context it is important to note that the region encompassing the Pa and the secondary molecular peak is not seen in H42 (Bendo et al. 2016). There the 3 H42 limit is 0.15 times its flux at the center of the galaxy. This may suggest that massive star formation is limited to the front side also accessible by near infrared spectroscopy.

The huge cavity, 5 (100 pc) in diameter, discovered by Marconi et al. (2000) a few arcseconds northwest of the center (see again their fig. 2 as well as Pérez-Beaupuits et al. 2011) might also be apparent in some of our channel maps. The northwestern part of the -shaped structure, encountered between barycentric 530 and 570 km s (see Table 2 and Fig. 3) may confine the bubble on its western, southern, and eastern sides. Knot C (Marconi et al. 2000) is not clearly seen in our molecular line maps.

For the connection between nuclear core, nuclear disk, and the outer arms mentioned in Sect. 3.2.2, see Sect. 4.7.

4.2 Why are some lines absorbed and others not?

There are several possible reasons why we measure some lines with signifiant amounts of absorption, while at least one of our molecular transitions (the CHCH 54 line) and the H42 recombination line (for this, see Bendo et al. 2016) appear to be seen entirely in emission. This may be caused by (1) different levels of line excitation, (2) different frequencies leading to changes in continuum flux density and morphology, (3) optical depth effects, (4) drastically different, directly detectable spatial distributions, (5) different chemical properties, or (6) a wide range of associated critical densities. In the following we briefly discuss these possibilities.

(1) Excitation: All molecular lines analyzed here connect levels up to 10 K above the ground state. This upper limit is likely much smaller than the kinetic temperature throughout the nuclear disk, so that no effect due to excitation is expected.

(2) Frequencies: The frequency range of our lines is too small to allow for a strong influence of variations in continuum morphology and intensity on the strength of the absorption features. Strongest absorption is seen in the HCO isotopologues and SiO. HCO =10 and SiO =21 are characterized by frequencies above that of CH 21 and below HCN =10 and CS =21, which are less affected by absorption. Considering HCO and HCN isotopologues, the amount of absorption appears to be molecule dependent and does not change significantly when considering different isotopologues at slightly different frequencies (Fig.10).

(3) Line saturation: Optical depth effects capable of reducing the effective critical density by a factor of in the optically thick case do not make a notable difference. This is most apparent in our data, when considering HCN, HCN, and HCN (Sect. 3.2.4). The same holds for HCO and HCO (our data), when a spectrum from the main species (HCO; fig. 6 in Cunningham & Whiteoak 2005) is added.

(4) and (5) Chemistry and morphology: Aspects related to spatial distribution and molecular chemistry have to be considered together, because the chemistry may strongly affect observed morphologies in different molecular lines, and similarly, different environments may exhibit different chemical compositions. According to Table 4, all tracers show a similar, elongated distribution with the main axis being consistent with the position angle of the nuclear disk and with a tendency for slightly smaller sizes for weaker lines. The latter is likely due to the noise level, which is (in absolute terms) similar for all lines (see Sect. 2). There is, however, one notable exception. While CS =21 is clearly stronger than c-CH 21 (compare Figs. 5 and 8), c-CH appears to be equally extended. This is consistent with the fact that c-CH is quite abundant in diffuse low density clouds (e.g., Thaddeus et al. 1985), while CS is more concentrated near sites of massive star formation (e.g., Mauersberger et al. 1989; Shirley et al. 2003). Possibly representing more diffuse gas and being more widespread than CS, c-CH should be less affected by the compact nuclear region with strongest continuum emission. Nevertheless, absorption in the c-CH 21 line is more pronounced than in CS = 21 (cf. Figs. 5 and 8).

(6) Critical densities: Using the Leiden Atomic and Molecular database (LAMBDA; Schöier et al. 2005; van der Tak et al. 2007) and for CHCH Askne et al. (1984), we can obtain critical densities (where rates of spontaneous radiative decay are matched by collisional de-excitation) for all observed molecular transitions. Accounting only for the Einstein- coefficient and the collision rate (downwards) related to the measured transition in the LAMBDA database, we obtain in order of increasing relevance of absorption for CHCH, CS, CH, HCN, SiO, and HCO density values of 110, 310, 110, 210, 210, and 210 cm, respectively. CS, for example, shows a critical density similar to that of HCO, but it is only HCO, which shows deep absorption, comparable to that of SiO (Sect. 3.3). Thus we do not see a dependence between critical density and absorption depth in most molecular lines. The inclusion of the HNC = 10 data by Cunningham & Whiteoak (2005) with 310 cm and a systemic velocity emission level of order 25% of the line peak does not make the problem more transparent. We note, however, that CHCH has a significantly lower critical density than all the other lines observed by us, like the CO = 21 line measured by Chou et al. (2007); in both cases, no absorption is apparent.

To summarize, notable absorption appears to be only found when critical densities are well in excess of 10 cm. This is likely due to different spatial distributions, which however cannot be assessed with the angular resolutions attained by Chou et al. (2007) and our study. With 84%10% of the continuum caused by bremsstrahlung (Bendo et al. 2016), the higher density tracers are likely concentrated close to the many individual continuum sources related to massive star formation, while those with low critical densities should emit from more extended regions. While we think that this is the most plausible explanation, there still may be a caveat. CHCH 54 line emission is weaker than that of HCN/HCO = 10 and CS = 21 emission and therefore its optical depth may be lower. More severe photon trapping in the stronger lines may reduce their effective critical densities below those given here (e.g., Shirley 2015). However, this is not supported by our low estimate of the HCN = 10 optical depth in Sect. 4.4. In any case, subarcsecond resolution is required to reveal the differences in the small scale distributions of the molecular tracers and the mm-wave continuum (for comparison, see, e.g., the high resolution 2.3 GHz continuum data presented by Lenc & Tingay 2009).

Species Peak flux
430 km s 710 km s
mJy/beam
HCN 235.0 200.0
CS 95.0 80.0
HCO 5.5 5.5
c-CH 27.0 30.0
CHCH 6.5 9.0
HCN 14.0 13.5
HCO 7.0 14.0
SiO 9.0 9.0

a) For the specific transitions and frequencies, see Table 1. Given flux densities may be accurate to 10% (see Sect. 1 and, for given values, also the end of Sect. 3.3). Flux density ratios should not be affected by calibration errors (estimated to be of order 10%; see Sect. 2) but only by noise. In Sect. 4.4 we therefore conservatively adopt for such ratios also a 10% uncertainty.

Table 5: Peak flux densities in single 1 sized pixels

4.3 Line intensity ratios

Single dish observations have already addressed a number of chemical fingerprints of the central region of NGC 4945 (e.g., Henkel et al. 1990, 1994; Curran et al. 2001; Wang et al. 2004; Monje et al. 2014). The second-to-last of these studies contains a particularly large collection of molecular tracers between 81 and 355 GHz. However, due to a lack of interferometric data these studies could not account for the influence of absorption onto overall line profiles and integrated intensities and, in addition, detailed comparisons with spectral properties of other galaxies were not yet possible more than a decade ago. Our study only covers a small frequency interval, but with superior sensitivity and angular resolution. And we can also refer to spectral line surveys of other galaxies. Here we focus on the line wings of NGC 4945 near barycentric = 430 and 710 km s, where absorption does not play a role.

With respect to the molecular species we have measured, the CS/HCN ratio has been claimed to be a possible starburst/AGN tracer (Meijerink et al. 2007; Izumi et al. 2013, 2016), while CHCH, c-CH, and HCO have been identified as tracers of UV fields in galaxies (e.g., García-Burillo et al. 2002; Martín et al. 2009; Aladro et al. 2015). SiO can highlight the presence of shocks (e.g., Gusdorf et al. 2008; Guillet et al. 2009; Duarte-Cabral et al. 2014).

In the following we compare line intensities (see Table 5) at the positions of strongest emission of the southwestern blue- and northeastern redshifted line peaks (note, that the values differ from those given in Table 1, where area integrated intensities are derived; taking flux density or temperature ratios is not critical in this context). Since a detailed discussion of chemical issues is well beyond the scope of this article, the following offers merely some highly superficial first insights into the chemical complexity of the gas in the nuclear disk of NGC 4945. Following fig. 17 of Meijerink et al. (2007), our (HCN)/(CS) line peak intensity ratios of 2.5 with an estimated uncertainty of 10% (see Table 6) are consistent with a photon dominated region (PDR) of density (H) 10 cm and 10 times the local interstellar radiation field as well as with an XDR of density 10 cm, but a more detailed discussion on physical and chemical boundary conditions is clearly still required.

Species NGC 4945 NGC 253
430 km s 710 km s
(HCN)/(CS) 2.50.3 2.50.3 2.00.2
(HCN)/(c-CH) 8.70.9 6.70.7 11.40.2
(HCN)/(HCN) 16.71.7 14.81.5 14.50.4
(c-CH/(HCO) 4.90.5 5.50.6 7.10.5
(HCN)/(SiO) 26.12.6 22.22.2 14.40.2
(CS)/(SiO) 10.61.1 8.90.9 7.10.1
(c-CH)/(CHCH) 4.20.4 3.30.3 2.30.1
(HCN)/(HCO) 2.00.2 1.00.1 1.20.1

a) For the specific 3 mm transitions and frequencies, see Table 1. For the NGC 4945 ratios, derived from Table 5, an uncertainty of 10% has been conservatively adopted. For NGC 253, see Table A.2 of Aladro et al. (2015).

Table 6: Integrated line intensity ratios

A comparison with the eight galaxies covered by the 3 mm line survey of Aladro et al. (2015) provides additional insights. Our 3 mm (HCN)/(CS) line intensity ratios are in line with seven of the eight targets studied by Aladro et al. (2015). Among these seven, NGC 253 is outstanding. It is the strongest molecular line emitter outside the Local Group and has a distance, an inclination and an infrared luminosity similar to that of NGC 4945 (see also Sect. 4.8). It is therefore selected as the representative one of the seven for Table 6. Only in M 51, the ratio is, at 5.20.1, significantly higher. M 51 is a Seyfert galaxy where the inner spiral arms were also part of the measurement. This adds complexity but the high ratio is likely related to the presence of a nuclear jet leading to shock enhanced HCN (Matsushita et al. 2015). Our (HCN)/(HCN) peak flux density ratios are, at 15, smaller than for most of the galaxies of the Aladro sample, indicating a relatively high optical depth. This makes sense because the data from Aladro et al. (2015) were taken with 25 resolution, while in NGC 4945 we consider exclusively those two compact 2 (40 pc) sized regions with highest unabsorbed intensity. Nevertheless, the ratios for the two ULIRGs in their study, Arp 220 and Mrk 231, are even smaller (4–7) in spite of particularly large C/C isotope ratios (Henkel et al. 2014), indicating that there the bulk of the molecular gas is assembled in a compact volume of very high opacity (e.g., Rangwala et al. 2011; Aalto et al. 2015a, b; Martín et al. 2016; Scoville et al. 2017).

The galaxy with a similar (HCN)/(HCN) ratio in the Aladro et al. (2015) sample is NGC 253, which is also supposed to have a similar C/C ratio (see Wang et al. 2004; Henkel et al. 2014) of 40 versus 40–50. Therefore opacity effects should not be significant when comparing the NGC 253 and NGC 4945 data sets. From Table 6 we find that ratios are qualitatively similar, but nevertheless there are some notable differences. Following Martín et al. (2006) the chemistry in the nuclear region of NGC 253 is dominated by low-velocity shocks, which can be traced by SiO. Our line ratios from NGC 4945 indicate that SiO is less prominent with respect to other tracers, so shocks may still play an important but slightly less dominant role. Surprisingly, the “ubiquitous” (Thaddeus et al. 1985) c-CH is less prominent in NGC 253 than in our NGC 4945 molecular hotspots. In view of the large region studied in NGC 253, this result appears to be counterintuitive. It can possibly be explained by massive star formation and a more pronounced dominance of PDRs in NGC 4945. Another interesting feature is the underabundance of CHCH with respect to c-CH in NGC 4945 (Table 6). If high temperatures support the formation of CHCH as suggested by Aladro et al. (2015), this may hint at higher temperatures in the central region of NGC 253 than in our two selected locations of NGC 4945. We conclude that the active galactic nucleus in NGC 4945 (NGC 253 is devoid of an AGN) has not a great influence. The greater dominance of PDRs in NGC 4945, a possibly higher in the nuclear region of NGC 253, and the small radius of the sphere of influence of the nuclear engine in NGC 4945 (Sect. 4.1.1) all suggest that the presence or absence of an AGN is not dominating the differences outlined above.

The (HCN)/(HCN) ratios at barycentric = 430 and 710 km s are, near 15, identical within uncertainties. If the gas is well mixed, this indicates similar opacities, in apparent disagreement with the strong gradient in (CO 21)/(CO 10) along the major axis, from 0.9 to 1.6, obtained by Dahlem et al. (1993, their fig. 6). CO 21/10 line ratios much larger than unity, as seen near 700 km s, are indicative of (1) highly excited gas with low CO opacity, (2) partially displaced CO = 10 and 21 emitting regions with different levels of excitation or (3) relative pointing errors between the CO =10 and 21 line data. Very high excitation is indicated by the detection of the 321 GHz line of HO near = 700 km s (but not at other velocities), about 1850 K above the ground state (Pesce et al. 2016). Similar (HCN)/(HCN) ratios and drastically different CO 21/10 ratios near barycentric = 400 and 710 km s (0.9 versus 1.6) would suggest in the absence of CO-pointing errors that unlike our data from high column density positions the single-dish CO data of Dahlem et al. (1993) include large amounts of more diffuse gas.

4.4 The nitrogen isotope ratio

As already mentioned in the introduction, Chin et al. (1999) reported a tentative detection of HCN =10 emission, suggesting an HCN/HCN ratio of order two. This has been interpreted as a hint for efficient stellar N ejection into the ISM of starbursts, due to rotationally induced mixing of protons into the helium-burning shells of young massive stars (e.g., Woosley & Weaver 1995; Timmes et al. 1995). NGC 4945, with its nuclear starburst, is an ideal target to test this scenario.

In view of the weakness of the HCN line, single pixel peak flux densities as displayed in Table 5 cannot be used. Covering instead the entire measured region, we find from Table 1 (Col. 6) and Table 4 (Col. 9) for the =10 lines velocity integrated intensities of (HCN)/(HCN) = 26.52.3, (HCN)/(HCN) = 22741, and (HCN)/(HCN) = 8.51.7 (the errors do not include calibration uncertainties; for these, see Sect. 2). For C/C 40–50 (Wang et al. 2004; Hitschfeld et al. 2008; for NGC 253, see Henkel et al. 2014), this would imply a plausible average HCN opacity of two and a N/N ratio of 400. The latter matches the value encountered in the local Galactic disk (Dahmen et al. 1995) and is much higher than that tentatively deduced by Chin et al. (1999).

Figure 11: The HCN = 10 moment 1 velocity field, clipped at the 5 level. Contours are 440, 520, 600, and 680 km s from lower right to upper left. Primary beam correction and Briggs weighting (robustness parameter +2.0) has been applied. The restored beam (252 217, PA = 78) is shown in the lower left corner of the image. The cross indicates the 3 mm radio continuum peak.
Figure 12: The CS = 21 moment 1 velocity field, clipped at the 5 level. Contours are 440, 520, 600, and 680 km s from lower right to upper left. Primary beam correction and Briggs weighting (robustness parameter +2.0) has been applied. The restored beam (222 199, PA = 78) is shown in the lower left corner of the image. The cross indicates the 3 mm radio continuum peak.

Nevertheless, in view of the possible influence of absorption and differences between the northeastern and southwestern edge of the inner disk (see, e.g., the huge difference in the CO = 21/10 ratio as a function of velocity (Dahlem et al. 1993, their fig. 6)), a closer inspection is warranted. Smoothing seven contiguous channels like in Fig. 8, we obtain peak flux densities of 13.9 (HCN-southwest), 13.5 (HCN-northeast), 3.5 (HCN-southwest), and 2.2 mJy beam (HCN-northeast). The latter is barely above the noise level and does not necessarily imply different abundance ratios in the southwest and northeast. Adding and dividing the respective flux densities, we obtain with this method an (HCN)/(HCN) ratio of order five. Again, this is well above the ratio deduced from the tentative detection of Chin et al. (1999) and would imply a N/N value similar or slightly below that of the solar system, 270 (e.g., Dahmen et al. 1995), where some enrichment due to massive star ejecta may have played a role (Chin et al. 1999). To summarize, the measured range of ratios, N/N 200 – 500, which we find in NGC 4945, is not in contradiction with the traditional view of N being mainly synthesized in low mass stars. More recently, this has been also proposed by Adande & Ziurys (2012), based on Galactic data (see also Ritchey et al. 2015 and Furuya 2017).

While our observations demonstrate that the ratio is 100, it remains to be seen whether it is 200 or 450. This will await ALMA measurements with higher sensitivity than those presented here with 30 minutes of on-source integration time. Furthermore, in subsequent studies C/N ratios should be quantified in more than one molecular species to account for possible fractionation effects as modeled by Roueff et al. (2015; see their fig. 3), even though they may not be significant because of the high kinetic temperatures likely characterizing the molecular gas in the nuclear disk of NGC 4945.

4.5 Inclinations

Assuming that the observed emission arises from a flat disk with azimuthal symmetry, the ratio of the lengths of the minor to the major axis of our 2-dimensional Gaussian fits provides a direct measure of the inclination. Such an estimate is not biased (as at optical or NIR wavelengths) by extinction due to dust. The resulting inclination angles are given in Col. 8 of Table 4. For the outer disk, = 783 (Ott et al. 2001), while for the nuclear region (hereafter “inner disk”) Chou et al. (2007) proposed = 622 on the basis of CO = 21 data. We also get for some lines inclinations below or even far below those of the large scale disk, i.e. 71 from CS =21 and values close to that of Chou et al. (2007) from HCN and HCN = 10. However, the lines also collect emission from outside the nuclear disk (Sects. 3.2.2, 4.1.3 and in particular 4.7) and thus broaden the extent of the emission along the minor axis, yielding unrealistically low inclinations. Our other four lines, HCO =10, c-CH 21, CHCH 54 and HCN =10, being too weak to trace much of the less intense emission outside the nuclear disk, yield instead = 75 – 78, i.e. values perfectly consistent with the inclination on larger scales. We conclude that the inner disk still retains the inclination of the outer one, while some molecular lines are contaminated by emission from regions outside the inner 102 sized region representing the nuclear disk. While this is an obvious explanation for CO (Chou et al. 2007), HCN, and CS, it is remarkable that the optically thinner HCN (but not HCN) is following HCN. We conclude that four of our seven analyzed transitions show the same inclination as the outer disk, which would not be the case if the inclinations were different. Furthermore, the radio continuum arising predominantly from bremsstrahlung related to star formation with associated molecular clouds (Bendo et al. 2016), provides with its extent given in Sect. 3.1 an inclination of = 76.90.1 (see also Sect. 4.6.2 for a modeled inclination). In this sense NGC 4945 is remarkable, because a correlation between nuclear and large scale disk inclination is absent in statistically relevant samples of active spirals (e.g., Ulvestad & Wilson 1984 and Sect. 4.7 below).

Figure 13: A map of the HCN = 10 moment 2 velocity dispersion, clipped at the 5 level. Primary beam correction and Briggs weighting (robustness parameter +2.0) has been applied. The restored beam (252 217, PA = 78) is shown in the lower left corner of the image. The cross indicates the 3 mm radio continuum peak, while the position of the secondary maximum in the south (also seen in Fig. 14) is inconspicuous in all other aspects.
Figure 14: A map of the CS = 21 moment 2 velocity dispersion, clipped at the 5 level. Primary beam correction and Briggs weighting (robustness parameter +2.0) has been applied. The restored beam (222 199, PA = 79) is shown in the lower left corner of the image. The cross indicates the 3 mm radio continuum peak, while the secondary peak in the south (also seen in Fig.13) is inconspicuous in all other aspects.

4.6 Kinematics

4.6.1 Observational constraints

Figs. 1114 provide images of the intensity () weighted line-of-sight velocity distributions ( = )

(moment 1) and the intensity weighted velocity dispersions

(moment 2) of the two strongest lines, those of HCN and CS. Emphasizing here not basic morphological features but trying to also account for as much large scale structure as possible, these figures present Briggs weighted images (weighting robustness parameter +2.0) to come close to natural weighting. While being affected by different degrees of absorption, the velocity fields of HCN and CS look quite similar and are at first order consistent with rotation of a disk around a central source. Furthermore, the systemic velocity contour passes, in spite of this absorption, within the limits of accuracy (1) through the center of the galaxy (see also Cunningham & Whiteoak 2005). However, our isovelocity contours are not quite as regular as those of the H42 recombination line. There, isophotes are approximately perpendicular to the main axis of the galaxy (PA –45; see Bendo et al. 2016), giving rise to the impression that the innermost gas is rotating like a rigid body. A comparison with the H91+H92 data of Roy et al. (2010, their fig. 2), obtained with an even smaller beam size (14 12 instead of the 20 to 25 used here) also reveals this difference. While in the innermost 4 (75 pc) their isovelocity contours are roughly parallel to each other and perpendicular to the position angle of the galaxy, their contours in the northeast (particularly the 680 km s contour) are displaced northwards from the main axis. This also holds in the case of CO = 21 (Chou et al. 2007). However, our HCN data (Fig. 11) show instead a twist to the south, while CS (Fig. 12) is neutral in this respect. While it might be that absorption is significantly affecting the observed velocity fields, the main cause of the differences may be emission from outside the nuclear disk (Sect. 4.7).

Figure 15: The CHCH 54 velocity field, clipped at the 5 level. Unlike in Figs. 11 and 12, here the emission is confined to the nuclear disk. Primary beam correction and Briggs weighting (robustness parameter +2.0) has been applied. The contours start at 470 km s and reach 695 km s; their spacing is 25 km s. The restored beam (265 228, PA = 80) is shown in the lower left corner of the image. The cross indicates the 3 mm radio continuum peak.
Figure 16: A map of the CHCH 54 velocity dispersion, clipped at the 5 level. Primary beam correction and Briggs weighting (robustness parameter +2.0) has been applied. The restored beam (265 228, PA = 80) is shown in the lower left corner of the image. The cross indicates the 3 mm radio continuum peak.

When considering the distribution of line widths, there is another notable feature: The region with largest dispersion (80 km s in Figs. 13 and 14) is, as expected, located at the center of the galaxy. Perhaps a part of this could be related to the non-thermal source with jet-like morphology detected by Lenc & Tingay (2009) at an offset of 1 from the dynamical center. However, a detailed inspection of the spectra indicates that the fits of the central region are questionable because of absorption near systemic velocities (see, e.g., Sect. 3.2.2 and Table 2). Toward the northwest and southeast slightly off the region exhibiting absorption (e.g., Fig. 3), wide line emission is seen, because here parts of the approaching and receding disk are seen simultaneously. Toward the absorbing region itself, velocities near the systemic ones are blocked, inhibiting meaningful fits of the line widths.

In view of the line contamination by absorption and, likely even more important, by gas arising from outside the nuclear disk (Sects. 3.2.2, 4.1.3, and 4.7), we thus prefer to choose the CHCH 54 line. While its emission is 35 times weaker than that of HCN =10, it is not significantly affected by absorption and traces solely the nuclear disk. Figs. 4 and 1516 show the distributions of velocity integrated intensity, velocity, and velocity dispersion exclusively for the nuclear disk. The velocity field turns out to be highly regular, similar to that derived for the H42 recombination line (Bendo et al. 2016). Twists to the north or south (e.g., Chou et al. 2007) are absent and the most blueshifted and redshifted velocities are found in the outermost regions of the nuclear disk along its major axis.

While the velocity field is quite regular, the velocity dispersion looks more complex. Instead of one peak extended along a southeast-northwest axis (Figs. 13, 14), we find three (Fig. 16). The central one, elongated along the main axis of the galaxy and not perpendicular to it, is accompanied by two additional regions with wider lines to the northeast and southwest, separated by a total of 5902. At these locations we observe the transition between rigid body rotation and steep rotation curve to a flat one (see also Chou et al. 2007). At a radius of 245 (45 pc) and with a rotation velocity of 140 km s we obtain with equation (1) of Mauersberger et al. (1996; = 1) an enclosed mass of = 2.1 10 M with an estimated error of 10%, in good agreement with Cunningham & Whiteoak (2005). For comparison, Roy et al. (2010) obtain with the H92 line 310 M for the mass inside a radius of 1 (19 pc). Furthermore, we note that the nuclear region is slightly lopsided: The center of the line connecting the two outer peaks of line width is located 14 southwest of the continuum peak (see Table 3). Systemic velocities ( 571 km s) are found about 09 (15–20 pc) southwest of the continuum peak and may be closer to the position of the maser disk (Greenhill et al. 1997) given in Sect. 3.1. While uncertainties with respect to the value of the systemic velocity (Sect. 2) and the relative positions of the continuum peak and the maser disk are significant, most of the star formation represented by the continuum emission (see Bendo et al 2016) appears to originate slightly northeast of the dynamical center. The lopsidedness of the central region of NGC 4945 also explains why molecular lines with significant absorption show integrated intensity peaks slightly shifted to the southwest: With the bulk of the continuum arising from the northeast, line emission is likely more quenched by absorption at this side of the center than in the southwestern part of the inclined nuclear disk (see Sect. 4.2 and Fig. 10).

Figure 17: Position-velocity plot of CHCH 54 along the major axis with the abscissa scaled relative to the continuum peak. Contour: 7 mJy beam. The plot emphasizes (a) the clumpiness of the gas, (b) the steepening of the rotation curve between barycentric 500 km s