The Molecular Gas in the NGC 6240 Merging Galaxy System at the Highest Spatial Resolution

Table 1 details the properties of each observation considered in this study. Despite not being used in the final imaging for the reasons listed below, the long baseline execution block (EB) Xac5575/X8a5f is also listed. All of the observations described, except Xb4da9a/X69a, share the same spectral setup. The total ALMA bandwidth of 7.5 GHz was divided into four spectral windows (SPWs) of 1875 MHz each and a spectral resolution of 1.953 MHz. In all observations, one of these SPWs was specifically tuned to study the redshifted ¹²CO(2–1) transition, centered at a frequency of 225.1 GHz. The remaining ones were used to trace the millimeter continuum centered at frequencies of 227.5, 239.9, and 241.9 GHz, with the second one partly covering the CS(5–4) line. The Xb4da9a/X69a EB covered the ¹²CO(2–1) and CS(5–4) transitions with two SPWs centered at 224.9 and 239.8 GHz with spectral resolutions of 1.953 and 3.906 MHz, respectively. The remaining two SPWs traced the continuum at 221.9 and 237.9 GHz with spectral resolutions of 7.813 MHz.

For the data analysis, the Common Astronomy Software Applications (CASA, v5.4; McMullin et al. 2007) was used. In general, the standard ALMA data reduction²⁵ approach was pursued, with some specific differences. More details about the array configuration properties, observations, processing method, and imaging parameters are reported in the Appendix.

Figure 1 shows the ¹²CO(2–1) emission, together with the optical imaging provided by the HST using the Wide Field Planetary Camera 2 (WFPC2) instrument. As described above, by combining low- and high-resolution observations, we are able to recover all of the emission out to scales of 30''. Indeed, the total measured flux is 1163 ± 28 Jy km s⁻¹ in the ALMA field of view of 25'' diameter, which accounts for 95% of the total CO(2–1) when compared to the Tacconi et al. (1999) measurements. We can further compare this total flux to single-dish observations. An unresolved, ¹²CO(2–1) flux of 1492 ± 253 Jy km s⁻¹ was reported by Greve et al. (2009), implying that our ALMA observations recover 67%–94% of the single-dish flux when considering the relatively large error bars of the Greve et al. (2009) measurements. As can be seen in Figure 1, a significant concentration of the molecular gas is found in the central regions, consistent with findings based on previous observations (e.g., Tacconi et al. 1999; Engel et al. 2010). However, extended molecular emission can be detected up to ∼10'' (∼5 kpc) away from the nuclear region. The bulk of the molecular gas emission in the nuclear region of the system appears to be directly connected to the material found between the nuclei.

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This molecular CO emission does not spatially coincide with the submillimeter continuum observed by ALMA at observed frame 235 GHz, as shown by the magenta contours in the left panel of Figure 1 and in more detail in Figure 2. Similarly, the bulk of the ¹²CO(2–1) emission does not appear to overlap with the stellar light traced by the optical and near-IR emission either, which is mostly centered around each nucleus (Max et al. 2005; Engel et al. 2010). This could be naturally explained by the effects of dust extinction in the nuclear regions. Indeed, Max et al. (2005, 2007) and Müller-Sánchez et al. (2018) reported the presence of significant dust lanes in the center of the NGC 6240 system. For our work, the nuclear positions were obtained from the VLBI observations of this source, as reported by Hagiwara et al. (2011). Specifically, the nuclear positions used are 16^h52^m58.924, 02°24'04776 and 16^h52^m58.890, 02°24'03350 for the northern and southern ones, respectively. The typical error in these positions is 0.6 mas. Only a relatively small fraction of the molecular gas, compared to the total amount in the central region of the system, is found inside the sphere of influence of each SMBH, which was computed by Medling et al. (2015) to have radii of 235 ± 8 and 212 ± 9 pc for the northern and southern nuclei, respectively, and are hence fully resolved in the ALMA cube. This can be used to improve existing measurements of the SMBH mass in each nucleus by separating the SMBH and molecular gas contributions to the total enclosed nuclear mass, which in turn indicates that these are closer to the usual M–σ correlation than previously indicated (Medling et al. 2019). Recently, Kollatschny et al. (2020) claimed the detection of a third nucleus in this system from optical integral field spectroscopy observations close to the southern nucleus at a position of 16^h52^m58.901, 02°24'0288. We do not detect any significant ¹²CO(2–1) emission at this position or its surroundings, in stark contrast with the two previously established nuclei.

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In contrast to previous submillimeter observations at lower spatial resolutions (Tacconi et al. 1999; Saito et al. 2018), the molecular gas in the nuclear region does not appear to follow a smooth distribution. Instead, substantial clumpiness can be found on the stream that connects both nuclei, as can be seen in Figure 1. Similar structures in the spatial distribution of the molecular emission, in this case from ¹²CO(3–2) observations, have been previously presented by U et al. (2011) and Wang et al. (2014). These clumpy regions are natural sites where vigorous star formation should take place. However, according to high-resolution, ∼01, near-IR observations at 2.2 μm obtained with the Very Large Telescope (VLT)/SINFONI integral field spectrograph, reported by Engel et al. (2010), the bulk of the nuclear stellar mass appears to precede the merger and is spatially concentrated around each nucleus, certainly much more than the molecular gas traced by the ALMA observations.

In order to take advantage of the superb spatial resolution of the ALMA observations, the focus of this work is mostly on the relatively small-scale, subarcsecond structures. That being said, we can place these in the context of more extended structures previously reported from CO observations. Specifically, Feruglio et al. (2013a) reported the detection of CO(1–0) emission extending up to 20'' (∼10 kpc) from the central region of the system in both the eastward and southwestward directions. The eastward extended structure is also detected in our ALMA data at velocities of ∼−300 to −400 km s⁻¹ and reaching up to ∼12'' at the edge of the field of view of our data cube; as such, we cannot verify a larger extension. We also detect the southwestward structure at a velocity of ∼−100 km s⁻¹. The extension of this structure is ∼13'' from the southern nucleus, where it reaches the edge of the field of view, as in the previous case. Hence, we confirm the detection of extended CO emission previously reported by Feruglio et al. (2013a). We can further identify an additional extended structure reaching up to ∼7'' to the northwest of the northern nucleus at velocities of ∼100–200 km s⁻¹. A detailed characterization of these large-scale structures is beyond the scope of this paper.

As described in Section 2, some of the ALMA SPWs were used to study the continuum emission at ∼235 GHz, corresponding to ∼1.25 mm, which we show in Figure 2. As can be seen there, and contrary to what was observed for the ¹²CO(2–1) transition that mostly traces the molecular gas, the millimeter continuum emission arising from dust is mostly concentrated around each nucleus. We measure continuum flux densities at an observed-frame frequency of 235 GHz of 9.92 ± 0.16 and 3.50 ± 0.037 mJy inside the corresponding SMBH spheres of influence for the southern and northern nuclei, respectively. Considering now a much larger aperture with a diameter of 22'', we measure a continuum flux of 13.0 ± 0.6 mJy. Hence, we can conclude that in this source, most of the total millimeter continuum emission detected in our ALMA data is directly associated with the two nuclear regions. Furthermore, as can be seen in Figure 2, the emission is strongly concentrated and marginally resolved. Similar to the nondetection of CO(2–1) emission discussed in the previous section, at the flux limit of our observations, we do not detect any submillimeter continuum emission at the location of the third nucleus reported by Kollatschny et al. (2020).

The continuum fluxes reported here are significantly higher than those reported by Tacconi et al. (1999) at 228 GHz, by 2.1 times for the southern nucleus and 3.5 times for the northern one, using similar apertures ∼08 in diameter. This discrepancy can likely be explained by sparse uv coverage by our long baseline observations combined with their poor phase stability, which particularly affects flux measurements on relatively small scales, ∼1'', after integrating over a relatively large wavelength range. Indeed, when continuum fluxes are measured using only the compact and intermediate array configurations described in Appendix A.1, we obtain values of 5.19 ± 0.06 and 1.73 ± 0.03 mJy for the southern and northern nuclei, respectively, roughly consistent with the Tacconi et al. (1999) values. The discrepancy with the maps containing the long baselines is much smaller if we use much smaller apertures, <05, as was done by Medling et al. (2019), or much larger ones. Indeed, on a 22'' diameter aperture, we measure a continuum flux of 10.1 ± 0.13 mJy, again showing that most of the millimeter continuum emission in this system is associated with the two nuclei. Therefore, in what follows, we only use these latter fluxes, as clearly better uv coverage and phase stability are required to measure continuum fluxes at relatively small scales. The continuum emission map presented in Figure 2 is, however, not significantly affected, as this "extra" flux has a relatively low surface density. Similarly, by comparing with previously reported ¹²CO(2–1) fluxes at a range of scales and resolutions, we concluded that our line measurements have not been significantly affected. Hence, this issue is only relevant for continuum studies that require one to integrate flux densities over a large wavelength range. Note that the use of feathering on a per-channel basis when retrieving the ¹²CO(2–1) emission map prevented this low surface density from being parsed to the final combined map.

Previously, Nakanishi et al. (2005) presented continuum observations at frequencies of 87 and 108 GHz, with synthesized beam sizes of ∼2'' and ∼4'' and total measured fluxes of 16.6 and 10.8 mJy at each frequency, respectively. Even at lower spatial resolution, both Tacconi et al. (1999) and Nakanishi et al. (2005) were able to identify an offset between the continuum and the line emission, which is further corroborated by our ALMA data. The spectral shape in these data is given by a power law with index α = −0.81, consistent with synchrotron emission from nonthermal sources such as AGNs or supernova remnants. Our ALMA observations reveal that most of the continuum emission arises from the surroundings of each SMBH and is significantly brighter than the extrapolation of the spectrum at low frequencies, as presented in Section 4.2, thus strongly suggesting that this radiation is dominated by dust heated by AGN activity.

The existing 880 μm single-dish dust continuum measurement of NGC 6240 is F₈₈₀ = 133 ± 40 mJy (Wilson et al. 2008). Adopting a dust emissivity index β = 1.8 (Planck Collaboration et al. 2011), this predicts a flux density at 235 GHz of 32 ± 10 mJy. Hence, our observations recover approximately 41% of the expected continuum flux. If we assume this missing flux is concentrated in a central 2'' diameter aperture, the predicted flux per beam is 5 μJy, below our detection threshold. The missing flux is hence consistent with being due to insufficient surface brightness sensitivity in these continuum observations.

The ALMA observations also yield information about the kinematics of the molecular gas. Figure 3 shows the velocity map for the central region of NGC 6240 as traced by the ¹²CO(2–1) emission. In contrast to previous claims (Bryant & Scoville 1999; Tacconi et al. 1999; Engel et al. 2010), there is no clear evidence for a rotating disk in the internuclear region. In fact, no well-defined structures are obviously visible in the moment 1 map. A similar conclusion, albeit using lower-resolution data, was reached before by Cicone et al. (2018). In this work, they reported that the typical butterfly pattern, a smoking gun for a rotating disk, could not be seen; it is not present in our data either. Instead, they concluded that the kinematics in the nuclear region were dominated by outflows. The higher-resolution ALMA data appear consistent with a significant high-velocity outflow, described in more detail in Section 4.3, visible directly south of the northern nucleus. A tentative velocity gradient can also be observed between the two nuclei. Indeed, such a velocity gradient, of ∼100 km s⁻¹, was previously reported based on subarcsecond ¹²CO(3–2) observations by Iono et al. (2007) and U et al. (2011). This gradient was also found in the H₂ observations of this region reported by Müller-Sánchez et al. (2018), who interpreted it as a perturbed turbulent rotational disk. To facilitate the visualization of the overall kinematical structure of the cold molecular gas in the nuclear region of NGC 6240, Figure 4 presents a snapshot of an animation available in the electronic version of this paper exploring the three-dimensional spatial and velocity cube traced by our ALMA observations.

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In order to better understand the kinematic structure of the molecular gas stream connecting the two nuclei, Figure 5 presents a position–velocity (p–v) diagram along a slit connecting the two nuclei. For illustration, we overplot the velocity field predictions of the respective stellar disk only and stellar disk plus black hole in the northern and southern nuclei using parameters from Medling et al. (2011, 2015) and assuming, for simplicity, that the PA and inclination of the inner gas disks around the north and south SMBHs follow the same kinematics as the larger-scale stellar disk.²⁶ First, near the southern nucleus, located inside the sphere of influence of the southern SMBH, we can see a strongly localized gradient over ∼02 spanning a range of approximately ±200 km s⁻¹. The CO velocities here do not follow the rotating stellar disk model posited by Medling et al. (2015). This gas thus either follows a different orientation or projection or traces outflowing nuclear gas (in which case, the near side of the southern stellar disk (S) is to the northeast). Interpreting the northeast side of the southern galaxy as the near side would be consistent with the model of Baan et al. (2007), in which the southern stellar disk is behind the northern one (N). Extrapolating this further, one could posit that the far side of the N stellar disk is to the southeast, as this overall orientation would most easily explain the observed internuclear bridge of gas. At offsets of 01–06 northeast of the southern nucleus, the CO roughly follows the rotating stellar disk model of Medling et al. (2011), except for the high-dispersion region at offset ∼03 north of the S nucleus. This structure is by far the most prominent feature seen in the p–v diagram, as it includes the highest flux peak for the internuclear ¹²CO(2–1) emission and extends over 400 km s⁻¹ in velocity space. It is strongly localized spatially, spanning roughly ∼02. This high-dispersion region is located at an offset where the slit passes very close to the radio supernova RS2 (which likely traces a starbursting region), detected at milliarcsecond-scale resolution in the radio by Hagiwara et al. (2011), and its peculiar molecular kinematics are discussed further below. At offsets between 09 and 15 north of the S nucleus, the gas roughly follows the predictions of the rotating northern stellar disk of Medling et al. (2015), though there are extra wiggles and a lack of Keplerian rotation inside the posited sphere of influence of the N nucleus black hole. Finally, at offsets between 05 and 08 north of the S nucleus, we see a clear transition region as the gas connects one of the claimed disks to the other. In summary, the gas bridge between the nuclei appears to connect them via the near side of the (behind) southern stellar disk and the far side of the (in front) northern one, and while some of the gas may be roughly consistent with the expected dynamics of the two nuclear core disks, there are clear streaming kinematics in the arm plus outflow signatures (in the S nucleus and near RS2) and a lack of Keplerian rotation near the N nucleus.

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Figure 6 illustrates the projected geometry of the inner stellar disks and the CO emission, following the interpretation from Baan et al. (2007), based on H i absorption, placing the N disk in front of the S one. In this case, the relatively continuous velocity structure of the gas as it transitions between the two stellar disks (Figure 5) suggests that the disks intersect each other; i.e., the northeast of the S disk is the near side, and the southeast of the N disk is the far side. This is what we present in Figure 6. In this case, both the N and S disks are rotating counterclockwise on the sky. The east–west extension in CO in the area between the two nuclei likely traces the area where both disks actually intersect (rather than just overlap to our sight line). In the Baan et al. (2007) model, the radial distance between the two nuclei is ∼3.3 kpc, so the stellar disks illustrated in the figure would overlap by line-of-sight projection but not actually physically intersect. If this radial distance was decreased to <2.4 kpc, the two illustrated stellar disks would actually intersect over the ridgeline at which we see the maximum east–west extension of the CO emission.

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There is significant evidence for symmetric superwind-driven outflows from the south nucleus (Heckman et al. 1990; Tecza et al. 2000; Baan et al. 2007). Our high-resolution molecular kinematics traces multiple outflow regions and shells. Figure 7 shows a p–v diagram that crosses both the south nucleus and the location of the radio supernova RS2. As in Figure 5, we clearly see high-velocity dispersions around the location of RS2 (dotted vertical line). Further, even higher velocities are seen to the northeast of RS2 (larger positive offsets in the p–v diagram). For illustration, we have overplotted the expected observed velocities from expanding spherical shells of gas with the expansion center at the southern nucleus (red) and the RS2 position (green). These are illustrative "toy" models fitted by eye; the true picture is certainly more complex, as the molecular gas is likely swept up from the internuclear bridge, which is likely highly inclined to our line of sight (see Section 1), and the shell morphology will depend strongly on the surrounding gas distribution with which it interacts. Further, the significantly different velocities required to explain the different shells in our toy model imply either that the periodic outflows are weakening over time or that we have ignored projection factors and local systemic velocities. If the shells are comprised of gas swept up from the internuclear bridge, then the inclination of this bridge to the line of sight and the fact that the bridge's systemic velocity gets more redshifted with larger offsets from the southern nucleus will both help to decrease the true difference in outflow velocities between shells. In any case, the dynamic timescales of these shells are significantly less than the estimated age of 10 Myr for the nuclear starburst (Tecza et al. 2000). As seen in Figure 7, the observed shells can be equally well explained as originating in either the southern nucleus or (near) RS2. However, there are two arguments that support an origin of the outflows from a region close to RS2. (1) The innermost shell extends over more than 180° in azimuth with a center very close to (but ∼002 south of) RS2. If the outflow were driven by the south nucleus, we would expect the arc to be present only on the far side of RS2 (with respect to the southern nucleus). (2) If the shells originated in the southern nucleus, one would expect to see high-velocity arcs over a larger range of offsets from the southern nucleus. What we observe, however, is that these are found at offsets consistent with them being driven out of the gas-rich region near RS2. Given that RS2 likely traces a more extended starburst region, a center slightly offset from RS2 is not unexpected.

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The unusual kinematics in the sphere of influence of the south nucleus also deserves comment. Here, as in Figure 5, the velocities within ± 01 of the southern nucleus do not follow those expected from the rotating model. Hagiwara et al. (2011) detected the young radio supernova RS1 35 mas southwest of the south nucleus. We note that the velocity pattern seen near the south nucleus appears symmetric about a negative offset (consistent with a shift toward the position of RS1) and a positive recessional velocity (consistent with the kinematics of the internuclear bridge) with respect to the southern nucleus. A detailed analysis of the outflows driven by RS1, RS2, S, and N are deferred to a forthcoming work.

We now analyze the kinematic properties of the molecular gas at larger scales by coarsely sampling the ALMA ¹²CO(2–1) SPW in 26 km s⁻¹ bins. Figure 8 shows the detected ¹²CO(2–1) emission at four specific velocities relative to the systemic velocity of the system, defined to be that of the southern nucleus: −148.2 (blueshifted stream), 7.9 (systemic), 164.0 (redshifted stream), and 528.1 (outflowing material) km s⁻¹. Most of the blueshifted material is found near the southern part of the stream, while the redshifted emission is seen near the northern end. This suggests that the stream can be considered as a dynamic structure that is interacting with and evolving relative to the two nuclei.

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The velocity dispersion, in particular relative to the kinematical structure of the system, provides important clues about the dynamical state of the molecular gas in the nuclear regions of this system. Previous observations, such as those presented by Feruglio et al. (2013b), found that at lower, ∼1'', spatial resolutions, velocity dispersions fell within the range ∼150–300 km s⁻¹ but reached up to ∼500 km s⁻¹ in the region between the nuclei and around the northern center. Our ALMA observations, shown in Figure 9, find consistent results, with typical velocity dispersion values in the region between the nuclei of ∼130–150 km s⁻¹ and higher values, reaching ∼250 km s⁻¹, mostly along the edges of the internuclear CO emission. In the northern nucleus, using a radius of ∼200 pc, we find an average velocity dispersion of 184 km s⁻¹, while for the southern one, in a similar area, we find an average velocity dispersion of 155 km s⁻¹. These values are significantly lower than those reported by Feruglio et al. (2013b) and highlight how the spatial resolution of the two observations, ∼1'' versus ∼003, plays a crucial role in the interpretation of the gas state near the SMBHs. Indeed, SMA observations of the ¹²CO(3–2) line carried out by U et al. (2011) at a resolution of ∼03 similarly find a modest velocity dispersion for the molecular gas, fully consistent with what is reported here. Comparable values were also found from ALMA observations of the [C i](1–0) line emission reported by Cicone et al. (2018).

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Albeit lower than previously reported, these still relatively high velocity dispersion values imply v/σ ∼ 1 and hence suggest that the kinematics of the internuclear gas is dominated by turbulence, possibly associated with the presence of significant shocks in the internuclear molecular gas. In addition, and as presented by U et al. (2011), the velocity dispersion distribution is not connected to the overall velocity gradient between the nuclei. This latter fact would indicate that the bulk of the molecular gas, traced by the CO emission, is not undergoing organized circular movements, as proposed previously (e.g., Engel et al. 2010), but instead is consistent with being dynamically supported by turbulence. This seems natural if this is indeed a transient structure created by the merger, which, after a relatively short time (compared to the duration of the merger), may (at least partially) collapse around the new coalesced center of the system, as was previously discussed by Engel et al. (2010) and others and predicted by numerical simulations (Bournaud et al. 2011).

While H₂ is the most abundant molecule in galaxies, it cannot be directly observed in emission due to its lack of dipolar rotational transitions at relatively low excitation temperatures (Bolatto et al. 2013, and references therein). In contrast, the CO molecule can be easily excited even in cold environments and is relatively abundant, making it an important tracer of the molecular gas contents of a galaxy. In order to convert from observed ¹²CO(2–1) line flux to total molecular gas mass, we follow the procedure described by Solomon & Vanden Bout (2005), as was done previously by Treister et al. (2018). Specifically, the line luminosity is computed from the observed line flux as

$$\begin{eqnarray*}&&{L}_{\mathrm{CO}}^{{\prime} }=3.25\times {10}^{7}{S}_{\mathrm{CO}}{\rm{\Delta }}v{\nu }_{\mathrm{obs}}^{-2}{D}_{L}^{2}{(1+z)}^{3}.\end{eqnarray*}$$

From the line luminosity, we estimate the molecular hydrogen mass as

$$\begin{eqnarray*}&&M({{\rm{H}}}_{2})={\alpha }_{\mathrm{CO}}{L}_{\mathrm{CO}}^{{\prime} },\end{eqnarray*}$$

where α_CO is the CO-to-H₂ conversion factor. Then, by incorporating the He contribution into this factor, we obtain the total molecular gas mass. The value of α_CO and its possible variations for different galaxies, and even regions inside a galaxy, are topics still extensively debated in the literature (e.g., Bolatto et al. 2013). The value typically assumed for galaxies like the Milky Way is 4.6 M_⊙ (K km s⁻¹ pc²)⁻¹ (Bolatto et al. 2013). However, in extreme environments, such as those observed in (U)LIRGs, lower α_CO values of ∼0.6–1 M_⊙ (K km s⁻¹ pc²)⁻¹ have been reported (e.g., Downes & Solomon 1998; Tacconi et al. 2008; Yamashita et al. 2017; Herrero-Illana et al. 2019), likely associated with higher velocity dispersions (Papadopoulos et al. 2012). For a source as complex as NGC 6240, it is dangerous to assume a single value of α_CO, as previously argued by Engel et al. (2010). Previously, for this galaxy, Tunnard et al. (2015) reported a very broad range of α_CO values spanning up to 3 orders of magnitude, from ∼0.1 to ∼100 M_⊙ (K km s⁻¹ pc²)⁻¹ (see their Figure 10, left panel). However, the reported range in the global value of α_CO is much narrower, finding α_CO = 1.5${}_{1.1}^{7.1}$ M_⊙ (K km s⁻¹ pc²)⁻¹. In parallel, using combined observations of HCN, CS, and HCO⁺ transitions, Papadopoulos et al. (2014) estimated a conversion factor of α_CO = 2–4 M_⊙ (K km s⁻¹ pc²)⁻¹. More recently, based on ALMA observations at lower spatial resolutions than those presented here, Cicone et al. (2018) reported a mean value of α_CO = 3.2 ± 1.8 M_⊙ (K km s⁻¹ pc²)⁻¹ for the systemic components and α_CO = 2.1 ± 1.2 M_⊙ (K km s⁻¹ pc²)⁻¹ for the outflowing material (presented in their Table 2). We adopt the Cicone et al. (2018) α_CO conversion factor for our calculations, together with their values of r₂₁ = 1.0 ± 0.2 for the systemic component and 1.4 ± 0.3 for the outflows, where r₂₁ is defined as the ratio between the ¹²CO(2–1) and ¹²CO(1–0) luminosities. When comparing to previous results, it is important to consider that fiducial values of α_CO = 1 M_⊙ (K km s⁻¹ pc²)⁻¹ were typically assumed (e.g., Engel et al. 2010).

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Table 2.  The Self-calibration (Selfcal) Results on the Check Source for Different Precalibration Strategies

WVRsc	Flagged Antennae	Peak	1 −	Integ.	1 −	Integ./Peak	Off
[Y/N]		(mJy)	(%)	(mJy)	(%)		(mas)
Y	DA61	132 ± 4		190 ± 10		1.4 ± 0.1	14
		261.9 ± 0.2	50 ± 1	261.3 ± 0.3	28 ± 2	0.998 ± 0.001	0
N	DA61	132 ± 5		190 ± 10		1.5 ± 0.1	13
		261.9 ± 0.2	50 ± 1	261.3 ± 0.3	27 ± 2	0.998 ± 0.001	0
Y	High phase	167 ± 4		196 ± 9		1.18 ± 0.06	15
		262.1 ± 0.2	36.3 ± 0.6	262.2 ± 0.4	25 ± 1	1.000 ± 0.002	0

Note. The first and second sets of rows consider water vapor radiometer (WVR) scaling being applied (Y) or not (N), respectively. The third set of rows shows the results when antennas with high phase rms (σ_ϕ > 1 rad or 57°) are flagged in addition to DA61 (see text for more details). In each set of rows, the first and second refer to the pre- and post-selfcal observed values, respectively. The decorrelation factors (in %) between the selfcal'ed and not selfcal'ed values are shown in the second row in each set. The last column shows the offset between the source and the phase centers in milliarcseconds.

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We adopt three distinct regions: the spheres of influence of the northern and southern SMBHs and a 1'' diameter aperture centered equidistantly along the line connecting the two nuclei (16^h52^m58.896, 02°24'0396). The velocity-integrated ¹²CO(2–1) flux measurements are 34.3 ± 3.7, 142.1 ± 3.1, and 389 ± 11 Jy km s⁻¹, respectively. The central 1'' flux density is similar, albeit ∼25% smaller, to the one reported by Tacconi et al. (1999), which might indicate that our interferometric observations may have resolved out a small fraction of the internuclear molecular mass; however, the lower-resolution value of Tacconi et al. (1999) likely includes substantial flux from the nuclei, at least partially accounting for the 25% difference. The fluxes for the two nuclei are fully consistent with those measured previously by Engel et al. (2010). As presented in Section 3, considering a much larger aperture of 25'' in diameter that should include most of the source ¹²CO(2–1) emission, we measure a flux of 1163 ± 28 Jy km s⁻¹, consistent with the 1220 Jy km s⁻¹ reported by Tacconi et al. (1999) and the ∼1000 Jy km s⁻¹ inferred from the ¹²CO(1–0) observations by Solomon et al. (1997).

Using the expressions derived above, these line fluxes correspond to gas masses of 7.4 × 10⁸ M_⊙ for the northern nucleus, 3.3 × 10⁹ M_⊙ for the southern nucleus, and 8.6 × 10⁹ M_⊙ for the central region. This is consistent with the values previously derived by Engel et al. (2010) and Feruglio et al. (2013b), once the differences in the assumed values of α_CO are accounted for. Finally, from the computed total flux in the 25'' diameter aperture, we derive a mass of 2.6 × 10¹⁰ M_⊙. This is fully consistent with the value of (2.1 ± 0.5) × 10¹⁰ M_⊙ reported by Cicone et al. (2018) in a central 12'' × 6'' region.

The total molecular gas mass can also be estimated from the warm molecular gas, as presented by Dale et al. (2005), even considering that the latter is not a direct tracer of the cold molecular gas. Specifically, for our comparison, we can use the H₂1–0 S(1) line luminosity at rest frame 2.12 μm. We use the empirically calibrated relation ${L}_{(1\mbox{--}0){\rm{S}}(1)}/{M}_{\mathrm{gas}}$ = 2.5 × 10⁻³, as measured specifically for NGC 6240 by Müller Sánchez et al. (2006). This luminosity can be estimated from the VLT/SINFONI observations of NGC 6240 presented by Engel et al. (2010). Using a nuclear 2'' diameter aperture, we estimate an ${L}_{(1\mbox{--}0){\rm{S}}(1)}$ luminosity of 2.8 × 10⁷ L_⊙, or, equivalently, 1.1 × 10¹⁰ M_⊙, and hence fully consistent with the values derived here based on the ¹²CO(2–1) observations. The comparison of the results obtained using the empirically calibrated ${L}_{(1\mbox{--}0){\rm{S}}(1)}/{M}_{\mathrm{gas}}$ ratio with the new ALMA observations presented in this paper indicates that we can use the H₂1–0 S(1) line to make at least an approximate estimate of the total (cold) molecular gas mass.

In the past, most studies of the ISM in the molecular gas phase have relied on the rotational transitions of the CO molecule to trace H₂ gas. However, as shown by Scoville et al. (2014, 2016), measurements of the dust continuum emission at the long-wavelength Rayleigh–Jeans (RJ) end can also provide accurate estimates of the ISM mass in a small fraction of the required observing time with ALMA. At these long wavelengths, the dust emission is thought to be optically thin; hence, the observed continuum flux density is directly proportional to the mass once the dust opacity coefficient and dust temperature are established. It is important to point out that this estimation depends on the assumption that the continuum emission at these wavelengths is dominated by thermal dust radiation, which might not be the case in the central region of NGC 6240. Evidence in this direction was presented by Tacconi et al. (1999) and later by Nakanishi et al. (2005) based on the observed spectral slope, which is consistent with that expected for synchrotron radiation. In addition, Meijerink et al. (2013) argued that the CO-to-continuum luminosity ratio in this galaxy is significantly larger than the value observed in other similar systems, which might be an indication of significant shock heating. However, as was presented by Medling et al. (2019), the estimated synchrotron emission in the northern and southern nuclei is 0.16 and 0.08 mJy, respectively, and hence negligible compared to the observed continuum fluxes reported in Section 3.2. Therefore, we can assume that the continuum emission is dominated by thermal emission from dust and hence compare this estimate to the molecular gas mass derived in the previous section. Following the procedure presented by Scoville et al. (2015), the ISM mass is given by their Equation (3):

$$\begin{eqnarray*}&&{M}_{\mathrm{ISM}}=\displaystyle \frac{0.868\times {S}_{\nu }(\mathrm{mJy}){d}_{\mathrm{Gpc}}^{2}}{{\left(1+z\right)}^{4.8}{T}_{25}{\nu }_{350}^{3.8}{{\rm{\Gamma }}}_{\mathrm{RJ}}}{10}^{10}{M}_{\odot },\end{eqnarray*}$$

where Γ_RJ is a correction factor that accounts for the departure from the RJ regime and is given by

$$\begin{eqnarray*}&&{{\rm{\Gamma }}}_{\mathrm{RJ}}=0.672\times \displaystyle \frac{{\nu }_{350}(1+z)}{{T}_{25}}\times \displaystyle \frac{1}{{e}^{0.672\times {\nu }_{350}(1+z)/{T}_{25}}-1}.\end{eqnarray*}$$

For our observations of NGC 6240, ν_obs corresponds to 235 GHz, or 0.67 in ν₃₅₀ units; z = 0.02448; and d_Gpc = 0.1. Furthermore, we assume T₂₅ = 1. From the ALMA compact and intermediate configuration Band 6 continuum maps of the southern and northern nuclei, as defined by the radius of the SMBH sphere of influence, we measure 235 GHz continuum fluxes of 5.19 ± 0.06 and 1.73 ± 0.03 mJy, respectively, as presented in Section 3.2. These correspond to ISM masses of 2.1 × 10⁹ and 6.9 × 10⁸ M_⊙, respectively. These masses are roughly consistent with those derived from the CO measurements in the previous section. Differences of that order are expected, given the typical uncertainties in the assumed values of α_CO and T₂₅, that cannot be determined directly from our ALMA data at these high resolutions.

To determine if the difference in molecular mass estimations could be caused by a surface brightness sensitivity issue, we use the Scoville et al. (2016) relation to calculate the expected 850 μm luminosity from the CO-derived molecular mass within the central 2'' aperture. We then convert this luminosity to an equivalent 235 GHz continuum emission (again assuming β = 1.8) and estimate the flux density per beam in our long baseline observations. We find that the predicted dust continuum emission per beam associated with the CO-detected molecular concentration in the internuclear 1'' region is an order of magnitude fainter than the sensitivity of our observations. This indicates that dust-based ISM mass estimates can be biased low in situations where galaxies are highly resolved.

The velocity map presented in Figure 3 shows a high-velocity, >500 km s⁻¹, component separated by ∼400 pc in projection to the south of the northern nucleus. A map showing the structure of this high-velocity material is presented in Figure 10. There seems to be a faint bridge of gas extending back to the northern nucleus. This structure was also presented and studied by Feruglio et al. (2013b) and later by Saito et al. (2018) from lower-resolution ¹²CO(2–1) observations. More recently, Cicone et al. (2018) reported the detection of this nuclear outflow in the CO(1–0), CO(2–1), and [C i](1–0) transitions, and Müller-Sánchez et al. (2018) reported it in H₂. Given its high velocity, we speculate that this is an outflow of molecular gas expelled from the nuclear region of the merger system. The total mass of this outflowing material is estimated at 9 × 10⁸ M_⊙, considering the α_CO factor computed specifically for the outflowing material by Cicone et al. (2018). This is very significant, as it represents ∼3.5% of the total molecular gas in the system and 10.5% of the molecular mass in the central region. If we assume that it is linked to the northern nucleus and has been moving at a constant speed, this material would have been expelled at least 0.78 million yr ago.

We estimate the mass outflow rate associated with this high-velocity structure using Equation (3) of Maiolino et al. (2012),

$$\begin{eqnarray*}&&{\dot{M}}_{\mathrm{out}}=3v\displaystyle \frac{{M}_{\mathrm{out}}}{{R}_{\mathrm{out}}},\end{eqnarray*}$$

where v is a representative velocity of the wind, M_out is its mass, and R_out is its size. This equation assumes a relatively simple geometry for the outflow, given by a spherical or, equivalently, conical or multiconical volume uniformly filled by outflowing clouds. We caution that the geometry of this outflow is likely more complex, and hence the values derived using this expression have uncertainties of factors of ∼3 (Lutz et al. 2020). Assuming values of v = 500 km s⁻¹, M_out = 9 × 10⁸ M_⊙ and R_out between 290 and 430 pc, depending on whether the wind arises from the northern nucleus for the latter or the extension of the high-velocity structure for the former. We then obtain values of 3200–4700 M_⊙ yr⁻¹, roughly consistent with the value of 2500 ± 1200 M_⊙ yr⁻¹ reported by Cicone et al. (2018).

We can then constrain the ultimate fate of the wind by comparing the wind velocity with the system's escape velocity. The stellar mass of NGC 6240 is 3.9 × 10¹¹ M_⊙ (Howell et al. 2010); with the abundance matching of Moster et al. (2010), the predicted halo mass is ∼12 × 10¹⁴ M_⊙. The escape velocity for a halo of this mass is ∼1000 km s⁻¹, computed at R₂₀₀ (∼1 Mpc). Assuming there is no further injection of energy, most of the mass will remain bound to the halo. If the material in the wind is efficiently heated and joins the X-ray halo (Nardini et al. 2013), future star formation may be delayed or prevented as a result of the wind's removal of material from the center of NGC 6240. The presence of this high-speed massive clump implies a very efficient process capable of accelerating high-velocity massive amounts of material and, consequently, due to the gas reservoir removal, significantly diminishing or even shutting down nuclear star formation episodes triggered by the major galaxy merger. Taking the previously derived total molecular gas mass of 2.6 × 10¹⁰ M_⊙ and the outflow rate of 3200 M_⊙ yr⁻¹, the implied gas removal time by the wind alone is 8.1 Myr.

In order to explore the kinematic properties of the aforementioned high-velocity material, in Figure 11, we present three p–v diagrams: one along the bulk of the high-velocity emission, roughly in the east–west direction; another tracing the material bridging this cloud to the northern nucleus, roughly oriented north–south; and another connecting the high-velocity emission to the southern nucleus, as shown in Figure 10. No evidence of rotation or obvious kinematic structures can be seen in the high-velocity emission. This emission extends from ∼300 to ∼800 km s⁻¹, spanning ∼03 in size, which corresponds to ∼150 pc in projection at the distance of NGC 6240. The north–south p–v diagram (middle panel of Figure 11) shows a rather steep gradient, which might be indicative of, for example, an expanding shell. A faint connecting bridge can be observed in Figure 10 between the northern nucleus and the high-velocity structure. This emission can also be seen in the middle panel of Figure 11 at a relatively constant velocity of ∼600 km s⁻¹ and spanning spatial offsets between 02 and 06, almost reaching the northern nucleus. No similar connection can be seen relative to the southern nucleus (bottom panel of Figure 11). As reported by Müller-Sánchez et al. (2018), the high-velocity outflow is clearly detected in the H₂ VLT/SINFONI map and spatially coincident [O iii] emission, as observed in HST Wide Field Camera 3 (WFC3) maps of the region. Considering as well the lack of Hα emission at this location, also reported by Müller-Sánchez et al. (2018), the physical origin of this structure remains uncertain. Given its complexity, a full kinematic analysis of this high-velocity structure and the models for its origin are beyond the scope of this paper.

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As shown by Medling et al. (2015), both SMBHs are currently more massive than expected based on the black hole mass–galaxy property correlations, requiring an increase in the total stellar mass of 1.7 × 10¹¹ M_⊙ to reach the scaling relation. Assuming that all of the molecular gas currently in the sphere of influence of each SMBH will be accreted by them, the SMBHs would only grow by ∼45% and hence not change too much in overall mass in the short term. The total mass reservoir is, however, significantly larger, considering the ∼9 × 10⁹ M_⊙ of molecular gas in the central 1'' (500 pc) diameter inferred from the ALMA ¹²CO(2–1) observations. Its precise fate, whether to be accreted, form stars, or be expelled from the system through outflows, is highly uncertain at this point. However, recent observations of a similar system, Mrk 463 (Treister et al. 2018), show that only a very small fraction, <0.01%, of the available mass is actually being accreted by the SMBHs. Therefore, we could expect that, after subtracting the outflowing material, most of the remaining molecular gas will be used as fuel for star formation. Assuming typical efficiencies for star formation for giant molecular clouds of ∼1%–10% (Murray 2011; Ochsendorf et al. 2017), this implies an expected increase of the stellar mass in the nuclear region of (0.3–3) × 10⁸ M_⊙, which is certainly not enough to make the resulting system consistent with the observed black hole–stellar mass correlation. Even if all of the molecular gas in the system is considered, these conclusions are not altered significantly, even if much higher star formation efficiencies are invoked.

The initial calibration of the short baseline (≲1 km) had some slight changes in EBs Xb4da9a/X69a and Xb5fdce/X79d. In the former, the spectral slope of J1550+0527, derived based on the Titan amplitude solutions, was considered in order to obtain an improved bandpass solution. In the latter, J1751+0939 was used as a flux calibrator instead of Pallas, which showed problematic amplitude visibilities inducing a general overestimate of the sources' fluxes. Also, despite the short times on source (1–4 minutes), a single self-calibration step considering phase solutions alone was adopted (no major improvement was found by considering the amplitude solutions). This was done by setting the solution interval parameter to "inf" and not adopting any scan average. Since each scan was ∼50–60 s long, in practice, this was the time range used.

Figure 12 shows the amplitude versus uvdistance visibility distribution for the pointlike calibrators used for flux, bandpass, or phase calibration, as well as the check sources. One can see that J1751+0939 shows modest large-scale structure being picked up by the shortest baselines (≲50 kλ), but this difference is about 1% above the core emission and thus is not expected to affect the calibration. Both J1550+0527 and J1751+0939 show visible variability in either flux, slope, or polarization. Comparing to the AMAPOLA²⁷ database, which makes use of ALMA's Grid Survey observations, one can see that J1550+0527 shows flux variability at the 10%–18% level from 2016 January to June and an increase of polarized flux in the same period. As for J1751+0939, there is an even larger variation in the flux and polarization fractions during the period of observations, with a peak in the first half of 2016 July (Xb5fdce/X79d). The fact that J1751+0939 was used as a flux calibrator in the execution Xb5fdce/X79d, which shows the highest flux density, does not affect the calibration. In fact, both J1651+0129 and J1659+0213, whose amplitude and bandpass calibrations depend on the previous sources, do not show the same flux and polarization fraction variations or any correlation. This indicates that the variation in J1751+0939 on that day was due to intrinsic source variability.

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Likewise, we also display in Figure 13 the phase versus uvdistance for the same sources. All of the sources in these configurations show scatter around 0° with no obvious trend with increasing baseline length, demonstrating that these sources are pointlike and have been observed at the phase center. Here J1550+0527 and J1751+0939 always show a phase scatter below 1°, while the phase calibrator J1651+0129 always appears below 10°, implying that the signal decorrelation is smaller than 1.5% ($\epsilon \equiv {e}^{-{\sigma }_{\phi }^{2}/2}$, where σ_ϕ is the phase rms in radians).

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A.1.1. High Resolution

A.1.2. Self-calibrating Xc5148b/X127e

A.1.3. Imaging Parameters