Free Access
Issue
A&A
Volume 617, September 2018
Article Number A20
Number of page(s) 20
Section Extragalactic astronomy
DOI https://doi.org/10.1051/0004-6361/201833338
Published online 20 September 2018

© ESO 2018

1. Introduction

Tracers of dense molecular gas are good probes of the central regions of active galaxies, where molecular regions are subjected to strong radiation fields (X-rays, cosmic rays, and ultraviolet (UV) fields) created by massive star formation and/or active galactic nuclei (AGNs). In particular, the rotational transitions of HCN and HCO+ have bright emission and high dipole moments (hence large critical density) and are therefore convenient tracers of the dense gas component in galactic centres.

In the particular case of ultra-luminous infrared galaxies (ULIRGs), the nuclei are embedded in large quantities of gas and dust produced by merging processes and massive star formation. In these conditions, the visual extinctions can be as high as Av > 1000 mag when the H2 column densities exceed > 1024 cm−2 in very compact regions of only a few pc (e.g. Costagliola et al. 2015; Aalto et al. 2017). In these extremely obscured environs, even the millimetre (mm)/sub-mm emission can be significantly attenuated by self- or continuum absorption, therefore probing only the gas emitted at distances ≥100 pc from the central engines (e.g. Aalto et al. 2015b; Martín et al. 2016). Complementing mm/sub-mm observations with radiatively pumped molecular lines emitted in the infrared (IR), where most of the ULIRGs luminosity is emitted, allows us to probe the regions of the dusty cores more deeply. In particular, H2O lines are excited via absorption and re-emission of IR photons produced very close to the central engines, and therefore provide essential information about the physical conditions of the hidden power sources of ULIRGs.

Mrk 273 (F13428+5608) is a ULIRG (LIR = 1.3 × 1012L, Gao et al. 1999) located at a distance of 157 Mpc (, z = 0.037780, 1″ = 761 pc1). Its complex morphology reflects a recent merger event between two or more galaxies. Near-infrared (NIR), radio emission, and HI images show two nuclei separated by ∼1″ in the northeast-southwest direction (hereafter the northern and southern nucleus) plus a weaker third source that might be another nucleus or a starburst region triggered by the merger (Majewski et al. 1993; Cole et al. 1999; Condon et al. 1991). The nature of the progenitors has been the subject of several studies that claim either AGN or starburst activities. Mrk 273 is classified as a Seyfert 2 galaxy in the optical and NIR (Vivian et al. 2013; Rodríguez Zaurín et al. 2014; Iwasawa 2018), having an AGN bolometric luminosity of log(LAGN) = 44.73 erg s−1, and a ratio between the bolometric luminosity of the AGN and the total bolometric luminosity of the galaxy of LAGN/LBol = 0.08 (Nardini et al. 2009). Nevertheless, far infrared (FIR) and mm data point to a compact ultra-luminous starburst region in the northern nucleus (Condon et al. 1991; Majewski et al. 1993; Downes & Solomon 1998). The total star formation rate (SFR) of the galaxy is 139 M yr−1 (Cicone et al. 2014).

Molecular observations reveal a complex structure in the centre of Mrk 273. Carbon monoxide maps show extended gas streamers in the north-south direction, a nuclear disc (of 2″ size) oriented east-west, and a very compact core (0.35″ × 0.2″) containing a powerful starburst (Downes & Solomon 1998). All these components belong to the northern nucleus (which is the strongest radio source), while there is no evidence of significant amounts of molecular gas in the southern objects.

A cool molecular outflow has been detected in Mrk 273 using H2, CO, and OH observations (Vivian et al. 2013; Veilleux et al. 2013; Cicone et al. 2014; González-Alfonso et al. 2017). From CO (1 − 0) observations, the outflow appears clearly in both line wings at high velocities |v| > 400 km s−1, and extends from the northern nucleus about 600 pc to the north. The CO channel-velocity maps also show a low velocity component (v > 150 km s−1) of the red-shifted outflow expanding to the north. The total mass of the outflow estimated from CO is ∼2 × 108M.

The multi-phase composition of the outflow has been revealed by infrared and optical observations. Hydroxide (OH) detections (Veilleux et al. 2013; González-Alfonso et al. 2017) probe a more compact and warmer phase of the outflow that expands at velocities of 100–700 km s−1. This component travels shorter distances (160 pc) before it presumably cools down, and its mass and mass-loss rate are 5 × 107M and 120 M yr−1. Colina et al. (1999) detected an ionised component of the outflow by using the [OIII] λ = 5007 Å optical line. The hot, ionised gas goes much further, as far as ±5″ (∼3.8 kpc) along the north-south direction, and reaches velocities as high as ±1200 km s−1. Other IR and optical lines, namely H2, HeI, Brγ, and [CII], also show the low-velocity component of the outflow (±200 km s−1), as well as a moderate velocity component (600 ± 300) km s−1 heading towards the north (Vivian et al. 2013; Janssen et al. 2016).

Despite being one of the most luminous ULIRGs in the local universe, molecular detections towards Mrk 273 were still scarce and limited to CO, OH, and H2. In this paper we present observations of molecules detected for the first time in this galaxy in the mm/sub-mm and FIR ranges. In Sect. 2 we present our observations with the NOEMA and Herschel telescopes and the data reduction. The continuum and spectroscopic data analyses are presented in Sect. 3, where we also describe our modelling of H2O. The asymmetric, double-peaked line profiles of the inner disc are discussed in Sect. 4.1. The detection of the Mrk 273 outflow and its properties are presented in Sect. 4.2. A brief discussion of the expansion of the galaxy core can be found in Sect. 4.3. The non-detection of vibrationally excited HCN and HC3N emission is addressed in Sect. 4.4. We also discuss the HCN/HNC, HCN/HCO+, and HCO+/HOC+ brightness temperature ratios (Sect. 4.5). The different origin of HOC+, the only species not peaking at the very centre, is discussed in Sect. 4.6. Finally, our conclusions are summarised in Sect. 5.

2. Observations and data reduction

2.1. NOEMA

The HCN, HCO+, HNC (1 − 0) and HC3N (10 − 9) lines were observed simultaneously with the NOEMA interferometer on April 7, 2015 (with seven antennae), and June 12 and 15, 2017 (with eight antennae during a total time of 8.9 h (precipitable water vapour (pwv) ∼ 5–9 mm), while the HCN, HCO+ and HOC+(3 − 2) transitions were observed, also simultaneously, for 5.6 h (pwv ∼ 1–2 mm). The receivers were tuned to centre their lower side bands at 85.8 GHz (λ ∼ 3 mm) and 257.5 GHz (λ ∼ 1 mm), respectively. The receivers were connected to the WideX correlator and provided a 3.6 GHz bandwidth in two orthogonal polarisations (which were averaged). The data were calibrated and imaged with CLIC and MAPPING within the GILDAS package2. Source 3C84 was used as a bandpass calibrator, and the phase and flux calibrations were done by observing 1418+546 and MWC349, respectively.

The observations were centred on RA (J2000) = 13h44m42ṣ1, Dec (J2000) = 55°53′13ʺ̣5. The original spectral resolutions of 6.8 km s−1 (3 mm) and 2.3 km s−1 (1 mm) were smoothed to 68–70 km s−1. The final rms of the cubes, averaged across the spectral channels that do not contain line emission, are 0.3 mJy beam−1 (3 mm) and 1.3 mJy beam−1 (1 mm). We used a natural weighting mode and 0ʺ̣8 (for the 3 mm data) and 0ʺ̣8 (for the 1.3 mm data) pixel sizes to create the data cubes, and the Hogbom deconvolution method to clean them. The sizes of the primary beams were 58ʺ̣6 and 19ʺ̣6, and the angular resolutions achieved were (4ʺ̣9 × 4ʺ̣5) and (0ʺ̣61 × 0ʺ̣55) with position angles (PA) of −80° and +34° for the 3 mm and 1 mm transitions, respectively. At the adopted distance of Mrk 273, these resolutions correspond to (3.7 × 3.4) kpc and (460 × 420) pc spatial scales.

Given the NOEMA configurations used, the maximum recoverable scales for our 1 mm and 3 mm observations are 2ʺ̣5 and 17″ respectively, which are well above (about one order of magnitude) the emission sizes of the molecules (see Sect. 3.3.1 and Table 3). Furthermore, our HCN and HCO+ fluxes are consistent with those obtained by Graciá-Carpio et al. (2008) with the IRAM 30 m single-dish telescope. We can, therefore, safely claim that there was no flux filtered out in our observations.

2.2. Herschel

The H2O data were taken with the Photodetector Array Camera and Spectrometer (PACS; Ott 2010; Poglitsch et al. 2010) and the Spectral and Photometric Imaging Receiver (SPIRE; Griffin et al. 2010) on December 16, 2012 and November 21, 2010 respectively. The absorption water lines observed with Herschel/PACS are new detections in Mrk 273, while the emission transitions detected with Herschel/SPIRE were already reported by Lu et al. (2017). The PACS observations3 (PI: González-Alfonso) were performed in high spectral sampling range spectroscopy mode in first and second orders of the grating, resulting in a velocity resolution of ∼170–265 km s−1. The spectra were reduced with the standard product generation pipeline version 14.2.0. The nuclear FIR emission from Mrk 273 was unresolved with the PACS 5spaxel × 5spaxel IFU with 9ʺ̣4 × 9ʺ̣4 spaxels, so the spectra were extracted using the point source correction task available in the Herschel interactive processing environment (HIPE; Ott 2010) version 14.0.1. The spectra were then scaled to the integrated flux level of the central 3 × 3 PACS spaxels to compensate for pointing offsets and jitter which act to move flux out of the central spaxel. The H2O spectra were sampled in velocities of 20–40 km s−1 per channel width. Polynomial baselines of orders lower than three were then removed, and the final root mean squares (rms) are 0.2–0.3 Jy km s−1. The lines were fitted with Gaussian line profiles using the GILDAS software CLASS (Fig. 3).

The SPIRE observation4 (PI: P. P. van der Werf) was conducted using a single pointing centred on Mrk 273 in high-spectral-resolution sparse image sampling mode with a resolution of 1.2 GHz (∼250–360 km s−1) in the two observing bands (447–989 GHz and 958–1545 GHz). In total, 99 repetitions (198 FTS scans) were performed for a total on source time of 13187 s. The data reduction was done with the standard single pointing pipeline available in HIPE version 14.0.1 and a bootstrap method was used to extract the line fluxes. The individual scans were averaged together and a polynomial baseline was extracted from each detector before all lines were fitted simultaneously using Gaussian profiles convolved with a sinc function (Fig. 4). After 1000 repetitions of this procedure, Gaussians were fitted to the resulting flux distribution of each line to get the mean line flux together with its standard deviation.

3. Results

3.1. Continuum

We first created the 3 mm (∼89 GHz) continuum visibilities including only the channels free of line emission after smoothing to the final spectral resolution. These visibilities were then subtracted from the total emission in the uv plane. Using the task uv_fit within the MAPPING package, we measured the size and flux of the 3 mm continuum. Our values were calculated by fitting an elliptical Gaussian in the Fourier plane. Circular and elliptical fits gave consistent values within the errors. Given that the images of the continuum and line emissions are quite round, but not entirely, we opted for using elliptical fits in order to take into account small asymmetries in the emission. We measured a spatially integrated flux density of the 3 mm continuum of 8.24 ± 0.07 mJy and a deconvolved size of (1.92 ± 0.06)″ × (1.83 ± 0.07)″ with a PA = (−90 ± 30)°. This is similar to the 111 GHz continuum flux density of 11 ± 2 mJy obtained by Downes & Solomon (1998).

Due to the very broad line widths at zero intensity (500–1100 km s−1) in the 1 mm band (∼265 GHz) and the narrow bandwidth of the correlator, almost all channels contain line emission. Therefore, the continuum visibilities were created using only nine line-emission-free channels (15% of the total number of channels) at both edges of the spectrum. Including more channels could potentially lead to an overestimation of the continuum flux. A fit of an elliptical Gaussian in the uv plane gives an integrated flux of 28.6 ± 0.9 mJy, with a deconvolved size (FWHM – full with at half maximum – of the Gaussian) of (0.36 ± 0.03)″ × (0.27 ± 0.03)″ ∼ (270 × 200 pc) with a position angle of (24 ± 12)°. The integrated intensities of the continuum at 1 mm and 3 mm are plotted in Fig. 1.

thumbnail Fig. 1.

Continuum maps at 3 mm and 1 mm. Contour levels start at a significance of 5σ with respect to the rms measured in both images (rms = 0.07 mJy channel−1 and 0.6 mJy channel−1 for the 3 mm and 1 mm maps respectively). The contour steps are 1 and 3 mJy km s−1 beam−1. The crosses at the centres mark the continuum peaks, which we take as the location of the northern nuclear source (see Sect. 1). The synthesised beams are shown in the bottom-left corner. The colour flux scales are in Jy km s−1 beam−1.

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3.2. Line profiles

3.2.1. Nuclear emission

Figure 2 shows the spectra of all observed lines with NOEMA extracted from the whole region (top panels), and from the central pixel of the observations (bottom panels). The total emission was integrated using masks in the moment zero maps of the HCO+(1 − 0) and (3 − 2) lines, which show the most extended emission at 1 mm and 3 mm (Table 3). These masks were also used to integrate the emission of the more compact species observed at similar frequencies (i.e. one mask over HCO+(1 − 0) for all 3 mm species, and other mask over HCO+(3 − 2) for all 1 mm species). While HOC+(1 − 0) is not detected, the (3 − 2) transition is seen arising from a very compact region near the nucleus. When integrating the overall flux of the map, the emission of the HOC+(3 − 2) line drops below the noise and, therefore, for comparison with the other molecules, we measure its flux in the pixel where it peaks (see Sect. 3.3 for details). There, HOC+(3 − 2) is detected with a signal-to-noise ratio (S/N) of seven. Additionally, its central velocity is blue-shifted compared to the other lines observed with NOEMA (Table 1).

thumbnail Fig. 2.

Emission lines detected with NOEMA (black histograms) and Gaussian fits (total: red lines; if two components are present, individual components are displayed in blue colour). The velocity resolution is 68 km s−1 in all cases. The labels in the top right corners indicate if the spectra were extracted from the integrated emission (“Integrated”, top panels), from the central pixel (bottom panels), or in the case of HOC+, from the pixel at (0ʺ̣2, −0ʺ̣05).

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Table 1.

Gaussian fit parameters to the emission lines detected with NOEMA.

The 3 mm lines (HCN, HCO+, HNC and HOC+(1 − 0), and HC3N(10 − 9)) have roughly Gaussian-like profiles, although their peaks are slightly flat-topped. These lines were fitted with single Gaussian velocity profiles (Fig. 2). On the other hand, the 1 mm lines (HCN, HCO+ and HOC+(3 − 2)) are double-peaked. To account for these profiles, we fitted two Gaussians (their parameters are listed in Table 1). The intensity of the dip between the double peaks is 19 mJy and 26 mJy for the HCN and HCO+(3 − 2) lines, respectively. The dip appears at slightly blue-shifted velocities, specifically at −30 km s−1 for HCO+(3 − 2), and −60 km s−1 for HCN(3 − 2). The nature of these double-peaked profiles is further discussed in Sect. 4.1.

Figures 3 and 4 show the H2O line profiles observed with Herschel PACS and SPIRE and the best Gaussians fits. All transitions were fitted by single Gaussian profiles (for the SPIRE lines, the Gaussians were convolved with a sinc function, Sect. 2.2). H2O (524 − 413) is partially blended with an OH line at 71 μm. To disentangle the emission of the two species, we fitted a double Gaussian profile to the observed spectrum.

It is important to note that all Gaussian fits to the water lines observed with PACS are blue-shifted to velocities between −20 km s−1 and −140 km s−1 (Table 2). Interestingly, these values are, to within the errors, the same as the velocities of the dips in the profiles of HCO+ and HCN(3 − 2). The connection between the two is discussed in Sect. 4.1.

thumbnail Fig. 3.

H2O absorption lines (black histograms) detected with Herschel/PACS and Gaussian fits (red lines). The velocity sampling is 20–40 km s−1.

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thumbnail Fig. 4.

H2O lines (black histograms) detected with Herschel/SPIRE, and Gaussians fits (convolved with sinc functions).

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Table 2.

Gaussian fit parameters to the water lines detected with Herschel.

3.2.2. High-velocity emission

None of the 3 mm and FIR water lines show obvious extended line wings, which would reveal the Mrk 273 molecular outflow previously detected with CO, OH, and H2 by Vivian et al. (2013), Veilleux et al. (2013), Cicone et al. (2014), and González-Alfonso et al. (2017). For the 1 mm observations, the spectrum of the overall integrated emission has no evident signatures of line wings either. However, in the central pixel, the HCN profile exhibits a line wing that extends between −400 and ∼1000 km s−1, while the red side shows no wing (Fig. 5). This emission is detected with an S/N of ∼5.

We used the JPL catalog (Pickett et al. 1998) to look for lines arising from 266.2 to 266.9 GHz, which correspond to the velocity range [−350,−1200] km s−1 where the HCN(3 − 2) wing-like feature is seen. CH2NH (41, 3 − 31, 2) is the most likely line arising at these frequencies, with an energy level of Elow = 19 K. This transition was detected in the LIRGs IC 860 and Zw 49-57, where its flux density is three to four times fainter than HCN(3 − 2) (Aalto et al. 2015b). Assuming a similar ratio and excitation conditions in Mrk 273, CH2NH would then have a peak flux between 4.5 and 6 mJy, that is, brighter than the emission we see. The HCN shoulder does not have a (single or double) Gaussian profile similar to the detected lines, but has the shape of a line wing. For these reasons, it seems unlikely that the emission comes from the CH2NH line, although a potential contamination cannot be ruled out.

thumbnail Fig. 5.

Spectrum of the central pixel for the (J − J′) = (3 − 2) lines of HCN, HCO+ and HOC+. The baseline of order 0 and the 3 × rms flux (calculated at the final velocity resolution of 68 km s−1) are marked with horizontal dashed lines. The outflow emission at the blue-shifted velocities of HCN is highlighted in yellow.

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The two nuclei of Arp 220 have HCN (3 − 2) and (4 − 3) blue-shifted wings very similar to what we observe here in Mrk 273 (Martín et al. 2016). The fact that these bumps appear in both HCN transitions, while there is no corresponding CH2NH line close to the HCN(4 − 3) frequency, strengthens our claim that the line wing in Mrk 273 comes from HCN(3 − 2), and that it is tracing gas outflowing at high velocities. In summary, given the feature intensity, spectral shape, and integrated emission (see following section), we conclude that the HCN(3 − 2) blue-shifted shoulder comes from the outflowing gas moving at approximate velocities between −400 and −1000 km s−1. In Sects. 3.3.2 and 4.2 we discuss in more detail the properties of this molecular outflow.

3.3. Moment maps of HCN, HCO+, HNC and HOC+

3.3.1. Position and size of the nuclear emission

Figure 6 and Table 1 show the integrated intensities of the HCN, HCO+, HNC and HOC+ lines. The deconvolved emission sizes and position angles of all lines were measured in the uv plane by fitting elliptical Gaussians with the task uv_fit within MAPPING, and are listed in Table 3. The continuum and the HCN, HCO+, and HNC lines have their peak intensities at the central pixel, which we take as the location of the northern nuclear source. On the contrary, the HOC+ maximum is found at an offset (0ʺ̣2, −0ʺ̣05), corresponding to (152, −38) pc to the south-east. The position accuracy of our observations can be calculated from(1)

thumbnail Fig. 6.

Integrated intensities (moment zero maps). Contour levels for HCN, HCO+ and HNC (1 − 0) go from 0 to 6 Jy kms−1 beam−1 with a step of 0.5 Jy kms−1 beam−1. For HCN (3 − 2) and HCO +(3 − 2) the levels range from 2 to 30 Jy kms−1 beam−1 with a step of 2 Jy kms−1 beam−1. For HOC+(3 − 2), the contours are from 0.4 to 0.6 Jy kms−1 beam−1 with a step of 0.04 Jy kms−1 beam−1. We highlight the different scales between the (1 − 0) and the (3 − 2) lines. The crosses in the centre of each panel indicate the location of the nuclear source. The synthesised beam is shown in the bottom-left corner of each panel. North is up, and east is to the left.

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Table 3.

Global properties of the lines detected with NOEMA.

where Δα, Δδ are the errors in RA and Dec, and θb is the synthesised beam size (Condon 1997; Ivison et al. 2007). For HOC+(3 − 2), we have a Δα, Δδ≃ 0ʺ̣1 (∼80 pc), confirming our claim that its shift is real, at least in RA (because the HOC+ offset in declination is smaller than our position accuracy). In Sect. 4.6, we further discuss the origin of HOC+.

Downes & Solomon (1998) estimated a CO (1 − 0) emission deconvolved size of (0ʺ̣9 × 0ʺ̣6) with a beam of (1ʺ̣4 × 1ʺ̣3). We obtain larger sizes for the HCN and HCO+(1 − 0) lines (∼(2″ × 2″) equivalent to ∼1.5 kpc, Table 3), most probably because we collect more emission within our larger beam of (4ʺ̣9 × 4ʺ̣5). In the following, we refer to this region as the outer disc. The HCN and HCO+(3 − 2) lines, however, are confined to a much smaller region of (0ʺ̣4 × 0ʺ̣3) (∼(300 × 230) pc, Table 3), implying that the denser, star-forming gas is more concentrated in the nucleus. We refer to this as the inner disc. Despite our high angular resolution at 267 GHz, the nucleus is still unresolved and does not show any structures in the moment zero maps, implying a very compact core that does not expand more than <0ʺ̣3 (230 pc). Indeed, from our H2O modelling, we estimate the size of the core to be ∼50 pc (see Sect. 3.4 for details).

HC3N (10 − 9) and HOC+(3 − 2) are faint (although we detect the latter with an S/N of ∼7 in the pixel where it peaks) and unresolved at our resolution. Therefore, we cannot measure their emission sizes.

3.3.2. Integrated intensity of the outflow

The moment-zero map of the outflow traced by HCN(3 − 2) (highlighted in yellow in Fig. 5) is shown in Fig. 7. After centering the spectrum at the HCN(3 − 2) rest frequency, we integrated the emission of all pixels in the velocity range [−1100, −400] km s−1. Its morphology is composed of two main peaks of emission; a stronger one with an elliptical shape around the centre, and a fainter and rounder feature at ∼0.8″ (∼600 pc) to the north. In the central pixel, the peak flux has a S/N of ∼5 (measured at the final velocity resolution), while the integrated flux density measured in the moment-zero map is 2.4 Jy km s−1.

thumbnail Fig. 7.

Integrated intensity of the outflow seen in HCN (3 − 2) between −1100 km s−1 and −400 km s−1 (yellow region marked in Fig. 5). The cross in the centre marks the location of the nuclear source. The contours start with a 5σ flux and go from 0.8 mJy kms−1 beam−1 to 2.5 Jy kms−1 beam−1 with steps of 0.2 mJy kms−1 beam−1. The magenta dashed ellipses show the regions fitted to the two main components. The synthesised beam is plotted in the bottom-left corner.

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We measured the size of the outflow by fitting two ellipses to the main peaks seen in the moment-zero map (see Fig. 7). We take the size of the outflow as the projected distance between the central pixel and the centre of the northern ellipse. We measure a size of 0ʺ̣61 ± 0ʺ̣05, equivalent to 464 pc, with a position angle of 10 ± 3°. The direction of the flow heading to the north is consistent with the outflowing gas detected with CO (1 − 0) by Cicone et al. (2014).

3.3.3. Velocity fields and position-velocity maps

Figures 8 and 9 show the velocity fields and position-velocity (p − v) maps of HCN, HCO+ and HNC. The velocity fields reveal the rotation of the Mrk 273 nuclear disc (see Downes & Solomon 1998 for a detailed study of the disc properties). There are significant differences in the morphologies of the iso-velocity contours of the various observed lines. HCN and HCO+(1 − 0) trace rotating gas in the outer disc showing a south-east to north-west direction. However, the less extended gas traced by HNC (1 − 0) shows a north-east to south-west rotation, similar to the gas in the inner disc traced by the HCN (3 − 2) and HCO+(3 − 2) lines. The rotation in the central <0ʺ̣5 follows the velocity fields traced by the warm and compact gas traced by H2, Brγ and [FeII] (Medling 2014; Vivian et al. 2013). This is consistent with the two kinematic systems of the disc already discovered by Downes & Solomon (1998) using the CO (1 − 0) and CO (2 − 1) lines with beam sizes (1ʺ̣4, 1ʺ̣3) and (0ʺ̣6, 0ʺ̣6). The agreement between the Downes & Solomon (1998) velocity fields and ours, despite the difference in angular resolution, suggests that this effect is not due to the larger beam size of our 3 mm data.

thumbnail Fig. 8.

Velocity fields (moment 1 maps). The coloured velocity scale (right axis) is in km s−1. The step in contours is 20 kms−1 for all lines. We note the blue-shifted velocities of HOC+(3 − 2). The crosses in the centre of each panel indicate the position of the nuclear source. The beam is shown in the lower left corner of each panel. North is up, and east is to the left. The dashed lines in the HCN plots indicate the cut for the p − v diagrams shown in Fig. 9.

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thumbnail Fig. 9.

Position-velocity (p-v) maps of HCN and HCO+. The cuts along the axes of rotation are shown by the dashed lines in Fig. 8. Contour steps of the (1 − 0) maps go from 0.9 mJy beam−1 (3σ) to 7.9 mJy beam−1 with steps of 1 mJy beam−1. For the (3 − 2) maps, contours go from 3.9 mJy beam−1 (3σ) to 59 mJy beam−1 with steps of 5 mJy beam−1.

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Figure 9 shows the p − v diagrams of the HCN and HCO+ lines along cuts through the centre and perpendicular to the axes of rotation as plotted with dashed lines in Fig. 8. The outer disc traced by the (1 − 0) lines shows maximum velocities ±300 km s−1, though most of the gas exhibits velocities within the smaller range of ±150 km s−1. The p − v diagrams of HCN and HCO+(3 − 2) show that the gas in the inner disc rotates faster, reaching maximum velocities of ±400 km s−1, although the average speed is ±200 km s−1 at the edges of the disc, which are separated by 0ʺ̣2 (∼150 pc).

The dynamical mass of the inner, starbursting disc, measured from the kinematics of the (3 − 2) lines, is (4 − 5) × 109M, as calculated from . Here, Vrot is the average rotational velocity corrected for the 45° inclination of the disc (Downes & Solomon 1998), R is the size of the emission in pc (Table 3), and G is the gravitational constant.

3.3.4. Velocity dispersions

The velocity dispersions of the NOEMA data were calculated as(2)

The moment-2 maps of HCN, HCO+ and HNC are shown in Fig. 10. All lines reach similar maximum dispersions of 500 km s−1. We explored the Toomre (1964) stability criterion, Q ≥ 1, for the inner, starbursting, gaseous disc to check its stability against gravitational perturbations, as(3)

thumbnail Fig. 10.

Velocity dispersions. Contours go from 0 to 500 kms−1 with steps of 50 kms−1 for all lines. The crosses in the centre of each panel indicate the position of the nuclear source. The synthesised beam is shown in the bottom-left corner of each panel. North is up and East is to the left.

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where κ is the epicyclic frequency, G is the gravitational constant, and Σ is the surface density of the gas. The surface density in the inner 380 pc of Mrk 273 is 37500 M/pc2 (Yun & Scoville 1995). We note that this value was calculated in a region slightly larger than the radius of the inner disc (of ∼300 pc), but it is still a good approximation if compared to the size of the outer disc (1.5 kpc). For the epicyclic frequency, we assumed a Keplerian disc. Therefore, κ = ω, ω being the angular velocity. For consistency, we calculated ω in the same radius of 380 pc, and used the average rotational velocities from our HCN and HCO+(3 − 2) velocity maps (200 km s−1, Figs. 8 and 9). We obtain Q = 0.5, which indicates that the inner disc is unstable and clumpy/turbulent enough to form further self-gravitating condensations of gas.

3.3.5. HCN and HCO+ channel-velocity maps

Figures 1114 show the channel maps of HCN and HCO+(1 − 0) and (3 − 2) between −500 and +500 km s−1 in steps of 50 km s−1, with a beginning intensity contour level of 5σ. The (1 − 0) lines show emission away from the nucleus in all directions. In particular, the HCN(1 − 0) channel map reveals gas in the northern direction as far as 10″ from the centre at negative velocities, as well as elongations towards the south at ±300 km s−1.

thumbnail Fig. 11.

Channel-velocity maps of HCN (1 − 0) in the velocity range [−500, 500] km s−1 with steps of 50 km s−1. Contours go from 1.5 mJy beam−1 (5σ) to 9.5 mJy beam−1 with a spacing of 2 mJy beam−1. The synthesised beam is plotted in the bottom-left corner. North is up and east is to the left.

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thumbnail Fig. 12.

Channel-velocity maps of HCO+(1 − 0) in the velocity range [−500, 500] km s−1 with steps of 50 km s−1. Contours go from 1.5 mJy beam−1 (5σ) to 9.5 mJy beam−1 with a spacing of 2 mJy beam−1. The synthesised beam is plotted in the bottom-left corner. North is up and east is to the left.

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thumbnail Fig. 13.

Channel-velocity maps of HCN (3 − 2) in the velocity range [−500, 500] km s−1 with steps of 50 km s−1. Contours go from 6.5 mJy beam−1 (5σ) to 56 mJy beam−1 with a spacing of 10 mJy beam−1. We highlight the significantly smaller spatial scale relative to the channel map of the (1 − 0) line shown in Fig. 11. The synthesised beam is plotted in the bottom-left corner. North is up and east is to the left.

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thumbnail Fig. 14.

Channel-velocity maps of HCO+(3 − 2) in the velocity range [−500, 500] km s−1 with steps of 50 km s−1. Contours go from 6.5 mJy beam−1 (5σ) to 56 mJy beam−1 with a spacing of 10 mJy beam−1. We highlight the significantly smaller spatial scale relative to the channel maps of the (1 − 0) line shown in Fig. 12. The synthesised beam is plotted in the bottom-left corner. North is up and east is to the left.

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Some extensions to the east and south-east are also seen in some channels. The HCO+(1 − 0) emission is similarly extended, but the structure towards the south is perhaps the most distinct (e.g. channels ± 150 km s−1). Nevertheless, extensions to the north (±200 km s−1 and ±300 km s−1), east (±250 km s−1), and west (+150 km s−1 and +250 km s−1) are also seen.

The HCN and HCO+(3 − 2) channel maps trace denser gas in the inner parts of the galactic disc (≤300 pc), although some emission can be seen in the central ±1″ ( ± 800 pc), and their elongations are even more obvious than in the (1 − 0) channel maps. HCN shows clear signs of emission towards the north in most channels (best seen between −250 km s−1 and +250 km s−1). On the other hand, the gas traced by HCO+ is more extended towards the south (e.g. −200 km s−1 and +50 km s−1).

The most significant extensions to the north (from the HCN maps) and to the south (from the HCO+ maps) are signatures of outflowing gas, and are further discussed in Sect. 4.2.

3.4. Modelling of H2O

We have used the library of H2O models generated by González-Alfonso et al. (2014) to fit the H2O emission and absorption observed in Mrk 273. The models assume spherical symmetry, are non-local, and include excitation by both the FIR field emitted by warm dust (which is mixed with the H2O molecules) and collisions with H2. The collisional rates were taken from Dubernet et al. (2009) and Daniel et al. (2011), and a gas-to-dust ratio of 100 was adopted (Wilson 2008). The models assume uniform physical properties (Tdust, Tgas, gas and dust densities, H2O abundance). The source is divided into a set of spherical shells where the statistical equilibrium level populations are calculated. We assume a H2O ortho-to-para ratio of three. Line broadening is simulated by including a micro-turbulent velocity (Vturb), with no systemic motions.

The modelled line fluxes and continuum flux densities scale as (R/DL)2 where R is the source radius and DL is the luminosity distance, so that they are easily scalable to the properties of any source. Following Falstad et al. (2017) and González-Alfonso et al. (2017), we have fit the line fluxes with a combination of NC model components by minimising χ2, with the radius R of each component being the only free parameter that is varied. We required NC = 2 components to properly fit the PACS and SPIRE H2O fluxes simultaneously. Since the models also make specific predictions for the spectral energy distribution (SED) of each component, and since H2O probes the galaxy FIR emission responsible for its excitation and more specifically the transition from the mid- to the far-IR (González-Alfonso et al. 2010; González-Alfonso et al. 2014), we also included in the fit the continuum flux densities at 30 and 60 μm.

Figure 15 compares the observed H2O fluxes and SED with the results of our best model fit, and Table 4 lists the properties of the two model components (shown with blue and green colours in Fig. 15). The two components show very different behaviours relative to the H2O absorption and emission. We require a very compact (effective radius R ∼ 50 pc), very warm (Tdust ∼ 95 K), and very optically thick (τ100 μm ≳ 4) component (referred to as the “core”) to account for the observed PACS absorption in several lines (primarily the 313 − 202 line at 138 μm, the 422 − 313 line at 58 μm, the 423 − 312 line at 79 μm, and the 524 − 413 line at 71 μm) and also the SPIRE emission in the 523 − 514 line at 212 μm (although this line is only marginally detected at the 3σ level). However, the core component predicts negligible emission in most SPIRE lines (and even absorption in the 220 − 211 line), which indicates the presence of a more extended component (R ∼ 280 pc), moderately warm (Tdust ∼ 55 K), and with lower column (τ100 μm ∼ 0.5). This extended component, mostly responsible for the H2O emission observed with SPIRE, is naturally identified with the inner disc traced by the J = 3 − 2 lines of HCO+ and HCN (∼300 pc).

thumbnail Fig. 15.

Comparison between the observed (black symbols) and modelled (coloured symbols and lines) H2O fluxes in Mrk 273 within panel a: PACS and panel b: SPIRE. The model includes two components: the core (in blue), which accounts for most absorption lines observed within PACS, and the inner disc (in green), which dominates the emission of the sub-mm lines with Eupper < 400 K observed with SPIRE (see Table 4). Red colours and symbols indicate total modelled fluxes. The numbers at the bottom of panels a and b indicate rounded-up transition wavelengths in μm. Panel c: the SED of Mrk 273, including the Spitzer/IRS spectrum, Herschel/PACS and SPIRE continuum data from observations of both H2O and OH lines, sub-mm data at 800 and 880 μm (Rigopoulou et al. 1996; Wilson 2008), and our measured flux density at 1 mm (starred-red symbol), is compared with the prediction of our composite model.

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Table 4.

Results for the two-component modelling of the H2O lines in Mrk 273.

The two water components together provide a good fit to the FIR emission, though the combined model underestimates, to some extent, the flux densities between 100 and 200 μm. This probably indicates a range in Tdust for the disc component, rather than a single value. It is also worth noting that the source luminosity is dominated by the disc (i.e. the starburst), with the core component accounting for LIR ∼ 4 × 1011L. The latter corresponds to ∼30% of the total IR luminosity, which is close to the estimated AGN contribution based on the 15–30 μm diagnostic (e.g. Veilleux et al. 2009). The blueshift of the absorption lines, which are dominated by the core component, is similar to that seen in the excited OH lines (González-Alfonso et al. 2017), suggesting that the core is expanding (more details in Sect. 4.3). Our model for the core also predicts significant emission at 265 GHz, ∼10 mJy, though this value is relatively uncertain because of its dependence on the actual continuum optical depth and the mass absorption coefficient of dust (κλ) at millimetre wavelengths5.

4. Discussion

4.1. A rotating disc with continuum absorption

The integrated intensities, velocity fields, and p − v diagrams of HCN and HCO+ show the typical pattern of a rotating disc (Figs. 6, 8 and 9). Such a rotating body should in principle be reflected in the spectra as double-peaked lines. However, due to our relatively large beam size used to observe the (1 − 0) lines (4ʺ̣9 × 4ʺ̣5), most of the gas is concentrated in the central pixel and the velocity gradients within the beam are not well reproduced. Therefore, our HCN and HCO+(1 − 0) lines show a single Gaussian-like profile. On the other hand, the (3 − 2) transitions were observed with a resolution significantly higher (0ʺ̣61 × 0ʺ̣55), so the velocity gradients are well traced and the spectral lines reveal the expected pattern of the inner rotating disc (see Sect. 3.3.1).

In addition, continuum absorption is also apparently contributing to the shape of the HCN and HCO+(3 − 2) profiles. Firstly, we note that the ∼20–25 mJy drop flux of the lines is very similar to the continuum flux density at these frequencies, of 29 mJy. Secondly, the velocity of the minimum flux, around −50 km s−1, roughly coincides with the peak absorption velocities of the H2O lines observed with PACS, and of the OH 84 μm and OH 65 μm lines presented in González-Alfonso et al. (2017). The foreground warm gas traced by H2O and OH is absorbing the continuum (e.g. González-Alfonso et al. 2017), and the coincidence in velocities and continuum values of the dip in HCN and HCO+ indicates that the dense gas might also be absorbing the background dust emission.

In this context, the two peaks of the HCN and HCO+(3 − 2) lines are probing the edges of the inner rotating disc, while the absorbed flux indicates the positions of the maximum column densities of the gas. The channel maps shown in Figs. 13 and 14 also show a minimum emission around −50 km s−1 and −100 km s−1 (better seen in the HCO+(3 − 2) map), velocities at which there is a maximum absorption of the continuum.

Self-absorption might also be an extra factor affecting the line shapes if there is cooler foreground gas with high-enough column densities. This, however, is difficult to disentangle from the continuum absorption in our data. In addition, it is possible that the opacity of the gas in the centre of Mrk 273 is high enough to result in flat-topped profiles such as some of those observed in the HCN and HCO+(1 − 0) transitions, even though these species are less abundant than CO. We note that any kind of absorption implies that our estimations of line fluxes, luminosities, and molecular gas mass are lower limits to the actual values.

4.2. The Mrk 273 molecular outflow

As commented in Sect. 3.2.1, all the water lines observed with Herschel/PACS are consistently blue-shifted with respect to the systemic velocity of Mrk 273. The absorption in the various lines peaks in the range [ − 20, −140] km s−1 and extends as far as −600 km s−1. Velocity shifts are also observed in OH 65 μm and OH 84 μm (González-Alfonso et al. 2017), and trace the low-velocity gas of the approaching component of the outflow (González-Alfonso et al. 2017).

The channel maps of HCN and HCO+ (Figs. 1114) show clear emission elongated to the north and south. This seems to correspond to the red-shifted low-velocity component (|v − vsys| < 400 km s−1) of the wind heading to the north observed with CO (Cicone et al. 2014). In addition, we also see emission in the north-south direction at higher velocities (up to ∼400 km s−1). In our data, we distinguish two velocity components of the outflow: one with relatively low velocities (|v − vsys| < 400 km s−1) that is seen only in the channel maps (because its emission is blended with that of the disc in the spectra); and a high-velocity component (|v − vsys| > 400 km s−1) that is seen in the spectrum of the central channel as a blue-shifted shoulder of the HCN (3 − 2) line (Fig. 5). The HCN (3 − 2) spectral bump spans approximately from −400 km s−1 to −1000 km s−1 (when centring the line at the rest frequency of HCN(3 − 2)), which is consistent with the outflow velocities measured with CO (Cicone et al. 2014). The channel maps of HCN and HCO+ (Figs. 1114) show clear emission elongated to the north and south. This seems to correspond to the red-shifted low-velocity component (|v − vsys| < 400 km s−1) of the wind heading to the north observed with CO (Cicone et al. 2014). In addition, we also see emission in the north-south direction at higher velocities (up to ∼400 km s−1). In our data, we distinguish two velocity components of the outflow: one with relatively low velocities (|v − vsys| < 400 km s−1) that is seen only in the channel maps (because its emission is blended with that of the disc in the spectra); and a high-velocity component (|v − vsys| > 400 km s−1) that is seen in the spectrum of the central channel as a blue-shifted shoulder of the HCN (3 − 2) line (Fig. 5). The HCN (3 − 2) spectral bump spans approximately from −400 km s−1 to −1000 km s−1 (when centring the line at the rest frequency of HCN(3 − 2)), which is consistent with the outflow velocities measured with CO (Cicone et al. 2014). As explained in Sect. 3.3.2, the size of the fast wind traced by HCN(3 − 2) is 0.″61 ± 0.″05 (∼460 pc) and the elongated shape towards the north can be seen clearly (Fig. 7). The CO (1 − 0) outflowing gas extends up to 550–600 pc in Mrk 273 (Cicone et al. 2014). This difference suggests that the expelled moderate-density gas (nH2 ≤ 103 cm−3) travels further than the dense gas (nH2 ≥ 104 cm−3). In the right panel of Fig. 16, we show a sketch of the dense and warm outflow properties derived from our observations of Mrk 273, as well as the comparison with the diffuse phase of the outflow detected in CO by Cicone et al. (2014).

thumbnail Fig. 16.

Left: sketch of the blue-shifted molecular outflow in the northern nucleus of Mrk 273. The dense and warm outflowing gas found in our HCN and H2O data, and its properties, is compared to the diffuse outflow found by Cicone et al. (2014) by observing CO (1 − 0). We note that the mass loss rate is calculated as Ṁ = v (MOF/ROF) (see Sect. 5.5 in González-Alfonso et al. 2017), which is a factor of three lower than the value given by Cicone et al. (2014). The position angles of the outflow (10°) and the inner disc (70°) are represented. Right: a face-on sketch of the three disc components identified in our data (outer disc, inner disc, and core), which are plotted in scaled sizes. The decoupled kinematics of the outer and inner discs are represented by the blue and red colours, which depict the orientation of the blue-shifted and red-shifted rotating gas. The rotation direction is indicated by the curved arrows. The intensity of the colours illustrate the velocities; the inner disc rotates at a higher speed than the outer disc. The stars in the inner disc aim to show the region where most of the starburst is located. The radial arrows represent the low-velocity expansion of the core.

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The HCN fast outflow luminosity is . To estimate the dense gas mass contained in the outflow, we used the relation Mdense = 10 × L′[HCN(3 − 2)] (Gao & Solomon 2004). We note that this formula might overestimate the actual value, because it assumes that all gas is virialised, which might not be the case in the outflow. Additionally, the HCN-to-H2 conversion factor, which is not well constrained in ULIRGs, could also lead to an overestimation of the gas mass. A detailed analysis of the factors affecting the gas mass estimation can be found in Gao & Solomon (2004). We obtain , which is consistent with the value obtained from OH observations (∼1.6 × 108M, González-Alfonso et al. 2017), and is also similar to the dense gas mass in the outflow of Mrk 231 (∼4 × 108M, Aalto et al. 2015a). We note that our result refers only to the high-velocity gas (|v − vsys| > 400 km s−1), as it is not possible to separate the slow component of the outflow from the disc emission in the HCN data.

In cases where the outflow and the disc kinematics cannot be distinguished well in the velocity maps (as in our Fig. 8), a difference in their position angles can help to probe the kinematic decoupling between the two, and can also yield information about the nuclear powering source. Gas in starburst-powered outflows is always expelled perpendicular to the plane of the galaxy (thus showing a change of 90° with respect to the PA of the starbursting disc), while AGN-powered outflows might have virtually any position angle, because the dusty torus and the accretion disc can be tilted with respect to the disc (García-Burillo et al. 2015). In Mrk 273, we measure a difference between the PA of the inner disc (of 71° ± 5°, from HCO+(3 − 2), Table 3) and the PA of the outflow (10° ± 3°), of 60° ± 8°. This seems to indicate that the outflow is powered by the central AGN. Indeed, Cicone et al. (2014) found that the high outflow mass-loss rate of Mrk 273 is consistent with the linear correlation they observed between the bolometric AGN luminosity and the mass outflow rate in local AGN-host galaxies.

A couple of interesting issues arising from our data are 1) the red-shifted component of the outflow traced by HCN (3 − 2) is seen only at low velocities in the channel maps (|v − vsys| < 400 km s−1). Indeed, a fast red-shifted wind, which should be visible in the spectra (with a similar intensity as the blue-shifted one), does not appear; and 2) the apparent absence of high-velocity winds in HCO+ (i.e. HCO+ line wings at |v − vsys| > 400 km s−1). Regarding the first issue, deeper observations are required to improve the S/N of the outflow, to make sure that the red-shifted component is not present. If that were the case, then it would imply that either there is a density or temperature gradient between the receding and the approaching gas (with the red-shifted gas being less dense and/or warm in general), or there is a chemical differentiation between the two. These are not uncommon characteristics in galactic winds, such as in Mrk 231 (Aalto et al. 2012; Aalto et al. 2015a). The chemical differentiation could also explain why there is not enough high-velocity gas traced by HCO+ in the outflow (see Lindberg et al. 2016 for details about this phenomena in the wind of Mrk 231).

4.3. Expansion of the core

In Mrk 273, the velocity peaks of the OH 65 μm and OI 63 μm lines are shifted in comparison to the emission peak of [CII] at 158 μm (González-Alfonso et al. 2017). The [CII] maximum is found at zero velocities (with respect to the systemic velocity of the galaxy), and arises from the bulk of the warm gas in the nucleus. The OH 65 μm and OI 63 μm lines, however, have their maximum (absorption/emission) peaks shifted by ±50 km s−1, and trace the blue component of the outflowing gas moving at low velocities. Apart from being a sign of a superwind, González-Alfonso et al. (2017) argue that this change in the redshift also indicates that the large columns of gas close to the central engine are expanding at low velocities. This effect has been found in a number of other ULIRGs where low-velocity outflows are also present (González-Alfonso et al. 2017).

As discussed in Sect. 4.1, the shift of approximately 50 km s−1 seen in OH 65 μm and OI 63 μm also appears in our H2O data. In particular, the H2O lines observed with PACS, which according to our models trace the compact core with a radius ∼50 pc, are also blue-shifted with respect to [CII] by [ − 20, −140] km s−1 (Table 2). This indicates that, apart from rotating, the gas in the core is expanding at low velocities, pushing outwards the boundaries with the inner disc traced by the HCN and HCO+(3 − 2) lines. Indeed, the expansion is also reflected in our HCO+ and HCN(3 − 2) data. While HCO+ and HCN(3 − 2) double-peaks probe the outer edges of the inner disc (at a radius of ∼300 pc), their peaks of absorption probe the kinematics of the bulk of the disc, which is also shifted by 50–100 km s−1. We interpret this as the gas in the inner disc being pushed by the expanding gas of the core. We illustrate this in the sketch shown in the right panel of Fig. 16.

4.4. Non-detection of vibrational emission

Rotational transitions within vibrationally excited levels of HCN and HC3N have been observed in the central < 100 pc of several (U)LIRGs, probing regions of high temperatures of ≥100 K (Sakamoto et al. 2010; Costagliola et al. 2015; Aalto et al. 2015b; Martín et al. 2011; Martín et al. 2016), and high column densities. Their excitation cannot be explained by collisional effects alone, and mid-IR pumping is necessary in order to populate the upper energy levels and fit their observed luminosities (Aalto et al. 2015b). In that case, the vibrational lines may be more suitable to study the optically thick dust cores of galaxies than the rotational transitions (see Aalto et al. 2015b for a detailed discussion).

In our NOEMA observations of Mrk 273, we do not detect the vibrationally excited line HCN(3 − 2) v2 = 1. The line is split into two components with energy levels of similar intensity, v2 = 1e and v2 = 1f, at frequencies 265.8 and 267.2 GHz. The first one is completely blended with the HCN (3 − 2) rotational line due to the galactic broad line widths, and it is not possible to estimate its peak temperature. However, the v2 = 1f transition would appear as a bump only partially blended with the red side of the HCO+(3 − 2) transition. We calculate a limit to its flux density in the central pixel of 1.1 mJy, which is slightly below the 1σ rms level (of 1.2 mJy, calculated at the final velocity resolution). Assuming that the vibrational line has the same FWHM as the HCN rotational transition (∼600 km s−1), the limit to the integrated intensity was estimated as I ≤ 3 × r m s × F W H M × Dv $ I\,{\le}\, 3\,{\times}\, rms \,{\times}\,\sqrt{FWHM \,{\times}\, \mathit{Dv}} $ = 4 K km s−1 (Dv, the final spectral resolution, is 68 km s−1). Under the conservative assumption that both rotational and vibrational lines have the same source size, the upper limit to the HCN-vib luminosity is L′ = πR2I ≤ 7.3 × 105 K km s−1 pc2 (R is the source size of the HCN (3 − 2) emission specified in Table 3). This is at least a factor of two fainter than the LHCNvib(3 − 2) in Mrk 231 (1.7 × 107 K km s−1pc2, Aalto et al. 2015a). This is interesting, because Lahuis et al. (2007) detect the 14 μm HCN-vib lines in absorption at the same level in both Mrk 273 and Mrk 231.

Aalto et al. (2015b) found a tentative correlation between the outflow velocity and the intensity of HCN vibrationally excited lines in a moderate sample of ten (U)LIRGs. Galaxies with fast outflows (i.e. when the outflow velocity exceeds the escape velocity of the nuclear region) tend to have fainter vibrational lines. With its non detection of HCN(3 − 2) v2 = 1, and its fast molecular outflow, Mrk 273 also follows this relationship. The reason for this correlation, if true, requires further investigation.

Similarly, we do not detect the vibrationally excited line HC3N (10 − 9) v6 = 1, even though the energy of its lower level is below that of the HCN-vib lines (Table 7). We obtain an upper limit to the peak flux density of 0.3 mJy. This line was first tentatively detected in the LIRG NGC 4418 with the IRAM 30 m telescope (Aalto et al. 2007; Costagliola & Aalto 2010) and later confirmed by ALMA observations (Costagliola et al. 2015) with a rotational-to-vibrational intensity ratio of ∼11. If we extrapolate that ratio to Mrk 273, then, based on our results, one would expect a flux density of ∼0.2 mJy. Under the assumption that, in Mrk 273, HC3N(10 − 9) v6 = 1 likely arises from a region smaller than HCN(1 − 0) (i.e. θ <  2.″0 × 1.″6), the limit to its luminosity is 6 × 106 K km s−1 pc2.

Apart from sensitivity issues, two possibilities can potentially explain why we do not detect vibrationally excited emission of HCN and HC3N in Mrk 273. One is that the mid-IR radiation density in the core of Mrk 273 is not as strong as in other galaxies hosting compact obscured nuclei, such as Mrk 231, Arp 220 or NGC 4418. The other is that, although the mid-IR radiation is strong, the obscuration in the centre of Mrk 273 is so high that even the emission of vibrationally excited lines is extinguished.

4.5. Molecular line ratios

Tables 5 and 6 list the brightness temperature ratios of the species detected with NOEMA. The ratios were obtained by comparing the temperatures of the lines calculated over the entire emission (Table 5) and in the central pixel (Table 6). We note that the 3 mm lines (HCN, HCO+, HNC(1 − 0) and HC3N(10 − 9)) and the 1 mm lines (HCN, HCO+, HOC+(3 − 2)) were observed with angular resolutions that differ in the covering area by a factor > 60. In order to avoid confusion and unphysical results, we only discuss ratios in which both lines were observed with the same resolution. In the following we discuss the most relevant ratios.

Table 5.

Brightness temperature ratios in Mrk 273 (row/column), evaluated over the entire emission.

Table 6.

Nuclear brightness temperature ratios in Mrk 273 (row/column), evaluated in the central pixel.

Table 7.

Main parameters of the undetected vibrational transitions of HCN and HC3N.

4.5.1. HCN/HCO+

The integrated intensity ratio between HCN and HCO+ has been widely used to discriminate between starburst-powered galaxies and AGNs (Kohno et al. 2001; Krips et al. 2008; Imanishi et al. 2009; Izumi et al. 2013; Izumi et al. 2016; Martín et al. 2015). An enhancement of the HCN abundance with respect to HCO+ is observed in a fairly large sample of Seyferts, which show a HCN/HCO+ ratio ≥1. Starbursts, on the other hand, are found to have values ≤1. The coexistence of both phenomena in the central few hundred parsecs can however lead to outlier values of the ratio if they are not resolved (Aladro et al. 2015; Privon et al. 2015).

The reason why the HCN/HCO+ ratio is a diagnostic of the galactic central activity is still debated. For more than a decade, the HCN enhancement was claimed to be due to X-ray irradiation in the vicinity of the SMBHs (e.g. Kohno et al. 2001). However, recent studies of NGC 1097 point to high temperatures and/or mechanical heating as a more plausible explanation (Martín et al. 2015; Izumi et al. 2016). Systematic high-angular-resolution observations of a larger sample of galaxies/AGNs is still needed to clarify this issue.

In the case of HCO+, both observations and models point to an enhancement in starburst galaxies, and in particular in photon-dominated regions (PDRs), as a result of strong UV irradiation. Unfortunately, the scenario could be more complex, since chemical models suggest that HCO+ can also be enhanced in regions heavily pervaded by cosmic ray/X-ray fields, such as in AGNs (Meijerink et al. 2006; Aladro et al. 2013). Yet this scenario has not been confirmed with observations, to our knowledge.

Bearing in mind that the reasons for the enhancements of HCN and HCO+ are difficult to assess, the fact is that this ratio diagnoses the power sources of many well-known active galaxies. Nevertheless, there are some sources which do not seem to be consistent with this scenario. For example, some apparent starbursts have HCN-to-HCO+ ratios as high as AGNs, indicating that the HCN enhancement cannot be exclusively attributed to AGNs (Privon et al. 2015).

In Mrk 273, we measure the global brightness intensity ratios HCN(1 − 0)/HCO+(1 − 0) = 1.0 ± 0.2 and HCN(3 − 2)/HCO+(3 − 2) = 1.1 ± 0.3 (Table 5). Evaluating the nuclear ratios in the central pixel, gives very similar values of 1.0 ± 0.2 and 1.3 ± 0.3 for the (1 − 0) and (3 − 2) transitions respectively (Table 6). These ratios place the galaxy in an ambiguous region of the diagnostic diagram, where a starburst, an AGN, or a combination of the two are possible. Depending on the diagnostic used in the literature, either an AGN or a starburst are claimed to power the northern nucleus of Mrk 273. Perhaps a mixture of both activities are present in the central hundred parsecs. Unfortunately, our HCN/HCO+ ratio does not help to elucidate the nature of the Mrk 273 nuclear source. Nevertheless, we note that both lines are significantly affected by opacity (Sect. 4.1), and that a more thorough analysis of this ratio including obscuration effects could further elucidate our results.

Using the Nobeyama Millimeter Array and RAINBOW interferometers at the Nobeyama Radio Observatory6, Imanishi et al. (2006) obtained a flux ratio HCN(1 − 0)/HCO+(1 − 0) > 1.8 in Mrk 273, within a beam of 1ʺ̣9 × 1ʺ̣5. This was interpreted as a signature of an AGN-dominated nucleus. This lower limit is much higher than our values, but we note that the latter authors did not detect HCO+(1 − 0), and that their HCN(1 − 0) and continuum fluxes are 70–80% lower than our values.

4.5.2. HCN/HNC and HCO+/HNC

Some (U)LIRGs, such as Arp 220 and Mrk 231, are extremely bright in HNC, with line intensities almost equal to or even surpassing those of HCN (Huettemeister et al. 1995; Aalto et al. 2002; Aladro et al. 2015). Several scenarios can produce this unusually high line ratio, such as chemical reactions at moderate densities and temperatures (n ∼ 104 − 105 cm−3, Tkin ∼ 50 K), high opacities of HCN with respect to HNC and, very likely, infrared pumping affecting HNC more than HCN (Aalto et al. 2002). The HCN/HNC intensity ratio can reflect the evolutionary stage of starburst regions in galactic centres, with faint HNC emission (i.e. HCN/HNC ≫ 1) being associated with shock-dominated regions, which are common in early starbursts (Aladro et al. 2015).

The brightness temperature ratio between HCN(1 − 0) and HNC(1 − 0) in Mrk 273 is 1.7 (2.1 in the central pixel). Previous observations of these lines with the OSO and SEST single-dish telescopes yielded an integrated intensity ratio ≥4 (Aalto et al. 2002), although that value is based on an upper limit to the HNC line. Taking our result of 1.7 as the minimum value, it is safe to say that this conservative ratio is moderately high, and that Mrk 273 cannot be classified as an HNC-luminous galaxy.

As discussed above (Sect. 4.5.1), the HCN/HCO+ brightness temperature ratio we obtain for Mrk 273 does not allow us to favour a starburst over an AGN-dominated nucleus. The HNC and HCO+ intensities appear to be anti-correlated in (U)LIRGs, as observed by Costagliola et al. (2011) in a fairly large sample of galaxies. From our data, we obtain an HCO+/HNC = 1.8 ± 0.2 (Table 6). If the northern nucleus of Mrk 273 is dominated by a starburst (as claimed by Condon et al. 1991; Majewski et al. 1993; Downes & Solomon 1998), the low HNC abundance would indicate that the gas comes from warm and dense phases in an early stage, and would explain why HCO+ is relatively abundant. However, models by Rodríguez Zaurín et al. (2009) indicate that most of the stellar population in Mrk 273 has an age of 0.7–10 Gyr (although there might be a significant fraction of stars younger than 50 Myr), which challenges an early starburst scenario. In the case of an AGN (as claimed by Vivian et al. 2013; Rodríguez Zaurín et al. 2014; Iwasawa 2018), HNC does not necessarily need to be faint, but HCN could be boosted. However, in the latter scenario one would expect a higher HCN/HCO+ ratio.

4.5.3. An extremely low HCO+/HOC+ ratio

HOC+, the metastable isomer of HCO+, is efficiently formed via the following ion-molecule reactions;(4)

(5)

(6)

and is mainly destroyed by reactions with H2 (Jarrold et al. 1986; Smith et al. 2002; Fuente et al. 2003):(7)

While typical values of the HCO+/HOC+ ratio in Galactic dense molecular clouds range between 300 and 6000 (Apponi & Ziurys 1997), values as low as 50–150 are found in several Galactic and extragalactic PDRs and XDRs (X-ray dominated regions), likely as a consequence of high ionisation rates created by ionisation fields (UV, cosmic rays, and/or X-ray radiation; Fuente et al. 2003; Usero et al. 2004). From our observations, we derive a global brightness temperature ratio HCO+(3 − 2)/HOC+(3 − 2) = 9 ± 4, and a nuclear ratio of 5 ± 3. Such low values have only been found in other ULIRGs hosting extremely compact obscured nuclei, namely IC 860, Zw 049.057, and Mrk 231 (Aalto et al. 2015a,b). The physical and chemical reasons for these low ratios are still not known and merit further study, but high opacities of HCO+ could be responsible.

4.6. A different origin of HOC+ emission

As mentioned in Sect. 3.3, HOC+(3 − 2) peaks at RA (J2000) = 13h44m42ṣ15, Dec (J2000) = 55°53′13ʺ̣45, which corresponds to an offset (0ʺ̣2, − 0ʺ̣05) south-east of the central pixel. This indicates that its emission has a different origin from the rest of the dense gas tracers. Furthermore, the two HOC+ components are blue-shifted in velocity with respect to HCN and HCO+(3 − 2) (Table 1). Figure 17 shows the HOC+ emission in contours plotted over the blue-shifted HCO+ emission integrated between −300 and 0 km s−1. This plot shows that HOC+ does not peak at the same position as the blue-shifted dense gas, and that there is not even detected emission of the species in the central pixel.

thumbnail Fig. 17.

HCO+(3 − 2) blue-shifted intensity colours (in units of Jy km s−1 beam−1) integrated from −300 km s−1 to 0 km s−1 with HOC+(3 − 2) contours (starting from 3σ = 3.3 mJy beam−1) over-plotted with steps of 10 mJy beam−1. The cross at the centre marks the continuum peak. The synthesised beam is shown in the bottom-left corner.

Open with DEXTER

Why is the emission of this species shifted with respect to the others? One possibility is that the central pixels are heavily obscured and HOC+ is completely absorbed there. However, in that case one would expect some kind of symmetric emission around the nucleus, with a ring-like shape, or at least other peaks of emission around the nucleus, which are not seen. This strongly suggests that HOC+ is not peaking at the very centre.

We checked the literature looking for OH masers, supernovae, radio continuum sources, or any source emitting at the HOC+ peak coordinates. High-angular resolution NIR observations conducted by Vivian et al. (2013) resolve the northern nucleus of Mrk 273 into three components, called N1, N2, and N3. Component N1 is the brightest of the three and is associated with the true nucleus (our (0″, 0″) position). N3 is found ∼0ʺ̣15 to the south-east of N1, very near the HOC+ maximum flux density. The peak of the HOC+ emission might correspond to N3, but better astrometry is needed in order to strengthen this association. The nature of N3 is, in any case, not clear, but it does have stronger emission from [FeII] than N1. It could be a separated clump, a compact star cluster, or a supernova remnant. N3 does not appear in radio continuum maps (Carilli & Taylor 2000; Bondi et al. 2005), leaving a supernova remnant as a less likely option. To further address the origin of this species, deeper high-resolution observations of HOC+ would be needed to determine its extent and to more accurately measure its position.

5. Conclusions

We have used the NOEMA interferometer to observe several spectral lines of HCN, HCO+, HOC+, HNC, and HC3N with angular resolutions of (4ʺ̣9 × 4ʺ̣5) and (0ʺ̣61 × 0ʺ̣55) (corresponding to spatial scales of ∼(3.7 × 3.4) kpc and ∼(460 × 420) pc). We also included multiple lines of H2O observed with the Herschel SPIRE and PACS instruments. Our observations, extending from the mm to the FIR regime, allowed us to study the properties of the gaseous disc in the northern nucleus of Mrk 273, as well as the connection between its cold and warm phases. We summarise the main results as follows.

Morphology and kinematics of the disc. The cold and dense gas in the nuclear disc (traced by HCN and HCO+) has two components with decoupled kinematics. The low-excitation gas in the outer parts of the disc extends up to a radius of ∼1.5 kpc and rotates from south-east to north-west (with a PA of approximately −40°), while the more excited dense gas arising from the central star forming region (< 300 pc) is characterised by a north-east to south-west rotation (PA ∼ 70°). This inner disc contains a dynamical mass of 3 × 109M, and a luminosity of LHCN = 3 × 108 K km s−2 pc2. The warm gas, traced by the FIR H2O lines, can also be separated into two components: a warm and very compact core with a radius of ∼50 pc and a temperature of 95 K, and a more extended and relatively cooler component, with R < 300 pc and T = 55 K.

The extended component of the warm gas and the compact component of the cold gas are co-spatial in the inner ∼300 pc. The H2 column densities and dust properties obtained from our water modelling, as well as the line profiles of the dense gas tracers, show that this region is significantly affected by dust obscuration. The blue-shifted emission of the bulk of gas (consistently seen in all observed lines) also indicates that the core is expanding outwards at low velocities (v − vsys ∼ 50–100 km s−1), likely affected by the outflow.

Outflow properties. We detected the cold (sub-mm) and warm (FIR) phases of the Mrk 273 molecular outflow. It is a compact outflow, being expelled to distances of ∼460 pc mostly towards the northern direction, but it reaches high velocities of ∼1000 km s−1. This fast outflow (|v − vsys| > 400 km s−1) has a luminosity of 8 × 107 K km s−1 pc2, and a mass of dense gas . The difference in PA between the major kinematic axis of the inner disc (71° ± 5°) and that of the outflow (10° ± 3°) suggests that the latter is probably driven by the AGN.

Chemistry. We explored the chemistry of Mrk 273 by means of molecular line ratios. The most outstanding ratio is that of HCO+/HOC+. We estimated it to be < 10, one of the lowest values ever measured in any galactic or extragalactic source. The reason for this value, however, is still not clear and should be further studied in detail with the help of chemical models. It is worth noting, however, that the origin of HOC+ is different from the rest of the detected molecular species, since its emission is spatially shifted from the centre.

Regarding the outflow, our non-detection of the high-velocity wind in HCO+, together with the non-detection of the red-shifted outflowing gas either in HCN and HCO+, suggests the possibility of chemical differentiation. However, we note that, despite our high sensitivity, the fast outflow of Mrk 273 is very faint, meaning that deeper observations would be necessary to better probe its chemistry.


3

OBSIDs:1342257290-1342257294.

4

OBSID:1342209850.

5

For the core component, we modified the kλ curve in Fig. 2 of González-Alfonso et al. (2014) in such a way that κ1.3 mm = 0.9 cm2 g−1 of dust, more similar to the value used by Downes & Solomon (1998).

Acknowledgments

This work is based on observations carried out under project numbers W14DD and E16AK with the IRAM NOEMA Interferometer. IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain). The research leading to these results has received funding from the European Union’s Horizon 2020 research and innovation program under grant agreement No 730562 [RadioNet]. This research has made use of NASA’s Astrophysics Data System, and the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. We are grateful to the referee for the fast and constructive report, as well as to the IRAM/NOEMA staff for their help during the observations and data reduction. RA would like to thank Leslie Hunt and Loreto Barcos-Muñoz for the useful discussions about Mrk 273.

References

  1. Aalto, S., Polatidis, A. G., Hüttemeister, S., & Curran, S. J. 2002, A&A, 381, 783 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  2. Aalto, S., Monje, R., & Martín, S. 2007, A&A, 475, 479 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  3. Aalto, S., Garcia-Burillo, S., Muller, S., et al. 2012, A&A, 537, A44 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  4. Aalto, S., Garcia-Burillo, S., Muller, S., et al. 2015a, A&A, 574, A85 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  5. Aalto, S., Martín, S., Costagliola, F., et al. 2015b, A&A, 584, A42 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  6. Aalto, S., Muller, S., Costagliola, F., et al. 2017, A&A, 608, A22 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  7. Aladro, R., Martín-Pintado, J., Martín, S., Mauersberger, R., & Bayet, E. 2011, A&A, 525, A89 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  8. Aladro, R., Viti, S., Bayet, E., et al. 2013, A&A, 549, A39 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  9. Aladro, R., Martín, S., Riquelme, D., et al. 2015, A&A, 579, A101 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  10. Apponi, A. J., & Ziurys, L. M. 1997, ApJ, 481, 800 [NASA ADS] [CrossRef] [Google Scholar]
  11. Armus, L., Charmandaris, V., Bernard-Salas, J., et al. 2007, ApJ, 656, 148 [NASA ADS] [CrossRef] [Google Scholar]
  12. Bondi, M., Pérez-Torres, M.-A., Dallacasa, D., & Muxlow, T. W. B. 2005, MNRAS, 361, 748 [NASA ADS] [CrossRef] [Google Scholar]
  13. Botinelli, L., Fraix-Burnet, D., Gouguenheim, L., Le Squeren, A. M., & Patey, I. 1985, IAU Circ., 4074, 1 [NASA ADS] [Google Scholar]
  14. Carilli, C. L., & Taylor, G. B. 2000, ApJ, 532, L95 [NASA ADS] [CrossRef] [PubMed] [Google Scholar]
  15. Cicone, C., Maiolino, R., Sturm, E., et al. 2014, A&A, 562, A21 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  16. Cole, G. H. J., Pedlar, A., Holloway, A. J., & Mundell, C. G. 1999, MNRAS, 310, 1033 [NASA ADS] [CrossRef] [Google Scholar]
  17. Colina, L., Arribas, S., & Borne, K. D. 1999, ApJ, 527, L13 [NASA ADS] [CrossRef] [PubMed] [Google Scholar]
  18. Condon, J. J. 1997, PASP, 109, 166 [NASA ADS] [CrossRef] [Google Scholar]
  19. Condon, J. J., Huang, Z.-P., Yin, Q. F., & Thuan, T. X. 1991, ApJ, 378, 65 [NASA ADS] [CrossRef] [Google Scholar]
  20. Costagliola, F., & Aalto, S. 2010, A&A, 515, A71 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  21. Costagliola, F., Aalto, S., Rodriguez, M. I., et al. 2011, A&A, 528, A30 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  22. Costagliola, F., Sakamoto, K., Muller, S., et al. 2015, A&A, 582, A91 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  23. Daniel, F., Dubernet, M.-L., & Grosjean, A. 2011, A&A, 536, A76 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  24. Davies, R. I., Tacconi, L. J., & Genzel, R. 2004, ApJ, 613, 781 [NASA ADS] [CrossRef] [Google Scholar]
  25. Downes, D., & Solomon, P. M. 1998, ApJ, 507, 615 [NASA ADS] [CrossRef] [Google Scholar]
  26. Dubernet, M.-L., Daniel, F., Grosjean, A., & Lin, C. Y. 2009, A&A, 497, 911 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  27. Falstad, N., González-Alfonso, E., Aalto, S., & Fischer, J. 2017, A&A, 597, A105 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  28. Fuente, A., Rodríguez-Franco, A., García-Burillo, S., Martín-Pintado, J., & Black, J. H. 2003, A&A, 406, 899 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  29. Gao, Y., & Solomon, P. M. 2004, ApJS, 152, 63 [NASA ADS] [CrossRef] [Google Scholar]
  30. Gao, Y., Gruendl, R. A., Hwang, C.-Y., & Lo, K. Y. 1999, Galaxy Interact. Low High Redshift, 186, 227 [NASA ADS] [CrossRef] [Google Scholar]
  31. García-Burillo, S., Combes, F., Usero, A., et al. 2015, A&A, 580, A35 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  32. González-Alfonso, E., & Cernicharo, J. 1997, A&A, 322, 938 [NASA ADS] [Google Scholar]
  33. González-Alfonso, E., & Cernicharo, J. 1999, ApJ, 525, 845 [NASA ADS] [CrossRef] [Google Scholar]
  34. González-Alfonso, E., Fischer, J., Isaak, K., et al. 2010, A&A, 518, L43 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  35. González-Alfonso, E., Fischer, J., Aalto, S., & Falstad, N. 2014, A&A, 567, A91 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  36. González-Alfonso, E., Fischer, J., Spoon, H. W. W., et al. 2017, ApJ, 836, 11 [NASA ADS] [CrossRef] [Google Scholar]
  37. Graciá-Carpio, J., García-Burillo, S., Planesas, P., Fuente, A., & Usero, A. 2008, A&A, 479, 703 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  38. Griffin, M. J., Abergel, A., Abreu, A., et al. 2010, A&A, 518, L3 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  39. Huettemeister, S., Henkel, C., Mauersberger, R., et al. 1995, A&A, 295, 571 [NASA ADS] [Google Scholar]
  40. Imanishi, M., Nakanishi, K., & Kohno, K. 2006, AJ, 131, 2888 [NASA ADS] [CrossRef] [Google Scholar]
  41. Imanishi, M., Nakanishi, K., Tamura, Y., & Peng, C.-H. 2009, AJ, 137, 3581 [NASA ADS] [CrossRef] [Google Scholar]
  42. Imanishi, M., Nakanishi, K., & Izumi, T. 2016, AJ, 152, 218 [NASA ADS] [CrossRef] [Google Scholar]
  43. Ivison, R. J., Greve, T. R., Dunlop, J. S., et al. 2007, MNRAS, 380, 199 [NASA ADS] [CrossRef] [Google Scholar]
  44. Iwasawa, K. U. V., Mazzarella, J. M., et al. 2018, A&A, 611, A71 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  45. Izumi, T., Kohno, K., Martín, S., et al. 2013, PASJ, 65, 100 [NASA ADS] [Google Scholar]
  46. Izumi, T., Kohno, K., Aalto, S., et al. 2016, ApJ, 818, 42 [NASA ADS] [CrossRef] [Google Scholar]
  47. Janssen, A. W., Christopher, N., Sturm, E., et al. 2016, ApJ, 822, 43 [NASA ADS] [Google Scholar]
  48. Jarrold, M. F., Bowers, M. T., Defrees, D. J., McLean, A. D., & Herbst, E. 1986, ApJ, 303, 392 [NASA ADS] [CrossRef] [Google Scholar]
  49. Kohno, K., Matsushita, S., Vila-Vilaró, B., et al. 2001, The Central Kiloparsec of Starbursts and AGN: The La Palma Connection, 249, 672 [NASA ADS] [Google Scholar]
  50. Krips, M., Neri, R., García-Burillo, S., et al. 2008, ApJ, 677, 262 [NASA ADS] [CrossRef] [Google Scholar]
  51. Lahuis, F., Spoon, H. W. W., Tielens, A. G. G. M., et al. 2007, ApJ, 659, 296 [NASA ADS] [CrossRef] [Google Scholar]
  52. Lindberg, J. E., Aalto, S., Muller, S., et al. 2016, A&A, 587, A15 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  53. Lu, N., Zhao, Y., Díaz-Santos, T., et al. 2017, ApJS, 230, 1 [NASA ADS] [CrossRef] [Google Scholar]
  54. Majewski, S. R., Hereld, M., Koo, D. C., Illingworth, G. D., & Heckman, T. M. 1993, ApJ, 402, 125 [NASA ADS] [CrossRef] [Google Scholar]
  55. Martín, S., Mauersberger, R., Martín-Pintado, J., Henkel, C., & García-Burillo, S. 2006, ApJS, 164, 450 [NASA ADS] [CrossRef] [Google Scholar]
  56. Martín, S., Krips, M., Martín-Pintado, J., et al. 2011, A&A, 527, A36 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  57. Martín, S., Kohno, K., Izumi, T., et al. 2015, A&A, 573, A116 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  58. Martín, S., Aalto, S., Sakamoto, K., et al. 2016, A&A, 590, A25 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  59. Medling, A. M. U. V., Guedes, J., et al. 2014, ApJ, 784, 70 [NASA ADS] [CrossRef] [Google Scholar]
  60. Meijerink, R., Spaans, M., & Israel, F. P. 2006, ApJ, 650, L103 [NASA ADS] [CrossRef] [Google Scholar]
  61. Nardini, E., Risaliti, G., Salvati, M., et al. 2009, MNRAS, 399, 1373 [NASA ADS] [CrossRef] [Google Scholar]
  62. Ott, S. 2010, Astron. Data Anal. Softw. Syst. XIX, 434, 139 [NASA ADS] [Google Scholar]
  63. Pickett, H. M., Poynter, R. L., Cohen, E. A., et al. 1998, J. Quant. Spectr. Rad. Transf., 60, 883 [NASA ADS] [CrossRef] [Google Scholar]
  64. Poglitsch, A., Waelkens, C., Geis, N., et al. 2010, A&A, 518, L2 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  65. Privon, G. C., Herrero-Illana, R., Evans, A. S., et al. 2015, ApJ, 814, 39 [NASA ADS] [CrossRef] [Google Scholar]
  66. Rangwala, N., Maloney, P. R., Wilson, C. D., et al. 2015, ApJ, 806, 17 [NASA ADS] [CrossRef] [Google Scholar]
  67. Rigopoulou, D., Lawrence, A., & Rowan-Robinson, M. 1996, MNRAS, 278, 1049 [NASA ADS] [Google Scholar]
  68. Rodríguez Zaurín, J., Tadhunter, C. N., & González Delgado, R. M. 2009, MNRAS, 400, 1139 [NASA ADS] [CrossRef] [Google Scholar]
  69. Rodríguez Zaurín, J., Tadhunter, C. N., Rupke, D. S. N., et al. 2014, A&A, 571, A57 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  70. Sakamoto, K., Aalto, S., Evans, A. S., Wiedner, M. C., & Wilner, D. J. 2010, ApJ, 725, L228 [NASA ADS] [CrossRef] [Google Scholar]
  71. Savage, C., & Ziurys, L. M. 2004, ApJ, 616, 966 [NASA ADS] [CrossRef] [Google Scholar]
  72. Smith, M. A., Schlemmer, S., von Richthofen, J., & Gerlich, D. 2002, ApJ, 578, L87 [NASA ADS] [CrossRef] [Google Scholar]
  73. Toomre, A. 1964, ApJ, 139, 1217 [NASA ADS] [CrossRef] [Google Scholar]
  74. Vivian, U., Medling, A., Sanders, D., et al. 2013, ApJ, 775, 115 [NASA ADS] [CrossRef] [Google Scholar]
  75. Usero, A., García-Burillo, S., Fuente, A., Martín-Pintado, J., & Rodríguez-Fernández, N. J. 2004, A&A, 419, 897 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  76. Veilleux, S., Rupke, D. S. N., Kim, D.-C., et al. 2009, ApJS, 182, 628 [NASA ADS] [CrossRef] [Google Scholar]
  77. Veilleux, S., Meléndez, M., Sturm, E., et al. 2013, ApJ, 776, 27 [NASA ADS] [CrossRef] [Google Scholar]
  78. Wilson, C. D., et al. 2008, ApJS, 178, 189 [NASA ADS] [CrossRef] [Google Scholar]
  79. Yun, M. S., & Scoville, N. Z. 1995, ApJ, 451, L45 [NASA ADS] [CrossRef] [Google Scholar]

All Tables

Table 1.

Gaussian fit parameters to the emission lines detected with NOEMA.

Table 2.

Gaussian fit parameters to the water lines detected with Herschel.

Table 3.

Global properties of the lines detected with NOEMA.

Table 4.

Results for the two-component modelling of the H2O lines in Mrk 273.

Table 5.

Brightness temperature ratios in Mrk 273 (row/column), evaluated over the entire emission.

Table 6.

Nuclear brightness temperature ratios in Mrk 273 (row/column), evaluated in the central pixel.

Table 7.

Main parameters of the undetected vibrational transitions of HCN and HC3N.

All Figures

thumbnail Fig. 1.

Continuum maps at 3 mm and 1 mm. Contour levels start at a significance of 5σ with respect to the rms measured in both images (rms = 0.07 mJy channel−1 and 0.6 mJy channel−1 for the 3 mm and 1 mm maps respectively). The contour steps are 1 and 3 mJy km s−1 beam−1. The crosses at the centres mark the continuum peaks, which we take as the location of the northern nuclear source (see Sect. 1). The synthesised beams are shown in the bottom-left corner. The colour flux scales are in Jy km s−1 beam−1.

Open with DEXTER
In the text
thumbnail Fig. 2.

Emission lines detected with NOEMA (black histograms) and Gaussian fits (total: red lines; if two components are present, individual components are displayed in blue colour). The velocity resolution is 68 km s−1 in all cases. The labels in the top right corners indicate if the spectra were extracted from the integrated emission (“Integrated”, top panels), from the central pixel (bottom panels), or in the case of HOC+, from the pixel at (0ʺ̣2, −0ʺ̣05).

Open with DEXTER
In the text
thumbnail Fig. 3.

H2O absorption lines (black histograms) detected with Herschel/PACS and Gaussian fits (red lines). The velocity sampling is 20–40 km s−1.

Open with DEXTER
In the text
thumbnail Fig. 4.

H2O lines (black histograms) detected with Herschel/SPIRE, and Gaussians fits (convolved with sinc functions).

Open with DEXTER
In the text
thumbnail Fig. 5.

Spectrum of the central pixel for the (J − J′) = (3 − 2) lines of HCN, HCO+ and HOC+. The baseline of order 0 and the 3 × rms flux (calculated at the final velocity resolution of 68 km s−1) are marked with horizontal dashed lines. The outflow emission at the blue-shifted velocities of HCN is highlighted in yellow.

Open with DEXTER
In the text
thumbnail Fig. 6.

Integrated intensities (moment zero maps). Contour levels for HCN, HCO+ and HNC (1 − 0) go from 0 to 6 Jy kms−1 beam−1 with a step of 0.5 Jy kms−1 beam−1. For HCN (3 − 2) and HCO +(3 − 2) the levels range from 2 to 30 Jy kms−1 beam−1 with a step of 2 Jy kms−1 beam−1. For HOC+(3 − 2), the contours are from 0.4 to 0.6 Jy kms−1 beam−1 with a step of 0.04 Jy kms−1 beam−1. We highlight the different scales between the (1 − 0) and the (3 − 2) lines. The crosses in the centre of each panel indicate the location of the nuclear source. The synthesised beam is shown in the bottom-left corner of each panel. North is up, and east is to the left.

Open with DEXTER
In the text
thumbnail Fig. 7.

Integrated intensity of the outflow seen in HCN (3 − 2) between −1100 km s−1 and −400 km s−1 (yellow region marked in Fig. 5). The cross in the centre marks the location of the nuclear source. The contours start with a 5σ flux and go from 0.8 mJy kms−1 beam−1 to 2.5 Jy kms−1 beam−1 with steps of 0.2 mJy kms−1 beam−1. The magenta dashed ellipses show the regions fitted to the two main components. The synthesised beam is plotted in the bottom-left corner.

Open with DEXTER
In the text
thumbnail Fig. 8.

Velocity fields (moment 1 maps). The coloured velocity scale (right axis) is in km s−1. The step in contours is 20 kms−1 for all lines. We note the blue-shifted velocities of HOC+(3 − 2). The crosses in the centre of each panel indicate the position of the nuclear source. The beam is shown in the lower left corner of each panel. North is up, and east is to the left. The dashed lines in the HCN plots indicate the cut for the p − v diagrams shown in Fig. 9.

Open with DEXTER
In the text
thumbnail Fig. 9.

Position-velocity (p-v) maps of HCN and HCO+. The cuts along the axes of rotation are shown by the dashed lines in Fig. 8. Contour steps of the (1 − 0) maps go from 0.9 mJy beam−1 (3σ) to 7.9 mJy beam−1 with steps of 1 mJy beam−1. For the (3 − 2) maps, contours go from 3.9 mJy beam−1 (3σ) to 59 mJy beam−1 with steps of 5 mJy beam−1.

Open with DEXTER
In the text
thumbnail Fig. 10.

Velocity dispersions. Contours go from 0 to 500 kms−1 with steps of 50 kms−1 for all lines. The crosses in the centre of each panel indicate the position of the nuclear source. The synthesised beam is shown in the bottom-left corner of each panel. North is up and East is to the left.

Open with DEXTER
In the text
thumbnail Fig. 11.

Channel-velocity maps of HCN (1 − 0) in the velocity range [−500, 500] km s−1 with steps of 50 km s−1. Contours go from 1.5 mJy beam−1 (5σ) to 9.5 mJy beam−1 with a spacing of 2 mJy beam−1. The synthesised beam is plotted in the bottom-left corner. North is up and east is to the left.

Open with DEXTER
In the text
thumbnail Fig. 12.

Channel-velocity maps of HCO+(1 − 0) in the velocity range [−500, 500] km s−1 with steps of 50 km s−1. Contours go from 1.5 mJy beam−1 (5σ) to 9.5 mJy beam−1 with a spacing of 2 mJy beam−1. The synthesised beam is plotted in the bottom-left corner. North is up and east is to the left.

Open with DEXTER
In the text
thumbnail Fig. 13.

Channel-velocity maps of HCN (3 − 2) in the velocity range [−500, 500] km s−1 with steps of 50 km s−1. Contours go from 6.5 mJy beam−1 (5σ) to 56 mJy beam−1 with a spacing of 10 mJy beam−1. We highlight the significantly smaller spatial scale relative to the channel map of the (1 − 0) line shown in Fig. 11. The synthesised beam is plotted in the bottom-left corner. North is up and east is to the left.

Open with DEXTER
In the text
thumbnail Fig. 14.

Channel-velocity maps of HCO+(3 − 2) in the velocity range [−500, 500] km s−1 with steps of 50 km s−1. Contours go from 6.5 mJy beam−1 (5σ) to 56 mJy beam−1 with a spacing of 10 mJy beam−1. We highlight the significantly smaller spatial scale relative to the channel maps of the (1 − 0) line shown in Fig. 12. The synthesised beam is plotted in the bottom-left corner. North is up and east is to the left.

Open with DEXTER
In the text
thumbnail Fig. 15.

Comparison between the observed (black symbols) and modelled (coloured symbols and lines) H2O fluxes in Mrk 273 within panel a: PACS and panel b: SPIRE. The model includes two components: the core (in blue), which accounts for most absorption lines observed within PACS, and the inner disc (in green), which dominates the emission of the sub-mm lines with Eupper < 400 K observed with SPIRE (see Table 4). Red colours and symbols indicate total modelled fluxes. The numbers at the bottom of panels a and b indicate rounded-up transition wavelengths in μm. Panel c: the SED of Mrk 273, including the Spitzer/IRS spectrum, Herschel/PACS and SPIRE continuum data from observations of both H2O and OH lines, sub-mm data at 800 and 880 μm (Rigopoulou et al. 1996; Wilson 2008), and our measured flux density at 1 mm (starred-red symbol), is compared with the prediction of our composite model.

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In the text
thumbnail Fig. 16.

Left: sketch of the blue-shifted molecular outflow in the northern nucleus of Mrk 273. The dense and warm outflowing gas found in our HCN and H2O data, and its properties, is compared to the diffuse outflow found by Cicone et al. (2014) by observing CO (1 − 0). We note that the mass loss rate is calculated as Ṁ = v (MOF/ROF) (see Sect. 5.5 in González-Alfonso et al. 2017), which is a factor of three lower than the value given by Cicone et al. (2014). The position angles of the outflow (10°) and the inner disc (70°) are represented. Right: a face-on sketch of the three disc components identified in our data (outer disc, inner disc, and core), which are plotted in scaled sizes. The decoupled kinematics of the outer and inner discs are represented by the blue and red colours, which depict the orientation of the blue-shifted and red-shifted rotating gas. The rotation direction is indicated by the curved arrows. The intensity of the colours illustrate the velocities; the inner disc rotates at a higher speed than the outer disc. The stars in the inner disc aim to show the region where most of the starburst is located. The radial arrows represent the low-velocity expansion of the core.

Open with DEXTER
In the text
thumbnail Fig. 17.

HCO+(3 − 2) blue-shifted intensity colours (in units of Jy km s−1 beam−1) integrated from −300 km s−1 to 0 km s−1 with HOC+(3 − 2) contours (starting from 3σ = 3.3 mJy beam−1) over-plotted with steps of 10 mJy beam−1. The cross at the centre marks the continuum peak. The synthesised beam is shown in the bottom-left corner.

Open with DEXTER
In the text

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