11institutetext:  Centro de Astrobiología (CAB), CSIC–INTA, Cra. de Ajalvir
Km. 4, 28850 – Torrejón de Ardoz, Madrid, Spain 22institutetext: Instituto de
Física Fundamental, CSIC, Calle Serrano 123, 28006 Madrid, Spain
33institutetext: Università di Firenze, Dipartimento di Fisica e Astronomia,
via G. Sansone 1, 50019 Sesto F.no, Firenze, Italy 44institutetext: INAF -
Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, I-50125 Firenze, Italy
55institutetext: University of Trento, Via Sommarive 14, I-38123 Trento, Italy
66institutetext: European Space Agency, c/o STScI, 3700 San Martin Drive,
Baltimore, MD 21218, USA 77institutetext: Kavli Institute for Cosmology,
University of Cambridge, Madingley Road, Cambridge, CB3 0HA, UK
88institutetext: Cavendish Laboratory - Astrophysics Group, University of
Cambridge, 19 JJ Thomson Avenue, Cambridge, CB3 0HE, UK 99institutetext:
Department of Physics and Astronomy, University College London, Gower Street,
London WC1E 6BT, UK 1010institutetext: Department of Physics, University of
Oxford, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH, UK
1111institutetext: Sorbonne Université, CNRS, UMR 7095, Institut
d’Astrophysique de Paris, 98 bis bd Arago, 75014 Paris, France
1212institutetext: European Space Agency, European Space Research and
Technology Centre, Keplerlaan 1, 2201 AZ Noordwijk, the Netherlands

# No evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFS

Michele Perna e-mail<EMAIL_ADDRESS>evidence of AGN features in
the nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFS Santiago Arribas No evidence of AGN
features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN
features in the nuclei of Arp 220 from JWST/NIRSpec IFS Isabella Lamperti No
evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo
evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFS Miguel
Pereira-Santaella No evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFS Lorenzo Ulivi No evidence of AGN features in the nuclei of
Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp
220 from JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220
from JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFS Torsten Böker No evidence of AGN features in the nuclei of
Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp
220 from JWST/NIRSpec IFS Roberto Maiolino No evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFS Andrew J. Bunker No evidence of AGN
features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN
features in the nuclei of Arp 220 from JWST/NIRSpec IFS Stéphane Charlot No
evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo
evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFS
Giovanni Cresci No evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFS Bruno Rodríguez Del Pino No evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN features in the
nuclei of Arp 220 from JWST/NIRSpec IFS Francesco D’Eugenio No evidence of AGN
features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN
features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN
features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo evidence of AGN
features in the nuclei of Arp 220 from JWST/NIRSpec IFS Hannah Übler No
evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo
evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo
evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFSNo
evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFS Katja
Fahrion No evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec
IFSNo evidence of AGN features in the nuclei of Arp 220 from JWST/NIRSpec IFS
Matteo Ceci No evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFSNo evidence of AGN features in the nuclei of Arp 220 from
JWST/NIRSpec IFS

###### Abstract

Context. Arp 220 is the nearest ultra-luminous infrared galaxy; it shows
evidence of 100 pc-scale molecular outflows likely connected with galaxy-scale
outflows traced by ionised and neutral gas. The two highly obscured nuclei of
Arp 220 are the site of intense star formation, with extreme (far-infrared
based) star-formation rate surface densities, $\Sigma_{SFR}\gtrsim 10^{3}$ M⊙
yr-1 kpc-2. Despite extensive investigations searching for active galactic
nucleus (AGN) activity in the Arp 220 nuclei, direct evidence remains elusive.

Aims. We present JWST/NIRSpec IFS observations covering the $0.9-5.1\ \mu$m
wavelength range of the innermost ($5\arcsec\times 4\arcsec$, i.e. $1.8\times
1.5$ kpc) regions of Arp 220. The primary goal is to investigate the potential
presence of AGN signatures in the nuclear regions by analysing the spectra
extracted from circular apertures of radius 55 pc (0.15″) around each of the
two nuclei.

Methods. The analysis aims to identify highly ionised gas emission lines (with
ionisation potential $>54$ eV) and other spectral features indicative of AGN
activity. Ionised and molecular gas kinematics are also taken into account to
study the outflow signatures at $<60$ pc scales.

Results. We identify $\sim 70$ ionised and $\sim 50$ molecular emission lines
in the nuclear spectra of Arp 220. We use recombination line ratios to measure
optical extinctions in the range A${}_{V}\sim 11-14$ mag. High ionisation
lines are not detected, except the [Mg iv] line at 4.49 $\mu$m which we
interpret as due to shocks rather than to AGN ionisation. We identify
broadening and multiple kinematic components in the H i and H2 lines caused by
outflows and shocks, with velocities up to $\sim 550$ km s-1. Significantly
higher velocities (up to $\sim 900$ km s-1) are detected in the off-nuclear
regions; however, they do not conclusively represent evidence for AGN
activity.

Conclusions. Even with the unprecedented sensitivity of JWST/NIRSpec IFS,
achieving an unambiguous identification or exclusion of the presence of an AGN
in the Arp 220 system remains challenging, because of its extreme dust
obscuration.

###### Key Words.:

galaxies: individual: Arp 220 – galaxies: active – galaxies: starburst –
galaxies: ISM

##  1 Introduction

Arp 220 is the closest ultra-luminous infrared galaxy (ULIRG), with a distance
of $\sim 78$ Mpc and log $L_{IR}/L_{\odot}$ = 12.2 (e.g. Pereira-Santaella et
al. 2021). This late-stage merger shows several tidal structures in the
outskirts (e.g. Arp 1966; Arribas et al. 2001; Perna et al. 2020) and contains
two compact ($<150$ pc), highly obscured nuclei separated by $\sim 370$ pc
(1″) that dominate the infrared luminosity (e.g. Scoville et al. 2007).

The nuclear region of Arp 220 is the site of an intense star formation, with a
star formation rate SFR $\sim 200-250$ M⊙ yr-1 (e.g. Nardini et al. 2010;
Varenius et al. 2016) and incredibly high star formation rate surface
densities ($\Sigma_{SFR}\sim 10^{3}-10^{4}$ M⊙ yr-1 kpc2, from radio and far-
infrared observations; e.g. Barcos-Muñoz et al. 2015; Pereira-Santaella et al.
2021). Long-term monitoring with very long baseline interferometry has
revealed dozens of supernovae and supernova remnants resulting from the
extreme star formation occurring in the compact nuclei (e.g. Varenius et al.
2019). The off-nuclear regions appear to have post-starburst properties, with
weak star formation produced in a few young and low-mass clusters (e.g. Perna
et al. 2020; Chandar et al. 2023), as expected for a late-stage merger.

Although Arp 220 has been extensively studied from radio to hard X-ray and
$\gamma$-ray wavelengths, there is still no convincing direct evidence for the
presence of active galactic nucleus (AGN) activity in either of the two
nuclei; however, indirect arguments provide suggestive indications (e.g. Teng
et al. 2015; Paggi et al. 2017; Yoast-Hull & Murray 2019). Accreting super-
massive black holes (SMBHs) could be buried in Compton-thick absorbers (i.e.
$N_{H}>1.5\times 10^{24}$ cm-2) at the position of the two nuclei (Teng et al.
2015; Sakamoto et al. 2021a) and, for instance, explain the excess of
$\gamma$-ray flux in comparison to all other diagnostics of star-forming
activity (Yoast-Hull et al. 2017).

Evidence of multi-phase (i.e. molecular, neutral, and ionised) outflows in Arp
220 have been inferred from the analysis of multi-wavelength observations.
Collimated outflows have been discovered in both nuclei on scales of $\sim
100$ pc using, for instance, ALMA observations of carbon monoxide transitions
(e.g. Wheeler et al. 2020; Ueda et al. 2022; Lamperti et al. 2022); large-
scale outflows have also been identified with optical integral field
instruments (IFS; e.g. Arribas et al. 2001; Colina et al. 2004; Perna et al.
2020) and associated with extended soft X-ray emission (e.g. McDowell et al.
2003). The possible connection between these multi-phase and multi-scale
outflows has not yet been investigated in detail. Consequently, the mechanisms
at the origin of these outflows, and hence whether they are due to starburst
or AGN winds, are still debated.

In this work, we revisit Arp 220 and its nuclear regions taking advantage of
JWST observations performed with the IFS unit of the NIRSpec instrument
(Jakobsen et al. 2022; Böker et al. 2022); these data cover the near-infrared
(NIR) wavelength range between 0.95 and 5.1 $\mu$m. The observations were
taken as part of the JWST/NIRSpec IFS Guaranteed Time Observations (GTO)
survey “Resolved structure and kinematics of the nuclear regions of nearby
galaxies” (program lead: Torsten Böker). The analysis presented in this paper
aims to identify AGN signatures in the Arp 220 nuclei, as the presence of
highly ionised gas and extreme kinematics unambiguously associated with
accreting SMBHs.

The rest of this paper is organised as follows. Section 2 presents the NIRSpec
IFS observations and data reduction. Section 3 describes the spectral fitting
analysis and the identification of emission line species. Section 4 reports a
qualitative description of the nuclear spectra as well as that of a nearby
bright stellar cluster, and some general properties of the interstellar medium
(ISM). The multi-phase outflow properties, the absence of highly ionised
species, and the comparison between H i and far-IR SFR tracers are discussed
in Section 5. Finally, Section 6 summarises our conclusions. Throughout this
work, we assume $\Omega_{m}=0.286$ and $H_{0}=69.9$ km/s/Mpc (Bennett et al.,
2014). With this cosmology, $1^{\prime\prime}$ corresponds to a distance of
0.368 kpc at $z=0.018$.

##  2 Observations and data reduction

Arp 220 was observed in NIRSpec IFS mode as part of the programme 1267 (PI: D.
Dicken). The present observations were combined in a single proposal with
independent MIRI acquisitions of the same target, with the aim of saving the
telescope slew overhead. The observations were collected on 6 March 2023,
using the NIRSpec IFS mode (Böker et al. 2022).

The dataset consists of two distinct sets of four-point medium cycling
dithering pattern with 15 groups per exposure. The first set is centred at the
mid-point between the coordinates of the western (W) and the eastern (E)
nuclei; the second set is offset $\sim 1\arcsec$ towards the north-west to
cover the shell of H$\alpha$ and [N ii] emission (e.g. Lockhart et al. 2015;
Perna et al. 2020). The total integration time is 933 seconds per set, in each
of the three high resolution (R $\sim 2700$) grating settings, to cover the
spectral range from 0.95 to 5.1 $\mu$m.

###  2.1 Data reduction

We used the STScI pipeline v1.8.2 with CRDS context 1063. A patch was included
to correct some important bugs that affect this specific version of the
pipeline (see details in Perna et al., 2023). In particular, we corrected the
count-rate images for $1/f$ noise by fitting a polynomial. During stage 2, we
masked pixels at the edge of the slices (one pixel wide) to conservatively
exclude pixels for which the correction for the throughput of the spectrograph
optics (contained in the SFLAT reference files) is unreliable. We also
implemented the outlier rejection of D’Eugenio et al. (2023). The combination
of the eight dithers (with drizzle weighting) allowed us to sub-sample the
detector pixels, resulting in cube spaxels with a size of 0.05″(corresponding
to $\sim 20$ pc per spaxel).

The resulting data cubes are known to show sinusoidal modulations in single
spaxel spectra, caused by the undersampling of the point-spread function
(PSF). To correct these so-called “wiggles”, we applied the methodology
presented in Perna et al. (2023), and adapted to the Arp 220 dataset as
reported in Ulivi et al. in prep.. However, we note that these wiggles do not
affect the results presented in this work, because we use spatially integrated
spectra, which are much less affected by PSF undersampling (e.g. Law et al.
2023).

In addition to the Arp 220 dataset, in this manuscript we present the
JWST/NIRSpec observations of another extra-galactic source, the nearby LIRG
VV114 (programme 1328, PI: L. Armus; e.g. Rich et al. 2023). They were reduced
following the same procedure reported above.

Therefore, for both Arp 220 and VV114 we obtained three data cubes,
corresponding to the grating settings G140H/F100LP (referred to as band1
hereinafter), G235H/F170LP (band2), and G395H/F290LP (band3), covering the
wavelength range 0.95 – 5.1 $\mu$m.

Figure 1:  Three-colour near-infrared image of Arp 220 obtained by combining
NIRSpec narrow-band images tracing the continuum emission at $\sim 1\ \mu$m
(blue), $3\mu$m (green) and $5\ \mu$m (red). The positions of the W and E
nuclei are marked with x symbols and match those of the nuclear pc-scale disks
observed in the mm with ALMA (Scoville et al. 2017). The box-circle symbol
marks the position of a bright cluster. The cyan diamond indicates a region
with fast multi-phase outflows.

###  2.2 Astrometric registration

Before analysing the Arp 220 data-cubes, we tested whether they are aligned
with each other. We extracted NIRSpec equivalent narrow-band images from the
wavelength range $1.85-1.87\ \mu$m covered by both band1 and band2 cubes;
these images are perfectly aligned in the pixel space, although a small
discrepancy in the WCS coordinates registration of the reference pixel
introduces an offset of $\delta$RA $=-0.0047\arcsec$ and $\delta$DEC
$=-0.0141\arcsec$. Similarly, we obtained additional narrow-band images
covering the range $3.00-3.02\ \mu$m from the band2 and band3 cubes; also in
this case, the images are perfectly aligned in the pixel space, but present a
discrepancy in WCS registrations, with $\delta$RA $=+0.0017\arcsec$ and
$\delta$DEC $=0.0071\arcsec$. Therefore, we corrected the WCS registrations of
the band2 and band3 cubes to match the coordinates in band1.

The absolute astrometric registration is performed using the ALMA 0.1″
resolution maps of continuum emission at 2.6 mm presented by Scoville et al.
(2017, project 2015.1.00113.S), and the 0.03″ resolution map of 1.3 mm
continuum (project 2017.1.00042.S). We matched the central position of the two
nuclei in the NIRSpec narrow-band image, obtained by collapsing the band3 cube
in the range $5.1-5.12\mu$m, and in the ALMA images. The narrow wavelength
range used to identify the nuclei in the NIR has been preferred to those at
shorter wavelengths, as the latter are significantly more affected by dust
extinction (see Sect. 3). The astrometry corrections we applied to the three
NIRSpec cubes are $\delta$RA $=-0.3028\arcsec$ and $\delta$DEC
$=-0.0209\arcsec$. This offset is consistent with the expected accuracy of
JWST pointing without a dedicated target acquisition procedure (Böker et al.
2023).

##  3 Analysis

The two Arp 220 nuclei separated by $\sim 1$″ are visible in the infrared
(e.g. Scoville et al. 1998), millimetre (e.g. Scoville et al. 2017; Lamperti
et al. 2022), radio (e.g. Baan & Haschick 1995), and X-ray (e.g. Lockhart et
al. 2015). Figure 1 shows the composite near-infrared image of Arp 220
obtained from the combination of three NIRSpec narrow-band images extracted
from the reduced data-cube at $\sim 1\mu$m (blue), $3\mu$m (green) and $5\mu$m
(red), covering a Field of View (FoV) of $5\arcsec\times 4\arcsec$. The
position of the W and E nuclei are marked for visual purposes.

The Arp 220 nuclear morphology is strongly influenced by dust extinction, even
in the NIR (see also Fig. 7). The continuum emission appears as a crescent arc
above the W nucleus at $\sim 1\mu$m and $\sim 3\mu$m. The E nucleus is barely
visible at these wavelengths, but a few bright, clumpy regions likely
associated with clusters can be identified (see also Scoville et al. 1998).
Among them, we marked with a box-circle the position of the brightest clump at
the north of the E nucleus; hereinafter, we refer to it as the NE cluster,
although it may be part of a more extended structure connecting the blue
clumps in the regions between the nuclei in Fig. 1. The emission at $\sim 5\
\mu$m is mostly dominated by the two unresolved nuclei, marked in the figure
with two crosses.

Figure 2:  Arp 220 NE cluster, W and E nuclear spectra from the three NIRSpec
gratings. The spectra are shown in luminosity as a function of the rest frame
wavelength. Vertical blue lines mark the positions of hydrogen transitions
(dark to light blue colours indicate distinct hydrogen series); orange lines
identify H2 lines; red lines are associated with [Fe ii] transitions; green
lines mark He i emission features. Strong broad-band absorption and emission
features are also labelled in the figure.

###  3.1 Spectral extraction

We extracted the spectra of the two nuclei and the NE cluster shown in Fig. 1
for each of the three NIRSpec gratings. We used apertures of three spaxels in
radius, corresponding to $0.15\arcsec$ or $55$ pc. This aperture size broadly
corresponds to two NIRSpec spatial resolution elements at 5 $\mu$m (FWHM $\sim
0.15\arcsec$; D’Eugenio et al. 2023), thus ensuring that the entire point-
source flux is captured at all wavelengths.

The three spectra are reported in Fig. 2; they are shown in rest-frame, taking
into account their relative redshifts: $z_{W}=0.01774$, $z_{E}=0.01839$, and
$z_{NE}=0.01870$ (see Table 1). These values are obtained by measuring the
centroid of the brightest hydrogen lines, and are comparable with those
inferred from sub-mm/mm ALMA spectroscopy at angular resolutions similar to
those of NIRSpec (see Sect. 4.1).

The spectra show a few gaps in wavelength coverage in the middle of the three
bands, due to the gap between the two NIRSpec detectors. The overlapping
wavelength ranges at $\sim 1.85\ \mu$m and $\sim 3\ \mu$m show a perfect match
between the different bands, demonstrating a consistent flux calibration. We
note that no aperture corrections are taken into account in this work.

Table 1: Properties of Arp220 nuclear regions

| W nucleus | E nucleus | NE cluster
---|---|---|---
RA (J2000) | 15:34:57.224 | 15:34:57.294 | 15:34:57.309
DEC (J2000) | +23:30:11.515 | +23:30:11.353 | +23:30:11.698
redshift (H i) | 0.01774 $\pm$ 0.00001 | 0.018396 $\pm$ 0.00001 | 0.018704 $\pm$ 0.00001
redshift (stars) | $0.0179\pm 0.0003$ | $0.0183\pm 0.0002$ | $0.0186\pm 0.0001$
SFR(radio/far-IR) [M⊙ yr-1 ] | 110 – 150 | 60 – 80 | –
SFR(Pa$\alpha$ corr) [M⊙ yr-1 ] | $0.52_{-0.06}^{+0.16}$ | $0.17\pm 0.02$ | $0.94\pm 0.02$
SFR(PAH $3.3\ \mu$m) [M⊙ yr-1 ] | $3.4\pm 0.2$ | $1.5\pm 0.2$ | $1.9\pm 0.2$
Av [mag] | $10.6_{-0.5}^{+1.7}$ | $14.1_{-0.3}^{+0.6}$ | $11.2\pm 0.1$
log ($n_{e}$/[cm-3] | $3.7\pm 0.1$ | $4.5\pm 0.4$ | $3.6\pm 0.1$

  * Notes: SFR measurements are corrected for dust extinction using the AV reported in the table; the latter are obtained from Eq. 1, assuming $R_{V}=3.1$. Radio and far-IR based SFR measurements come from Varenius et al. (2016) and Pereira-Santaella et al. (2021), respectively.

###  3.2 Identification of the lines

The identification of the multitude of emission lines detected in NIRSpec
spectra of Arp 220, and marked in Fig. 2 with vertical lines, was performed
using compilations of NIR features in star-forming and active galaxies from
the literature (e.g. Lutz et al. 2000; Koo et al. 2016; Lamperti et al. 2017;
Lee et al. 2017; Riffel et al. 2019). We also use the NIST atomic spectra
database111https://www.nist.gov/pml/atomic-spectra-database (Ralchenko 2005)
and the HITRAN (high-resolution transmission molecular absorption
database222https://hitran.org/, Rothman 2021) to provide probable
identifications of faint lines that, to our knowledge, have not previously
been reported in the literature.

Figures 8-13 report all species identified in the spectra of the Arp 220
nuclei and the NE cluster, in comparison with those in the nuclear regions of
another nearby interacting LIRG observed with NIRSpec, VV114 (e.g. Rich et al.
2023). Firmly detected line transitions are indicated with solid vertical
lines, and are also reported in Table 2. Probable identifications are marked
in the figures in the Appendix with dotted vertical lines. All emission lines
are well visible in the spectra of the south-west nucleus of VV114 (named SWc2
by González-Alfonso et al. 2023), at higher signal-to-noise and with narrower
profiles than in Arp 220. All three nuclear regions of VV114, namely NE, SW,
and SWc2, have been suggested to harbor deeply obscured AGN (e.g. Rich et al.
2023; González-Alfonso et al. 2023). Hence, they serve as an optimal benchmark
for comparison with the nuclei of Arp 220, which are also presumed to contain
obscured AGN.

Table 2: Arp 220 emission line list in the W and E nuclei, and NE cluster
(band1 spectra)

line | $\lambda_{vac}$ | $f_{W}$ | $f_{E}$ | $f_{NE}$
---|---|---|---|---
| ($\mu$m) | ($10^{-18}$ erg s-1 cm-2) | ($10^{-18}$ erg s-1 cm-2) | ($10^{-18}$ erg s-1 cm-2)
[S iii] | 0.953321 | 598.0${}_{-0.4}^{+0.4}$ | 54.9${}_{-1.2}^{+1.1}$ | 438.0${}_{-1.0}^{+1.8}$
Pa$\epsilon$ | 0.954865 | 87.9${}_{-1.0}^{+1.5}$ | 7.8${}_{-0.8}^{+1.0}$ | 41.7${}_{-0.8}^{+1.9}$
[C i] | 0.982412 | 11.3${}_{-1.1}^{+1.2}$ | 3.3${}_{-0.6}^{+0.8}$ | 8.8${}_{-0.9}^{+1.8}$
[C i] | 0.985295 | 63.7${}_{-0.5}^{+1.2}$ | 15.5${}_{-0.9}^{+1.4}$ | 28.8${}_{-0.3}^{+0.0}$
Pa$\delta$ | 1.005219 | 112.0${}_{-1.5}^{+1.4}$ | 15.8${}_{-1.0}^{+1.2}$ | 67.6${}_{-0.8}^{+1.6}$
[S ii] | 1.028955 | 48.3${}_{-2.7}^{+5.2}$ | 5.7${}_{-1.5}^{+4.2}$ | 11.6${}_{-1.6}^{+2.5}$
[S ii] | 1.032332 | 62.6${}_{-2.2}^{+2.5}$ | 0.9${}_{-0.0}^{+0.0}$ | 13.6${}_{-1.6}^{+2.7}$
[S ii] | 1.033924 | 63.1${}_{-1.0}^{+2.6}$ | 5.3${}_{-1.9}^{+2.4}$ | 17.2${}_{-1.4}^{+2.6}$
[S ii] | 1.037334 | 39.3${}_{-0.3}^{+0.1}$ | 0.9${}_{-0.0}^{+4.4}$ | 11.1${}_{-2.7}^{+6.1}$
$[$N i] | 1.040059 | 42.8${}_{-0.0}^{+0.1}$ | 4.2${}_{-3.2}^{+2.9}$ | 3.8${}_{-2.8}^{+2.5}$
He i | 1.083331 | 628.0${}_{-1.1}^{+2.2}$ | 41.2${}_{-0.4}^{+1.3}$ | 185.0${}_{-0.7}^{+1.0}$
Pa$\gamma$ | 1.094116 | 167.0${}_{-1.6}^{+1.7}$ | 25.5${}_{-0.6}^{+1.1}$ | 205.0${}_{-1.9}^{+1.7}$
[P ii] | 1.147134 | 101.0${}_{-0.9}^{+1.4}$ | 9.4${}_{-0.5}^{+1.0}$ | 45.7${}_{-0.7}^{+1.2}$
[P ii] | 1.188605 | 328.0${}_{-0.7}^{+1.3}$ | 27.8${}_{-0.3}^{+0.7}$ | 126.0${}_{-0.8}^{+0.9}$
$[$Ni ii] | 1.191000 | 31.5${}_{-0.6}^{+1.3}$ | 5.0${}_{-0.7}^{+0.6}$ | 19.6${}_{-0.8}^{+1.1}$
[Fe ii] | 1.257021 | 711.0${}_{-4.3}^{+10.6}$ | 78.0${}_{-1.8}^{+3.6}$ | 371.0${}_{-3.8}^{+7.2}$
[Fe ii] | 1.270690 | 44.5${}_{-2.2}^{+5.1}$ | 0.8${}_{-0.0}^{+1.7}$ | 10.8${}_{-3.6}^{+6.6}$
[Fe ii] | 1.278789 | 75.6${}_{-1.6}^{+3.7}$ | 3.3${}_{-2.5}^{+6.0}$ | 26.7${}_{-3.1}^{+5.5}$
Pa$\beta$ | 1.282167 | 1170.0${}_{-7.0}^{+6.7}$ | 124.0${}_{-2.6}^{+4.6}$ | 822.0${}_{-4.1}^{+6.5}$
[Fe ii] | 1.294453 | 71.7${}_{-2.5}^{+6.6}$ | 7.1${}_{-2.0}^{+2.8}$ | 24.6${}_{-2.0}^{+5.3}$
[Fe ii] | 1.320911 | 289.0${}_{-4.5}^{+7.0}$ | 26.9${}_{-2.0}^{+3.2}$ | 130.0${}_{-3.8}^{+4.6}$
$[$O I] | 1.316871 | 94.2${}_{-2.5}^{+5.3}$ | 4.0${}_{-2.4}^{+3.0}$ | 22.6${}_{-2.2}^{+5.3}$
[Fe ii] | 1.372200 | 185.0${}_{-4.3}^{+6.1}$ | 25.0${}_{-1.2}^{+2.9}$ | 102.0${}_{-2.0}^{+4.3}$
Br 13 | 1.611378 | 51.6${}_{-3.6}^{+6.9}$ | 9.2${}_{-3.0}^{+3.5}$ | 54.0${}_{-3.8}^{+5.2}$
Br 12 | 1.641174 | 71.3${}_{-4.2}^{+6.6}$ | 11.0${}_{-4.5}^{+2.9}$ | 7.6${}_{-3.8}^{+4.0}$
[Fe ii] | 1.644002 | 1340.0${}_{-5.7}^{+11.1}$ | 209.0${}_{-3.9}^{+8.5}$ | 787.0${}_{-8.6}^{+9.6}$
[Fe ii] | 1.677300 | 160.0${}_{-5.2}^{+6.6}$ | 20.3${}_{-2.7}^{+6.4}$ | 67.3${}_{-5.1}^{+5.8}$
Br 11 | 1.681118 | 166.0${}_{-7.0}^{+7.2}$ | 27.7${}_{-2.5}^{+5.1}$ | 117.0${}_{-3.3}^{+5.5}$
H2 | 1.714881 | 14.9${}_{-11.3}^{+72.8}$ | 2.6${}_{-0.1}^{+40.5}$ | 2.5${}_{-0.0}^{+0.0}$
H2 | 1.714958 | 88.8${}_{-59.5}^{+4.6}$ | 30.3${}_{-20.1}^{+12.3}$ | 29.4${}_{-7.0}^{+5.0}$
H2 | 1.748101 | 176.0${}_{-8.6}^{+5.6}$ | 43.8${}_{-8.3}^{+18.9}$ | 66.1${}_{-5.9}^{+1.2}$
Br 10 | 1.736692 | 202.0${}_{-6.4}^{+6.3}$ | 17.4${}_{-2.5}^{+4.5}$ | 111.0${}_{-7.2}^{+6.9}$
Br$\epsilon$ | 1.817916 | 245.0${}_{-5.7}^{+11.5}$ | 34.0${}_{-4.5}^{+8.3}$ | 168.0${}_{-4.5}^{+12.7}$
H2 | 1.834123 | 23.1${}_{-4.2}^{+11.9}$ | 57.8${}_{-22.0}^{+12.6}$ | 37.6${}_{-7.0}^{+8.2}$
H2 | 1.835913 | 370.0${}_{-6.9}^{+12.4}$ | 109.0${}_{-13.4}^{+39.8}$ | 145.0${}_{-3.7}^{+10.8}$

  * Notes: This list comprises all emission lines detected in band1; see Tables 5 and 6 for the ones in band2 and band3, respectively. The reported fluxes are obtained within apertures of 0.3″ in diameter from single Gaussian fits, without dust- and aperture-corrections.

###  3.3 Modelling of the continuum emission

To fit the stellar continuum, we used the penalised-pixel-fitting routine pPXF
(Cappellari 2017). We used the MARCS synthetic library of stellar spectra to
model the stellar continuum (Gustafsson et al. 2008); these templates have a
constant spectral resolution $\Delta\lambda/\lambda$ = 20 000 ($\sigma\sim 6$
km s-1) and cover the range $0.13-20\ \mu$m with wavelengths in vacuum.
Although not as reliable as an empirical stellar library, MARCS templates were
preferred to E-MILES (Vazdekis et al. 2016), which have lower spectral
resolution than our NIRSpec data, and XSL (Verro et al. 2022), which have
narrower wavelength coverage and are affected by telluric features.

We also adopted a sixth-order multiplicative polynomial to better reproduce
the spectral continuum shape. In addition, we used a set of gas emission line
templates (parametrised with Gaussian profiles) together with the stellar ones
to guide the fit. We conducted separate fits for the band1 ($0.95-1.86\mu$m)
and band2 ($1.63-3.11\ \mu$m) cubes, with the latter restricted to wavelengths
below 2.4 $\mu$m, corresponding to the cutoff of the detector gap. For band1,
the wavelengths falling within the NIRSpec detector gap were appropriately
masked during the fit procedure.

Since no prominent stellar features are observed at wavelengths above $2.4\
\mu$m, and because of the presence of dust emission and ice absorption (e.g.
Donnan et al. 2024), we preferred to model the pseudo-continuum with a
Savitzky–Golay filter after masking all prominent emission and absorption
features.

In Figs. 8-10 we report the reconstructed pPXF best-fit continuum emission,
while Figs. 11-13 show the pseudo-continuum shape in the two Arp220 nuclei and
the NE cluster. We used the (pseudo-)continuum-subtracted spectra to model the
emission lines, and hence derive the ISM kinematic and physical properties in
these three regions.

###  3.4 Modelling of the emission lines

After subtracting the (pseudo-)continuum, we modelled the most relevant gas
emission lines using a combination of Gaussian profiles. In particular, we
modelled a set of H i recombination lines - Paschen (Pa), Brackett (Br), Pfund
(Pf), and Humphreys (Hu) series - and molecular hydrogen features, as well as
singly ionised iron transitions.

The sets of H i, H2 and [Fe ii] transitions are modelled using independent
fits, to account for the fact that these three species can trace different gas
phases and can have therefore different systemic velocities and line widths
(e.g. May & Steiner 2017). All transitions were fitted using the
Levenberg–Markwardt least-squares fitting code CAP-MPFIT (Cappellari 2017).

The transitions of a given species were fitted simultaneously (e.g. Perna et
al. 2015), as it allows to better constrain the properties of the transitions
associated with lower S/N. Namely, we constrained the wavelength separation
between emission lines in accordance with atomic physics (considering their
vacuum wavelengths); moreover, we fixed their widths (in km s-1) to be the
same for all the emission lines. Some emission lines are covered by two cubes
(e.g. Br$\epsilon$ at $\sim 1.82\ \mu$m, in band1 and band2 spectra); for
these lines we obtained an average profile by considering the errors from
individual cubes as weights.

For each species we performed two fits, with one and two Gaussian components
per emission line respectively. The number of Gaussian sets (i.e. distinct
kinematic components) used to model the lines of a specific species was
determined on the basis of the Bayesian information criterion (BIC; Schwarz
1978).

The measured fluxes of all identified emission line transitions are reported
in Tables 2, 5 and 6. To streamline the presentation, all tabulated fluxes are
derived from single Gaussian fits; detailed Gaussian decomposition of bright
line profiles is discussed in Sect. 4.

##  4 Results

In this section, we report the general characterisation of the nuclei and the
NE cluster, in terms of spectral shape, extinction, electron density, and
presence of outflows. These properties are discussed in Sect. 5 to investigate
the possible presence of AGN signatures.

###  4.1 Systemic redshift of the two nuclei

The comparison of the three selected regions in Arp 220 was performed after
reporting the individual spectra in luminosity ($\lambda L_{\lambda}$) as a
function of the rest-frame wavelength in vacuum. We used the centroid of the H
i lines to determine the redshifts for the three spectra. In fact, these lines
are quite symmetric and narrow and, therefore, are likely less affected by
outflows and perturbed motions than other lines (see Perna et al. 2020;
Wheeler et al. 2020). In contrast, the brightest H2 and [Fe ii] lines show
prominent blue wings and are blue-shifted by a few 10s km s-1 with respect to
the H i lines (Table 3).

The redshifts inferred from the atomic hydrogen lines are reported in Table 1.
These values are consistent with those obtained from the stellar continuum
modelling, within the uncertainties (see Table 1). Moreover, our measurements
match the ones obtained from the study of sub-mm/mm molecular line transitions
as reported in Sakamoto et al. (2021b), $z_{W}\sim 0.01768$ and $z_{E}\sim
0.01808$, and Pereira-Santaella et al. (2021), $z_{W}\sim 0.01765$ and
$z_{W}\sim 0.01797$. In particular, for the W nucleus, all measurements are
consistent within $\sim 20$ km s-1. The molecular lines in the E nucleus are
quite asymmetric (see Fig. 1 in Sakamoto et al. 2021b), resulting in slightly
larger discrepancies compared to our NIR redshift. However, all measurements
remain consistent within $\lesssim 100$ km s-1, suggesting minimal relative
motions between the various components.

###  4.2 Qualitative comparison between Arp 220 nuclear regions and the bright
cluster

The spectra of the two nuclei and the NE cluster contain a high number of
emission lines: among them hydrogen lines (from the Pa, Br, Pf, and Hu
series), as well as helium, iron, and molecular hydrogen transitions. Notably,
H i lines appear a few times fainter in the E spectrum compared to the W and
NE spectra, likely due to higher extinction levels (see Sect. 4.3) and lower
SFR (Sect. 5.1). The W nucleus displays the most intense H2 lines,
approximately $2\times$ stronger than those in the E and NE spectra.

At 3.3 $\mu$m, all three spectra show a broad emission feature produced by
polycyclic aromatic hydrocarbons (PAHs); the 3.4 $\mu$m aliphatic features are
also present in all three locations. PAHs and aliphatic transitions are
stronger in the W nucleus, by a factor of $\sim 3$ (2) with respect to the E
nucleus (NE cluster; see Sect. 5.1). This suggests that the SFR is higher in
the W nucleus (as PAHs trace recent star formation activity; Sect. 5.1).

The prominent CO absorption bands at $\sim 1.6\ \mu$m and between 2.3 $\mu$m
and 2.4 $\mu$m, produced in the atmospheres of giant and supergiant stars
(Oliva et al. 1995) are clearly visible in the two nuclei as well as in the NE
cluster.

Various other broad-band absorption features are visible as well, mostly
produced by solid state molecules of water-ice (H2O, at 3 and 4.5 $\mu$m), and
carbon monoxide (CO, at 4.67 $\mu$m) transitions. These absorption features
are stronger in the nuclear regions than in the nearby cluster; the
$3.3-3.8\mu$m “ice band wing” (Gibb et al. 2004) is observed only in the two
nuclei. The strength of these absorption bands is therefore likely correlated
with the abundance of dust, given that the two nuclear regions are the most
obscured in the mm range (Sakamoto et al. 2021a). Conversely, the carbon
dioxide ice transitions 12CO2 at 4.27 $\mu$m and 13CO2 at 4.38 $\mu$m have
rather peculiar shapes in the three spectra: the former is stronger in the E
nucleus, but more asymmetric in the W nucleus (with a prominent blue
shoulder), while the NE cluster shows intermediate strength with respect to
the nuclei. The 13CO2 is deeper and narrower in the W nucleus.

The W and E nuclei of Arp 220 show 12CO ro-vibrational lines up to Jlow = 23
(Fig. 13), indicating the presence of high excitation gas. Such transitions
could provide a further indirect evidence of AGN activity. Detailed modelling
is required to infer the CO gas properties, and hence test the more likely
mechanisms responsible for such high excitation (e.g. Baba et al. 2022;
González-Alfonso et al. 2023; Buiten et al. 2023; García-Bernete et al. 2024;
Pereira-Santaella et al. 2024). This analysis will be presented in a
forthcoming paper (Buiten, van der Werf et al., in prep.).

Differences in the three spectra can also be appreciated from a qualitative
comparison between their continuum shapes. Going from the shortest wavelengths
to $\sim 1.8\ \mu$m, all three sources show a pronounced steepening in their
spectra, likely dominated by the diminishing effect of dust obscuration with
increasing wavelength (see also Engel et al. 2011). In the range $1.8-2.3\
\mu$m the W and NE spectra flatten, while the E spectrum still shows the same
steepening observed at shortest wavelengths; this could suggest an higher
extinction at the position of the E nucleus, due to the dust lane visible in
reddish colours in Fig. 1. Between 2.3 and 5$\mu$m, the nuclear regions show a
mild increase in luminosity while NE is almost constant. At $\sim 5\mu$m, the
spectrum of the NE cluster reveals a modest rise in luminosity, while the two
nuclei exhibit a much more pronounced increase likely due to the presence of
warm dust (Armus et al. 2007). It is possible therefore that dust-absorbed
light is re-emitted at wavelengths $\gtrsim 2\ \mu$m in the two nuclei; the
higher dilution of CO bands at $\sim 2.3-2.4\ \mu$m in the nuclei compared to
the NE cluster (Fig. 10) also supports this possibility, as suggested by Engel
et al. 2011.

In the next sections, we will briefly discuss the general properties of some
of these components. We do not pursue the full spectral energy distribution
modelling further in this work, but instead will use a subsequent paper to
focus on a more detailed investigation.

###  4.3 Extinction

Extinction can have an important effect on the estimated source luminosity and
SFR, and we therefore consider different methods for deriving its value. We
compared pairs of emission lines at wide wavelength separation, for which the
intrinsic ratio of line strengths is known. We calculated the colour excess in
terms of Pa$\beta$ and Br$\gamma$,

$\begin{split}E(B-V)&=\frac{2.5}{k(Pa\beta)-k(Br\gamma)}\
log_{10}\left(\frac{\left(Pa\beta/Br\gamma\right)_{obs}}{\left(Pa\beta/Br\gamma\right)_{int}}\right)\\\
&=5.22\
log_{10}\left(\frac{\left(Pa\beta/Br\gamma\right)_{obs}}{5.88}\right)\end{split}$
(1)

where the intrinsic ratio Pa$\beta$/Br$\gamma$ is set to 5.88 assuming an
electron temperature and a density of 104 K and 103 cm-3 respectively (e.g.
Osterbrock & Ferland 2006). We also used the Pf$\gamma$ and Br$\delta$ to
obtain an independent estimate for the colour excess

$\begin{split}E(B-V)&=\frac{2.5}{k(Br\delta)-k(Pf\gamma)}\
log_{10}\left(\frac{\left(Br\delta/Pf\gamma\right)_{obs}}{\left(Br\delta/Pf\gamma\right)_{int}}\right)\\\
&=8.95\
log_{10}\left(\frac{\left(Br\delta/Pf\gamma\right)_{obs}}{1.71}\right)\end{split}$
(2)

where the intrinsic ratio Br$\delta$/Pf$\gamma$ is set to 1.71 (with the same
assumptions as reported above). The coefficients in Eqs. 1 and 2 were derived
assuming a Cardelli et al. (1989) extinction law.

We derived an additional estimate of the colour excess from the flux ratio of
[Fe ii] at 1.257 and 1.644 $\mu$m (e.g. Riffel et al. 2014), which are the two
strongest iron lines in the NIR regime,

$E(B-V)=8.95\
log_{10}\left(\frac{1.34}{\left([FeII]_{1.257}/[FeII]_{1.644}\right)_{obs}}\right)$
(3)

where 1.34 is the intrinsic ratio (we note, however, that this is still a
matter of some debate, and different values are reported in the literature,
from 1.04 to 1.49, see Eriksen et al. 2009).

We measured a colour excess of the order of a few magnitudes ($\sim 3-5$) in
the three regions of interest, corresponding to A${}_{V}\sim 10-14$ (assuming
R${}_{V}=3.1$). We do not observe significant differences in the measurements
obtained from the three diagnostic ratios above mentioned (Eqs. 1, 2, 3).
Therefore, hereinafter we will refer to the results obtained from Eq. 1. The
most extreme extinction is associated with the E nucleus, with A${}_{V}\sim
14$, consistent with the presence of the strongest steepening in continuum
emission at wavelengths $\lesssim 2\ \mu$m with respect to the W nucleus and
the NE cluster (with A${}_{V}\sim 11$, see Table 1).

In previous studies, an extinction of $\approx 6$ mag in the innermost nuclear
regions of Arp 220 was reported, employing diagnostics such as the H i flux
ratios H$\alpha$/H$\beta$ (Perna et al. 2020), H$\alpha$/Pa$\beta$ (Giménez-
Arteaga et al. 2022), and Pa$\alpha$/Br$\gamma$ (Engel et al. 2011). More
extreme measurements were obtained from mid-infrared diagnostics: for
instance, Haas et al. (2001) reported an extinction of a few tens of
magnitudes. This suggests that NIRSpec observations may only probe the outer
gaseous layers, and that the total extinction towards the core of these dusty
regions can be much higher than $A_{V}\sim 10-14$.

###  4.4 Electron density

We derived the electron density in the line-emitting regions surrounding the
E, W, and NE regions of Arp 220 via the comparison of observed line strength
ratios of [Fe ii] transitions with theoretical values.

We used the ratio [Fe ii] 1.677 $\mu$m/1.644 $\mu$m to infer electron
densities in the range log ne $\approx 3.5-4.5$ cm-3 for the three regions of
interest (from Eq. 6 in Koo et al. 2016). For the two nuclei, for which two
kinematic components are required to fit the total profiles, we consider the
narrow components (as the broad one is undetected in the faint [Fe ii] 1.677
$\mu$m). These measurements, reported in Table 1, are discussed in the context
of previous measurements in Sect. 5.2.

Figure 3:  Continuum-subtracted integrated spectra of the three regions of
interest in the vicinity of the Pa$\alpha$, [Fe ii] 1.644 $\mu$m, and H2 2.12
$\mu$m. The original spectra (in black) are reported in velocity space. The
best-fit emission line profiles are shown in orange; the individual Gaussian
profiles used to reproduce the line shapes are reported in blue and green.
Broad and asymmetric profiles are found in the nuclear spectra. The vertical
lines mark the zero-velocity as inferred from the narrow component of H i
lines.

###  4.5 Nuclear outflows

A distinctive feature of Arp 220’s nuclear environment is the presence of
powerful and complex outflows. Compact and collimated cold molecular outflows,
with velocities of a few 100s km s-1, a bipolar morphology and an extension of
$\sim 120$ pc, are observed in both nuclei (e.g. Wheeler et al. 2020). Broader
line profiles are instead observed in the ionized and atomic gas: for
instance, Perna et al. (2021) reported $W80$ velocities (defined as the
difference between the 90th- and 10th-percentile velocities of the fitted line
profile) of $\sim 800$ km s-1. However, these measurements have been obtained
from observations at lower spatial resolution ($\sim 200$ pc), and likely
cannot be used to properly resolve the outflow kinematics on smaller scales.

Figure 3 shows the integrated spectra in the vicinity of the Pa$\alpha$ (top),
[Fe ii] 1.644 $\mu$m (centre), and H2 2.12$\mu$m (bottom) lines for the three
regions of interest. The NE cluster shows symmetric and narrow profiles;
conversely, broad and asymmetric (preferentially blue-shifted) line shapes are
observed in both nuclear regions, with outflowing gas at velocities up to
$-500$ km s-1. In this work, we consider as outflow velocity $v_{max}=\Delta
V_{broad}+2\sigma_{broad}$, where $\sigma_{broad}$ is corrected for
instrumental resolution (e.g. Fluetsch et al. 2021). In Table 3 we report the
best-fit kinematic parameters, considering both non-parametric (e.g. $W80$),
and parametric outflow velocity measurements (e.g. $v_{max}$; see caption of
Table 3 for details).

A more detailed discussion about the presence of AGN winds possibly
responsible for the broad and asymmetric profiles observed in the Arp 220
nuclei is presented in Sect. 5.3.

Table 3: Kinematic parameters

W nucleus
---
| H i | H2 | [Fe ii] | [Mg iv]
$\Delta V_{1}$ | $0\pm 1$ | $-1_{-4}^{+26}$ | $0_{-1}^{+2}$ | $-111\pm 3$
FWHM1 | $238\pm 2$ | $250_{-6}^{+24}$ | $242_{-4}^{+6}$ | $234_{-10}^{+6}$
$\Delta V_{2}$ | $-72\pm 2$ | – | $-110_{-8}^{+16}$ | –
FWHM2 | $480_{-4}^{+6}$ | – | $525_{-9}^{+24}$ | –
$W80$ | $370\pm 3$ | $270\pm 20$ | $380\pm 3$ | $260\pm 10$
$V10$ | $-250\pm 2$ | $-180\pm 15$ | $-270\pm 6$ | $-265\pm 10$
$V_{max}$ | $-480_{-4}^{+9}$ | $-220_{-3}^{+30}$ | $-560\pm 25$ | $-309\pm 8$
E nucleus
| H i | H2 | [Fe ii] | [Mg iv]
$\Delta V_{1}$ | $0_{-2}^{+10}$ | $-40\pm 5$ | $-8_{-3}^{+8}$ | $-130_{-8}^{+14}$
FWHM1 | $205_{-27}^{+9}$ | $230\pm 3$ | $180_{-70}^{+45}$ | $370_{-14}^{+20}$
$\Delta V_{2}$ | $-70_{-10}^{+60}$ | $-201_{-7}^{+13}$ | $-60\pm 30$ | –
FWHM2 | $425_{-60}^{+50}$ | $405\pm 10$ | $420_{-60}^{+120}$ | –
$W80$ | $370\pm 30$ | $440\pm 50$ | $380\pm 42$ | $400\pm 30$
$V10$ | $-250\pm 20$ | $-370\pm 30$ | $-250\pm 42$ | $-370\pm 20$
$V_{max}$ | $-430_{-45}^{+70}$ | $-540_{-10}^{+20}$ | $-430_{-90}^{+140}$ | $-440\pm 20$
NE cluster
| H i | H2 | [Fe ii] | [Mg iv]
$\Delta V_{1}$ | $0_{-3}^{+4}$ | $-40\pm 5$ | $-14_{-8}^{+7}$ | $-81_{-3}^{+6}$
FWHM1 | $260_{-3}^{+30}$ | $225_{-60}^{+10}$ | $210\pm 30$ | $260\pm 10$
$\Delta V_{2}$ | – | – | $11_{-6}^{+33}$ | –
FWHM2 | – | – | $340_{-40}^{+170}$ | –
$W80$ | $250\pm 10$ | $270_{-50}^{+10}$ | $290\pm 40$ | $305\pm 40$
$V10$ | $-160\pm 10$ | $-180_{-30}^{+10}$ | $-170\pm 30$ | $-240\pm 30$
$V_{max}$ | $-230_{-3}^{+25}$ | $-230_{-60}^{+10}$ | $-190\pm 20$ | $-300\pm 6$
NW outflow
| H i | H2 | [Fe ii] | [Mg iv]
$\Delta V_{1}$ | $-175_{-10}^{+30}$ | $-180_{-7}^{+16}$ | $-290_{-10}^{+15}$ | -
FWHM1 | $130_{-20}^{+80}$ | $80_{-45}^{+20}$ | $640_{-30}^{+50}$ | -
$\Delta V_{2}$ | $-340_{-30}^{+40}$ | $-260_{-30}^{+70}$ | – | –
FWHM2 | $650_{-50}^{+105}$ | $650_{-70}^{+60}$ | – | –
$\Delta V_{3}$ | – | $85_{-6}^{+11}$ | – | –
FWHM3 | – | $210_{-30}^{+50}$ | – | –
$W80$ | $680_{-30}^{+90}$ | $715\pm 10$ | $715\pm 40$ | –
$v10$ | $-690\pm 30$ | $-590\pm 15$ | $-670\pm 35$ | –
$v_{max}$ | $-890_{-45}^{+130}$ | $-820\pm 35$ | $-810_{-45}^{+120}$ | –

  * Notes: all kinematic parameters are given in km s-1. Velocity shifts $\Delta V_{i}$ and line widths FWHMi refer to the individual Gaussian components ($i=$ 1 to 3, at maximum) required to model the emission line profiles. $V10$ is defined as 10th-percentile velocity of the fitted line profile, while $W80$ is defined as the difference between the 90th- and 10th-percentiles. $v_{max}$ is defined as $\Delta V_{broad}+2\sigma_{broad}$.

##  5 Discussion

We used the aperture-extracted spectra of the obscured nuclei E and W and the
bright cluster NE to identify all gas emission line features in the NIRSpec
spectral range, as well as the main absorption features due to dust and stars.
We identified $\sim 70$ ionised gas and $\sim 50$ molecular emission lines;
Tables 2, 5, and 6 give the (vacuum) wavelengths and fluxes for all the firmly
identified features in Arp220, while Table 3 displays the kinematic parameters
of the main optical lines discussed in this paper. We note that a few well
detected emission lines are not tabulated, as (i) their modelling would
require detailed deblending from other lines, or (ii) their observed
wavelengths are very close to the edges of the NIRSpec detector gaps.

In this section we discuss the main spectral features observed in the spectra
reported in Fig. 2 in relation with the possible presence of intense episodes
of starburst activity and AGN accretion in the two nuclei.

###  5.1 Star-formation tracers

We derived the instantaneous ($\sim 10$ Myr) SFR from the Pa$\alpha$ dust
corrected luminosity, following Giménez-Arteaga et al. (2022, and assuming a
initial mass function):

$\begin{split}SFR\ [M_{\odot}\ yr^{-1}]&=4.4\times
10^{-42}\times\left(\frac{H\alpha}{Pa\alpha}\right)_{int}\times
L(Pa\alpha)_{corr}\ [ergs^{-1}]\\\ &=3.7\times 10^{-41}\times
L(Pa\alpha)_{corr}\ [ergs^{-1}]\end{split}$ (4)

where the dust-corrected Pa$\alpha$ luminosity can be obtained with
A(Pa$\alpha$) $=k($Pa$\alpha$ $)\times$ E(B-V). In deriving Eq. 4 we assumed
electron density of 103 cm-3 (the dependence from ne is very weak) and
temperature of 104 K. Moreover, for the E and W nuclear emission we consider
only the narrow components (dashed blue Gaussian lines in Fig. 3), thus
excluding the contribution of outflowing material.

The use of E(B-V) inferred from Pa$\beta$/Br$\gamma$ imply dust corrections of
$\sim 5-8\times$ factors for the Pa$\alpha$ flux, resulting in a SFR $\sim
0.5$ M⊙ yr-1 for the W nucleus, $\sim 0.2$ M⊙ yr-1 for the E nucleus, and
$\sim 0.9$ M⊙ yr-1 for the NE cluster. Therefore, these three regions account
for a total SFR(Pa$\alpha$) $\sim 2$ M⊙ yr-1. Considering the extraction
regions of $\sim 0.01$ kpc2, these values would be translated in star-
formation surface densities $\Sigma_{SFR}\sim 20-100$ M⊙ yr-1 kpc-2.

We also use the calibration of SFR for the 3.3 $\mu$m PAH from Kim et al.
(2012, see their Eq. 5) to convert the measured luminosities (L(PAH) $\sim
2\times 10^{40}$ erg s-1 for the W, L(PAH) $7\times 10^{39}$ erg s-1 for the
E, and L(PAH) $10^{40}$ erg s-1 for NE) in L(IR), and then the Kennicutt &
Evans (2012)333Slightly smaller (but consistent) values would be obtained
using the Piqueras López et al. 2016 relation. relation to derive dust-
corrected SFR(PAH) in the range $\sim 1.5-3.4$ M⊙ yr-1 (see Table 1).

In previous optical and NIR studies of Arp 220, slightly higher SFR
measurements were obtained when integrating over the whole galaxy: Perna et
al. (2020) reported SFR(H$\alpha$) $\lesssim 10$ M⊙ yr-1 (using the Balmer
decrement to correct for extinction), while Giménez-Arteaga et al. (2022)
obtained SFR(Pa$\beta$) $\sim 19$ M⊙ yr-1 from HST narrow-band images (using
Pa$\beta$/H$\alpha$ for the dust-correction). Giménez-Arteaga et al. (2022)
also reported the total SFR inferred from the 24 $\mu$m luminosity, 76 M⊙
yr-1; this value is significantly less affected by dust (A${}_{V}\sim 90$
would imply a correction of a factor 2 at 24 $\mu$m). This value is more
consistent with the total SFR(PAH) $\sim 40$ M⊙ yr-1 we measure from NIRSpec
data using an integrated spectrum extracted over a very large aperture
($r=1.5\arcsec$, i.e. 0.56 kpc).

Therefore, our NIRSpec-based outcomes are definitely in line with previous
estimates obtained from optical, NIR, and mid-infrared bands. However, it is
worth noting that the SFR inferred from the radio and sub-mm/mm regime appears
to be several times higher, reaching 200-250 M⊙ yr-1 (e.g. Varenius et al.
2016; Pereira-Santaella et al. 2021). In particular, these authors reported
SFR $\sim 60-80$ M⊙ yr-1 and $\sim 110-150$ M⊙ yr-1 for the E and W nuclei
separately, hence significantly higher than our NIR-based nuclear
measurements. This discrepancy may stem from a combination of factors: (i) The
far-IR and radio emission might not be solely attributable to star formation,
but could also include contributions from AGN, which are totally extinguished
in the NIR. (ii) The SFR derived from FIR measurements traces the star
formation history over a longer period (up to 100 Myr) compared to the
Pa$\alpha$ and PAH features ($\lesssim 10$ Myr).

Summarising, the NIRSpec-based SFR measurements do not offer a solution to the
existing discrepancy among various star-formation tracers identified in
previous studies.

###  5.2 High electron densities

The [Fe ii]-based electron densities at the position of the two nuclei and the
NE cluster (log (ne/cm-3) $\sim 3.7-4.5$) consistently exceed those inferred
from optical [S ii] 6718, 32 $\AA$ lines, which are of the order of 200 cm-3
(Perna et al. 2020). This suggests that [Fe ii] lines, having critical
densities higher than [S ii] (n${}_{\rm crit}\sim 1600$ cm-3 for [S ii] 6716
$\AA$; n${}_{\rm crit}\sim 10^{4}$ cm-3 for iron lines), trace post-shock
regions with higher compression. An alternative explanation might be that the
higher spatial resolution of NIRSpec data ($\sim 0.1\arcsec$ at the
wavelengths of the ne-sensitive [Fe ii] lines) compared to the MUSE
observations ($\sim 0.6\arcsec$) enables a better isolation of regions with
higher degree of ionisation.

At higher redshift (up to $z\sim 9$), [S ii] and [O ii] 3727, 30 $\AA$ line
ratios are systematically used to measure the ISM electron densities; values
in the range $\sim 300-1000$ cm-3 are generally reported for both AGN (e.g.
Perna et al. 2017; Cresci et al. 2023) and star-forming galaxies (e.g. Förster
Schreiber et al. 2019; Davies et al. 2021; Rodríguez Del Pino et al. 2023;
Isobe et al. 2023). Significant correlations between $\Sigma_{SFR}$ and ne
have been reported in the literature (e.g. Shimakawa et al. 2015), although
they cover relatively narrow ranges in both $\Sigma_{SFR}$ ($\sim$ 0.1 – 1 M⊙
yr-1 kpc-2) and ne ($\sim$ 10 – 103 cm-3). From the extrapolation of these
relations (e.g. Eq. 3 in Reddy et al. 2023) to the very high $\Sigma_{SFR}$
measured in Arp 220 (see Sect. 5.1), we would expect log(ne/cm-3) $\sim 3$.
However, this value stands roughly 30 times lower than the one measured at the
position of the Arp 220 W nucleus using [Fe ii] lines. Instead, our
measurements more closely resemble electron densities observed in supernova
remnants (e.g., Lee et al. 2017) and at the bases of protostar jets (e.g.,
Davis et al. 2011), with [Fe ii]-based densities spanning log (ne/cm-3)
$\approx 3.5-4.5$.

Therefore, on the basis of available information, we cannot discriminate
between the two above mentioned scenarios: the extremely high electron
densities we measured in the nuclei and the NE cluster could be due to higher
levels of compression and/or ionisation with respect to [S ii] gas emitting in
the optical. We can definitely exclude any potential contamination from high-
density AGN Broad Line Region (BLR) in the [Fe ii]-based electron densities
obtained for the Arp 220 nuclei, as the forbidden [Fe ii] line transitions
cannot originate from BLR high-density regions ($N_{\rm e}({\rm BLR})\gg
N_{c}$([Fe ii])).

###  5.3 Presence of multi-phase outflows

The NIRSpec nuclear spectra of Arp220 show ionised and molecular gas
transitions with broad profiles and prominent blue wings associated with
outflows. In particular, high-velocity components are detected in Pa$\alpha$
and [Fe ii] 1.644 $\mu$m at the position of the E and W nuclei. Prominent H2
blue wings are solely observed in the E nucleus. This suggests that the high-
velocity gas in the W nucleus is sufficiently hot (because of shocks, or
intense ultraviolet radiation by hot stars or AGN) to disassociate the H2
molecules.

The nuclear material can reach velocities of up to $v_{max}\sim 550$ km s-1
(Fig. 3). We anticipate here that such outflows reach distances of $\sim 1$
kpc and velocities of $-900$ km s-1 in the off-nuclear regions covered by
NIRSpec (Ulivi et al., in prep.). Figure 4 shows the Pa$\alpha$ and H2 2.12
$\mu$m line profiles extracted from a circular aperture of $r=0.15\arcsec$ at
a distance of $\sim 1.3\arcsec$ to the west from the W nucleus, as defined in
Fig. 1 (cyan diamond); at this position, both ionised and molecular gas have
outflow velocities up to $900$ km s-1 (Table 3). Similar velocities have also
been observed in the ionised and atomic gas phases at kpc-scales in the
optical regime (e.g. Perna et al. 2020).

Nuclear outflows in NIR H i, H2 and iron lines have been observed in nearby
ULIRGs and active galaxies. While sources lacking an active SMBH typically
exhibit less extreme outflow velocities (on the order of a few hundred km
s-1), AGN can manifest outflows across a broad velocity range, from a few
hundred to over a thousand km s-1 (e.g. Emonts et al. 2017; Perna et al. 2021;
Speranza et al. 2022; Villar Martín et al. 2023). Consequently, distinguishing
between starburst- and AGN-driven outflows in Arp 220, based solely on
velocity estimates, proves to be a challenging task. A more comprehensive
analysis of the multi-phase and multi-scale outflows will be necessary to
elucidate whether the outflow energetics in Arp 220 are more consistent with
AGN or starburst launching mechanisms; this work will be presented in a
forthcoming paper, Ulivi et al., in prep..

Figure 4:  Observed profiles of Pa$\alpha$ and H2 2.12 $\mu$m, showing fast
outflows. The spectra are extracted from the region in Fig. 1 marked with a
cyan diamond, and reported in velocity space. The best-fit emission line
profiles are shown in orange; the individual Gaussian profiles used to
reproduce the line shapes are reported in blue, purple, and green. The
vertical grey lines mark the zero-velocity of the E nucleus (corresponding to
the kinematic centre of Arp 220); the dashed red vertical lines in the top
panel identify the He i lines close to Pa$\alpha$ (see Table 2).

###  5.4 Absence of high-ionisation lines

High excitation emission lines with ionisation potentials IP $>54$ eV (the He+
ionisation energy) can be formed either in gas photoionised by hard
ultraviolet radiation or in very hot, collisionally ionised plasma. Therefore,
they can be used as sign for AGN activity (e.g. Oliva et al. 1994; Moorwood et
al. 1997).

Most of these highly ionised lines have previously been identified in ground-
based (e.g. Rodríguez-Ardila et al. 2011; Lamperti et al. 2017; Cerqueira-
Campos et al. 2021; den Brok et al. 2022) and space-based (e.g. Lutz et al.
2000; Sturm et al. 2002) NIR spectra of some nearby Seyfert galaxies (Table
4). No quantitative studies of the coronal lines prevalence in the AGN
population have been made yet; however, there are indications for a more
likely presence of such lines in unobscured type-1 sources rather than in
obscured sources. This is consistent with the fact that some coronal lines
have been recently detected with NIRSpec in the nearby Seyfert 1.5 nucleus of
NGC7469 (Bianchin et al. 2023) but not in the highly obscured active nuclei of
VV114 (see Appendix B; but see also Speranza et al. 2022; Álvarez-Márquez et
al. 2023). The lack of high ionisation lines in AGN sources may reflect an AGN
ionising continuum lacking photons below a few keV: this could explain the
absence of such line emission in obscured sources (e.g. Rodríguez-Ardila et
al. 2011).

High ionisation lines are not detected in the nuclear spectra of Arp 220.
However, we note the presence of a bright feature near $\sim 4.488\mu$m, hence
close to expected wavelength of the [Mg iv] line (IP $\sim 80$ eV). Previous
studies reported the detection of the [Mg iv] line in a few nearby active
galaxies alongside other highly ionised species; for instance, Sturm et al.
(2002) reported [Mg iv]/[Ar vi] $\sim 0.5$ in Circinus and NGC1068; Bianchin
et al. (2023) obtained [Mg iv]/[Ar vi] $\sim 1.5$. Notably, [Ar vi] is not
detected in Arp 220. Given its similar IP to [Mg iv] (see Table 4), the
observed feature at the expected wavelength of [Mg iv] may involve alternative
transitions. For instance, the low-ionisation line SI ($\lambda 4.4855$ in
vacuum, IP $\sim 10$ eV) could be responsible of the line features in the Arp
220 spectra.

Figure 5 shows the emission line at the wavelength position of the [Mg iv] in
the three spectra of interest: these features are broad (FWHM $\sim 300$ km
s-1) and blue-shifted ($\Delta V\sim-100$ km s-1) with respect to the [Mg iv]
expected wavelength (grey solid line at $v=0$ km/s). Moreover, their peaks are
not at the systemic of the SI transition neither. Pereira-Santaella et al.
(subm.) show that this feature can be associated with shocked [Mg iv] emission
due to supernova explosions. In support of this explanation we note that this
line feature is bright at the position of intense clusters in Arp 220, but it
is faint in the surroundings of the two nuclei.

In Table 4, we report the main high ionisation lines in the range covered by
our NIRSpec spectra; for all but [Mg iv] we report the $3\sigma$ upper limits
inferred from NIRSpec data. These upper limits, and in particular the [Si vi]
one, can be used to compare the NIRSpec measurements with those from the
literature. In fact, [Si vi] is detected in the largest number of AGN spectra
(Lamperti et al. 2017; den Brok et al. 2022).

Figure 6, top panel, shows the distribution of L([Si vi]) as a function of the
redshift for a sample of nearby active galaxies selected in the hard X-ray
band (14-195 keV) from the Swift/Burst Alert Telescope (BAT) survey, as part
of the BASS sample (Koss et al. 2017; Ricci et al. 2017). We display both [Si
vi] detections and 3$\sigma$ upper limits, alongside the Arp 220 non
detections. It is evident that NIRSpec enables us to obtain more stringent
upper limits compared to previous observations of other systems at $z\sim
0.018$ (by $\sim 1$ dex).

Figure 6, bottom panel, shows the correlation between intrinsic X-ray and [Si
vi] luminosity (e.g. Rodríguez-Ardila et al. 2011; Lamperti et al. 2017) for
the same sources. The upper limits for [Si vi] in Arp 220 nuclei are indicated
at the intrinsic X-ray (2–10 keV) luminosities of $<1\times 10^{42}$ erg s-1
(W) and $<4\times 10^{41}$ erg s-1 (E), computed by Paggi et al. (2017) from
the $3\sigma$ upper limits on the neutral Fe–K$\alpha$ line. The Arp 220 [Si
vi] upper limits are approximately 1 dex below the L([Si vi]) – L${}_{\rm
2-10\ keV}$ relation (dashed magenta line). However, it should be noted that
measuring intrinsic X-ray luminosity for highly obscured sources can be highly
uncertain, and that the Arp 220 X-ray luminosities are merely upper limits.
Consequently, it is possible that the nuclei of Arp 220 may not exhibit
notable distinctions from other AGN with undetected [Si vi] reported in the
figure, for which the presence of accreting SMBHs is confirmed in the hard
X-ray band.

To conclude, there is no obvious highly ionised line emission associated with
AGN from either nucleus of Arp 220, despite the many bright H i, He i, H2, and
low excitation [Fe ii] lines detected at very high signal-to-noise. Moreover,
the very stringent upper limits derived from NIRSpec observations cannot be
used to discard the presence of AGN in Arp 220 nuclei when we take into
account the extremely high nuclear extinction (Sect. 4.3).

Table 4: High ionisation lines

line | $\lambda_{vac}$ | IP | ref | log $L_{line}$
---|---|---|---|---
| ($\mu$m) | (eV) |  | (erg s-1)
(1) | (2) | (3) | (4) | (5)
$[$S viii] | 0.9916 | 281 | L17, C21 |  $<$37.0
Cr viii | 1.0109 | 160 |  |  $<$36.9
$[$Fe xiii] | 1.0750 | 331 | L17, C21 |  $<$37.0
$[$Fe xiii] | 1.0801 | 331 | L17, C21 |  $<$37.0
$[$S ix] | 1.2523 | 329 | B23, C21 |  $<$36.9
$[$Si x] | 1.4305 | 351 | C21 | –
Ti vi | 1.7156 | 99 |  |  $<$37.3
$[$S xi] | 1.9201 | 448 | L17, C21 |  $<$37.1
$[$Si xi] | 1.9327 | 401 | L17, C21 |  $<$37.1
[Si vi] | 1.9646 | 167 | L17, C21, B23 |  $<$37.1
$[$Al ix] | 2.0450 | 285 | C21 |  $<$37.1
Sc v | 2.3118 | 73 |  |  $<$37.0
$[$Ca viii] | 2.3211 | 127 | C21,R06 |  $<$37.0
$[$Si vii] | 2.4810 | 205 | B23 |  $<$37.0
$[$Si ix$]$ | 2.5850 | 303 | S02, L00 |  $<$37.1
Al v | 2.9050 | 120 |  |  $<$37.0
$[$Mg viii$]$ | 3.0280 | 225 | S02,B23, L00 |  $<$37.0
$[$Ca iv] | 3.2070 | 51 | L00 |  $<$36.7
$[$Al vi] | 3.6603 | 154 |  |  $<$36.9
$[$Al viii$]$ | 3.6900 | 242 |  |  $<$36.9
$[$Si ix] | 3.9290 | 303 | S02, L00 |  $<$37.0
$[$Ca v] | 4.1585 | 67 |  |  $<$37.0
[Mg iv] | 4.4880 | 80 | S02, B23, L00 |  $38.03\pm 0.01$ (W)
|  |  |  |  $37.72\pm 0.02$ (E)
[Ar vi] | 4.5292 | 75 | B23, L00 |  $<$37.0

  * Notes: references of papers reporting the detection of high ionisation lines: Sturm et al. (2002, S02), Lamperti et al. (2017, L17), Bianchin et al. (2023, B23), Cerqueira-Campos et al. (2021, C21), Lutz et al. (2000, L00). [Mg iv] luminosities are reported for both E and W nuclei; for all other high ionisation lines we report 3$\sigma$ upper limits for the luminosities, assuming FWHM $=250$ km s-1. [Si x] is not covered by NIRSpec observations (it falls in detector gap).

Figure 5:  Continuum-subtracted integrated spectra of the three regions of
interest in the vicinity of the [Mg iv] 4.488 $\mu$m. The original spectra (in
black) are reported in velocity space. The best-fit emission line profiles are
shown in orange. Broad and asymmetric profiles are found in the nuclear
spectra. The vertical grey lines mark the zero-velocity as inferred from the
narrow component of H i lines; the dashed red lines marks the expected
position of the S I line. Figure 6: [Si vi] luminosity for Arp 220 and a
sample of nearby AGN. (Top panel:) [Si vi] luminosity as a function of
redshift for the sample of BASS AGN with detected (purple circles) and
undetected (green squares) [Si vi] line, from Lamperti et al. (2017). Arp 220
measurements are represented by red stars. For sources with undetected [Si vi]
lines, we provide three-sigma upper limits (illustrated with arrows). (Bottom
panel): [Si vi] luminosity as a function of intrinsic X-ray (2–10 keV)
luminosity for the same sources reported in the top panel. The dashed line
indicate the linear relation between L([Si vi]) and L${}_{2-10\ keV\ int}$
obtained by Lamperti et al. (2017, not considering upper limits). X-ray
luminosity upper limits for the nuclei of Arp 220 are taken from Paggi et al.
(2017).

###  5.5 Near-infrared emission line diagnostic diagrams

We used the diagnostic diagrams presented by Riffel et al. (2013), Colina et
al. (2015), and Calabrò et al. (2023), to investigate the ionisation
mechanisms responsible of the emission line ratios observed in the nuclei and
the NE cluster. Prior to analysis, all fluxes required for the diagnostics
were corrected for dust extinction, using the $A_{V}$ values reported in Table
1.

According to the diagram presented by Riffel et al. (2013), based on the H2
2.12$\mu$m/Br$\gamma$ vs. [Fe ii] 1.26$\mu$m/Pa$\beta$ line ratios, both
nuclei and the NE cluster are in the AGN region; the NE cluster has slightly
lower line ratios and it is therefore closer to the ‘star-forming’ region of
the diagram. In the diagram proposed by Colina et al. (2015), based on the H2
2.12$\mu$m/Br$\gamma$ vs. [Fe ii] 1.64$\mu$m/Br$\gamma$, the two nuclei and NE
cluster are in the region dominated by AGN and supernovae. Therefore, a
similar ionisation mechanism should be present in the W, E, and NE regions of
Arp 220 according to both Riffel et al. (2013) and Colina et al. (2015).

We also explored the diagnostic diagram recently proposed by Calabrò et al.
(2023), that use combinations of ratios of the [S iii], [Fe ii], [P ii] and [C
i] lines over hydrogen recombination lines. In all the diagrams, the two
nuclei are slightly above the maximum starburst line (by 0.1 – 0.5 dex). The
NE clump lies on (or just below) the maximum starburst line. We note that the
two nuclei are well above the maximum starburst line in the [C i]/Pa$\beta$
diagram, compatible with both AGN and shock ionisations (see Fig. 13 in
Calabrò et al., 2023).

In summary, these diagnostics do not allow us to distinguish between
supernovae and AGN ionisation mechanisms. More detailed investigations of the
flux ratios in the circum-nuclear region of Arp 220 will be presented in a
future work.

##  6 Conclusions

This paper provides an overview of the JWST/NIRSpec IFS observations of Arp
220. It introduces the data reduction, and the continuum and emission line
data modelling. It shows the near-infrared spectra of the E and W nuclei, in
comparison with that of a bright stellar cluster at $\sim 0.4\arcsec$ (0.15
kpc) north-east from the E nucleus. These spectra are extracted using circular
apertures of radius 55 pc (0.15″) for each region. In the following, we
summarise our main findings.

We identify broadening and multiple kinematic components in H i, H2 and [Fe
ii] lines caused by outflows (Fig. 3). The emission line features show maximum
velocities of up to $\sim 550$ km s-1 in the nuclear spectra. Even higher
velocities ($\sim 900$ km s-1) are detected in the off-nuclear regions (Fig.
4). However, they do not conclusively represent a direct evidence of AGN
activity. A companion paper will report the spatially resolved gas kinematics,
and the outflow energetics will be used to test the theoretical predictions
for both starburst- and AGN-driven winds (Ulivi et al., in prep.).

High ionisation lines associated with AGN activity are not detected in the
nuclear spectra of Arp 220. We observe the presence of bright [Mg iv] 4.488
$\mu$m lines (80 eV) in the two nuclear spectra and the NE stellar cluster
(Fig. 5). However, no other high ionisation lines are detected; in particular
the [Ar vi] 4.529 $\mu$m line, which has an IP and wavelength close to those
of [Mg iv], is not detected. The [Mg iv] lines are broader (FWHM $\sim 300$ km
s-1) and blue-shifted ($\Delta V\sim-100$ km s-1) compared to those of the H i
lines; moreover, [Mg iv] is bright at the position of clusters, but it is
faint in the surroundings of the two nuclei. These arguments suggest that this
emission line is not produced by an AGN, but rather due to stellar activity.
This result is in line with recent findings of Pereira-Santaella et al.
(subm.), who recently reported the presence of [Mg iv] in other nearby ULIRGs
not hosting accreting SMBHs, with kinematic properties similar to those in Arp
220; they also provided grids of photoionization and shock models to determine
the conditions where [Mg iv] originates, concluding that shocks by supernova
explosions can explain the observed [Mg iv] lines.

The luminosity upper limits inferred for the high ionisation (IP $=167$ eV)
[Si vi] emission line at the position of the Arp 220 nuclei are $\sim 1$ dex
below the L([Si vi]) – L${}_{2-10\ keV,int}$ observed for X-ray AGN
(Rodríguez-Ardila et al. 2011; Lamperti et al. 2017), but still broadly
consistent with other sources hosting accreting SMBHs but missing [Si vi]
emission (Fig. 6).

We also report the presence of high excitation 12CO gas transitions (up to
J${}_{low}=23$) in the Arp 220 nuclei (Fig. 13), possibly due to AGN activity
(e.g. González-Alfonso et al. 2023). A companion paper (Buiten et al., in
prep.) will provide detailed modelling of CO gas conditions to assess which
mechanism among star formation and AGN contributes to the observed high
excitation CO lines.

To conclude, an unambiguous identification of the presence of AGN in the Arp
220 system remains elusive. At the same time, we cannot exclude the presence
of accreting SMBHs, because of the extremely high extinction measured at the
position of the two nuclei, with recombination lines-based A${}_{V}\sim
11-14$. NIRSpec observations may only probe the outer gaseous layers, and the
total extinction towards the core of the Arp 220 nuclear regions could be much
higher (see e.g. Haas et al. 2001; Sakamoto et al. 2021a). This could explain
the absence of highly ionised gas transitions, as also reported in previous
works on obscured AGN (e.g. Rodríguez-Ardila et al. 2011; Lamperti et al.
2017).

## Acknowledgements

We are grateful to Paul van der Werf, Victorine Buiten, and Angèle Taillard
for discussing various aspects of this work

MP, SA, and BRP acknowledge support from the research project
PID2021-127718NB-I00 of the Spanish Ministry of Science and Innovation/State
Agency of Research (MCIN/AEI/10.13039/501100011033). IL acknowledges support
from PID2022-140483NB-C22 funded by AEI 10.13039/501100011033 and BDC 20221289
funded by MCIN by the Recovery, Transformation and Resilience Plan from the
Spanish State, and by NextGenerationEU from the European Union through the
Recovery and Resilience Facility. MPS acknowledges support from grant
RYC2021-033094-I funded by MICIU/AEI/10.13039/501100011033 and the European
Union NextGenerationEU/PRTR. RM acknowledges support by the Science and
Technology Facilities Council (STFC), by the ERC Advanced Grant 695671
“QUENCH”, and by the UKRI Frontier Research grant RISEandFALL; RM is further
supported by a research professorship from the Royal Society. GC, MP, and LU
acknowledge the support of the INAF Large Grant 2022 ”The metal circle: a new
sharp view of the baryon cycle up to Cosmic Dawn with the latest generation
IFU facilities”.

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##  Appendix A NIRSpec narrow-band images

Figure 7:  NIRSpec narrow-band images tracing the continuum emission at $\sim
1\mu$m (top), $3\mu$m (centre) and $5\mu$m (bottom). The position of the W and
E nuclei are marked with x symbols; the box-circle symbol marks the position
of a bright cluster (see also Fig. 1)

##  Appendix B Integrated spectra

Figure 8:  $0.95-1.42\ \mu$m portion on the spectra of the Arp 220 and VV114
nuclear regions. The spectra are shown in luminosity as a function of the rest
frame wavelength, considering the redshift of each region. The grey, cyan and
orange curves represent the pPXF best-fit. The vertical blue lines mark the
position of hydrogen transitions; the orange lines identify H2 lines; the red
lines are associated with [Fe ii] transitions; the green lines mark He i
features; solid vertical lines mark the position of faint metal lines;
potential identifications are indicated with dotted vertical lines. The
positions of highly ionised gas transitions are indicated with narrow vertical
bands in purple. In the figure, we also indicate the position of strong
stellar absorption features.  Figure 9:  $1.45-1.86\ \mu$m portion on the
spectra of the Arp 220 and VV114 nuclear regions. See 8 for details.  Figure
10:  $1.85-2.40\ \mu$m portion on the spectra of the Arp 220 and VV114 nuclear
regions. See 8 for further details.  Figure 11:  $2.4-3.1\ \mu$m portion on
the spectra of the Arp 220 and VV114 nuclear regions. The pseudo-continuum
described in Sect. 3.3 are shown with dashed curves. See 8 for further
details. Figure 12:  $3.1-4.0\ \mu$m portion on the spectra of the Arp 220 and
VV114 nuclear regions. See 8 for further details.  Figure 13:  $4.1-5.2\ \mu$m
portion on the spectra of Arp 220 and VV114 nuclear regions. The E and W
nuclei show 12CO ro-vibrational lines up to J${}_{low}=23$ (close to the [Mg
iv] line transition, at $\sim 4.48\ \mu$m, and at $\sim 4.88\ \mu$m; see e.g.
González-Alfonso et al. 2023). Detailed modelling of the CO transitions will
be presented in Buiten et al., in prep.. See 8 for further details.

##  Appendix C Arp 220 emission line list (band2 and band3 cubes)

Table 5: Arp 220 emission line list in the W and E nuclei, and NE cluster
(band2 spectra)

line | $\lambda_{vac}$ | $f_{W}$ | $f_{E}$ | $f_{NE}$
---|---|---|---|---
| ($\mu$m) | ($10^{-18}$ erg s-1 cm-2) | ($10^{-18}$ erg s-1 cm-2) | ($10^{-18}$ erg s-1 cm-2)
He i | 1.869044 | 273.0${}_{-3.3}^{+6.8}$ | 62.3${}_{-2.8}^{+4.7}$ | 193.0${}_{-3.9}^{+5.0}$
He i | 1.870234 | 112.0${}_{-3.5}^{+5.8}$ | 35.1${}_{-2.9}^{+4.6}$ | 36.7${}_{-3.3}^{+6.3}$
Pa$\alpha$ | 1.875627 | 7260.0${}_{-2.9}^{+5.4}$ | 1510.0${}_{-2.6}^{+2.1}$ | 6390.0${}_{-4.0}^{+5.8}$
H2 | 1.892094 | 253.0${}_{-5.0}^{+5.1}$ | 77.7${}_{-11.9}^{+32.3}$ | 85.6${}_{-1.0}^{+0.8}$
Br$\delta$ | 1.945095 | 455.0${}_{-5.7}^{+4.5}$ | 77.2${}_{-2.8}^{+4.5}$ | 334.0${}_{-3.2}^{+6.6}$
He i | 1.954840 | 27.2${}_{-1.8}^{+3.1}$ | 46.6${}_{-3.1}^{+2.7}$ | 30.5${}_{-2.7}^{+2.1}$
H2 | 1.957722 | 659.0${}_{-5.4}^{+6.9}$ | 243.0${}_{-45.6}^{+86.2}$ | 278.0${}_{-3.3}^{+8.8}$
H2 | 2.033927 | 268.0${}_{-3.4}^{+3.4}$ | 149.0${}_{-3.4}^{+5.1}$ | 116.0${}_{-3.1}^{+3.8}$
He i | 2.058600 | 247.0${}_{-1.3}^{+1.4}$ | 62.6${}_{-0.6}^{+1.3}$ | 215.0${}_{-1.1}^{+2.2}$
H2 | 2.065756 | 81.7${}_{-0.7}^{+0.5}$ | 28.9${}_{-2.0}^{+4.7}$ | 33.3${}_{-4.2}^{+4.2}$
H2 | 2.073655 | 126.0${}_{-4.6}^{+3.0}$ | 69.6${}_{-3.4}^{+4.0}$ | 59.1${}_{-4.3}^{+3.1}$
H2 | 2.122011 | 901.0${}_{-5.1}^{+7.9}$ | 427.0${}_{-3.6}^{+6.0}$ | 345.0${}_{-5.4}^{+6.6}$
Br$\gamma$ | 2.166129 | 762.0${}_{-4.1}^{+5.5}$ | 156.0${}_{-4.9}^{+9.2}$ | 686.0${}_{-2.8}^{+5.9}$
H2 | 2.223475 | 263.0${}_{-5.5}^{+3.4}$ | 123.0${}_{-4.1}^{+4.2}$ | 100.0${}_{-2.8}^{+5.5}$
H2 | 2.247903 | 147.0${}_{-3.3}^{+2.4}$ | 76.2${}_{-3.5}^{+3.0}$ | 69.6${}_{-0.9}^{+1.0}$
H2 | 2.406793 | 1010.0${}_{-5.6}^{+9.7}$ | – | –
H2 | 2.413640 | 402.0${}_{-3.9}^{+8.0}$ | – | –
H2 | 2.423932 | 940.0${}_{-4.7}^{+9.2}$ | – | –
H2 | 2.454956 | 633${}_{-3.0}^{+5.4}$ | 375${}_{-4.4}^{+5.0}$ | 223${}_{-3.8}^{+5.3}$
H2 | 2.475765 | 173.0${}_{-1.8}^{+3.1}$ | 124.0${}_{-4.0}^{+3.0}$ | 62.4${}_{-0.8}^{+1.0}$
H2 | 2.500174 | 332.0${}_{-2.1}^{+4.0}$ | 207.0${}_{-3.1}^{+4.4}$ | 106.0${}_{-2.6}^{+4.9}$
H2 | 2.501650 | 50.9${}_{-2.4}^{+3.0}$ | 19.4${}_{-2.7}^{+5.0}$ | 39.0${}_{-3.1}^{+2.9}$
H2 | 2.528247 | 87.8${}_{-1.2}^{+0.6}$ | 58.6${}_{-1.7}^{+3.9}$ | 34.6${}_{-1.7}^{+3.9}$
H2 | 2.551198 | 103.0${}_{-1.6}^{+5.0}$ | 72.7${}_{-1.7}^{+4.1}$ | 33.8${}_{-2.9}^{+4.1}$
H2 | 2.558725 | 96.0${}_{-2.7}^{+4.6}$ | 61.4${}_{-3.7}^{+5.3}$ | 54.8${}_{-1.7}^{+2.2}$
H2 | 2.560064 | 193.0${}_{-2.6}^{+3.3}$ | 121.0${}_{-2.6}^{+5.4}$ | 48.9${}_{-2.1}^{+3.0}$
H2 | 2.570043 | 124.0${}_{-2.7}^{+4.5}$ | 89.6${}_{-2.7}^{+3.3}$ | 44.8${}_{-2.5}^{+2.8}$
H2 | 2.585185 | 31.3${}_{-1.6}^{+2.5}$ | 20.1${}_{-3.6}^{+4.1}$ | 1.76${}_{-0.7}^{+3.6}$
H2 | 2.604197 | 66.5${}_{-2.4}^{+3.7}$ | 46.1${}_{-3.0}^{+4.8}$ | 25.5${}_{-3.1}^{+4.1}$
Pf 14 | 2.612655 | 121.0${}_{-0.1}^{+0.3}$ | 70.9${}_{-2.3}^{+4.7}$ | 84.6${}_{-2.2}^{+6.0}$
Br$\beta$ | 2.625878 | 1840.0${}_{-3.9}^{+6.2}$ | 749.0${}_{-4.4}^{+6.5}$ | 1610.0${}_{-4.8}^{+5.9}$
H2 | 2.635329 | 103.0${}_{-3.0}^{+2.9}$ | 89.8${}_{-3.2}^{+4.0}$ | 45.8${}_{-3.8}^{+3.5}$
H2 | 2.654084 | 80.5${}_{-4.3}^{+3.1}$ | 65.1${}_{-2.6}^{+3.7}$ | 24.2${}_{-1.6}^{+5.8}$
Pf 13 | 2.675139 | 58.9${}_{-6.0}^{+5.1}$ | 17.1${}_{-3.5}^{+4.0}$ | 51.4${}_{-2.5}^{+6.4}$
H2 | 2.710478 | 48.9${}_{-3.0}^{+0.4}$ | 19.7${}_{-1.4}^{+5.4}$ | 12.7${}_{-2.7}^{+3.6}$
H2 | 2.718847 | 1.19${}_{-0.004}^{+0.006}$ | 2.29${}_{-0.74}^{+5.5}$ | 4.55${}_{-2.2}^{+1.6}$
H2 | 2.720371 | 103.0${}_{-0.94}^{+2.4}$ | 81.3${}_{-2.2}^{+4.9}$ | 26.4${}_{-1.9}^{+3.7}$
H2 | 2.726984 | 87.1${}_{-2.6}^{+1.0}$ | 59.6${}_{-1.5}^{+3.4}$ | 33.2${}_{-1.5}^{+3.0}$
H2 | 2.731438 | 17.3${}_{-2.6}^{+5.4}$ | 10.0${}_{-2.1}^{+2.9}$ | 14.8${}_{-1.6}^{+3.6}$
H2 | 2.748299 | 32.0${}_{-3.4}^{+4.0}$ | 20.3${}_{-2.8}^{+3.4}$ | 11.2${}_{-3.9}^{+2.5}$
Pf 12 | 2.758276 | 110.0${}_{-1.9}^{+6.3}$ | 48.0${}_{-2.7}^{+4.6}$ | 109.0${}_{-3.0}^{+4.3}$
H2 | 2.786396 | 12.5${}_{-2.0}^{+4.4}$ | 13.6${}_{-2.8}^{+4.0}$ | 6.99${}_{-2.4}^{+3.7}$
H2 | 2.802750 | 601.0${}_{-3.8}^{+3.6}$ | 313.0${}_{-1.3}^{+4.3}$ | 244.0${}_{-1.7}^{+3.2}$
H2 | 2.804052 | 1.25${}_{-0.004}^{+0.007}$ | 1.41${}_{-0.02}^{+0.03}$ | 1.08${}_{-0.01}^{+0.01}$
Pf 11 | 2.873004 | 58.7${}_{-2.4}^{+5.7}$ | 8.35${}_{-1.8}^{+3.2}$ | 74.4${}_{-2.6}^{+3.5}$
H2 | 2.974311 | 54.0${}_{-2.5}^{+2.5}$ | 22.7${}_{-2.4}^{+2.7}$ | 27.2${}_{-2.2}^{+2.9}$
H2 | 3.004118 | 132.0${}_{-2.8}^{+3.7}$ | 48.5${}_{-2.1}^{+2.7}$ | 53.5${}_{-1.9}^{+3.5}$
Pf$\epsilon$ | 3.039211 | 128.0${}_{-2.9}^{+7.3}$ | 21.3${}_{-4.1}^{+3.5}$ | 110.0${}_{-3.3}^{+6.6}$

  * Notes: This list comprises all emission lines detected in band2; see Tables 2 and 6 for those in band1 and band3, respectively. The reported fluxes are obtained from single Gaussian fits, without dust- and aperture-corrections.

Table 6: Arp 220 emission line list in the W and E nuclei, and NE cluster
(band3 spectra)

line | $\lambda_{vac}$ | $f_{W}$ | $f_{E}$ | $f_{NE}$
---|---|---|---|---
| ($\mu$m) | ($10^{-18}$ erg s-1 cm-2) | ($10^{-18}$ erg s-1 cm-2) | ($10^{-18}$ erg s-1 cm-2)
H2 | 3.164034 | 15.2${}_{-2.4}^{+2.1}$ | 8.13${}_{-1.6}^{+2.0}$ | 13.4${}_{-2.3}^{+3.0}$
H2 | 3.190077 | 28.2${}_{-1.8}^{+2.4}$ | 13.4${}_{-1.4}^{+1.1}$ | 11.9${}_{-2.2}^{+2.6}$
H2 | 3.235257 | 310.0${}_{-2.2}^{+3.3}$ | 117.0${}_{-1.1}^{+2.8}$ | 117.0${}_{-2.2}^{+2.2}$
Pf$\delta$ | 3.297001 | 409.0${}_{-0.6}^{+0.8}$ | 74.3${}_{-3.5}^{+5.5}$ | 293.0${}_{-2.3}^{+1.7}$
H2 | 3.626469 | 152.0${}_{-2.7}^{+3.5}$ | 70.4${}_{-1.6}^{+2.5}$ | 47.6${}_{-1.4}^{+4.1}$
H2 | 3.663464 | 23.2${}_{-2.6}^{+2.8}$ | 29.0${}_{-1.7}^{+1.9}$ | 12.2${}_{-1.0}^{+3.2}$
Hu 18 | 3.692642 | 30.2${}_{-1.6}^{+2.3}$ | 9.63${}_{-0.8}^{+2.1}$ | 23.5${}_{-1.9}^{+2.8}$
H2 | 3.723999 | 5.49${}_{-4.7}^{+4.8}$ | 0.898${}_{-0.006}^{+0.021}$ | 3.75${}_{-2.2}^{+4.8}$
H2 | 3.724736 | 110.0${}_{-4.0}^{+4.7}$ | 59.3${}_{-1.2}^{+2.6}$ | 30.4${}_{-4.3}^{+3.8}$
Pf$\gamma$ | 3.740568 | 414.0${}_{-3.3}^{+3.8}$ | 111.0${}_{-3.5}^{+4.7}$ | 398.0${}_{-1.7}^{+2.6}$
Hu 17 | 3.749402 | 12.6${}_{-2.1}^{+3.2}$ | 2.87${}_{-2.0}^{+2.9}$ | 19.5${}_{-2.4}^{+3.1}$
H2 | 3.807736 | 215.0${}_{-2.7}^{+2.1}$ | 93.5${}_{-16.9}^{+33.1}$ | 70.6${}_{-1.6}^{+2.2}$
Hu 16 | 3.819460 | 16.9${}_{-2.0}^{+3.3}$ | 3.96${}_{-2.1}^{+3.2}$ | 25.4${}_{-2.3}^{+2.7}$
C I (?) | 3.841473 | 24.2${}_{-0.9}^{+1.7}$ | 26${}_{-1.6}^{+2.4}$ | 13.3${}_{-1.3}^{+1.9}$
H2 | 3.846434 | 349.0${}_{-2.6}^{+3.1}$ | 186.0${}_{-2.4}^{+2.9}$ | 99.6${}_{-1.2}^{+3.4}$
Hu 15 | 3.907557 | 41.7${}_{-2.5}^{+4.3}$ | 23.6${}_{-1.7}^{+3.2}$ | 36.7${}_{-1.8}^{+2.7}$
H2 | 4.162772 | 63.1${}_{-2.5}^{+5.4}$ | 47.0${}_{-2.2}^{+3.5}$ | 24.1${}_{-1.5}^{+1.7}$
Hu 13 | 4.170803 | 26.8${}_{-3.6}^{+3.6}$ | 77.4${}_{-10.3}^{+18.8}$ | 45.0${}_{-1.9}^{+3.4}$
H2 | 4.181426 | 941.0${}_{-2.6}^{+5.4}$ | 580.0${}_{-3.3}^{+4.8}$ | 248.0${}_{-1.7}^{+3.0}$
Hu 12 | 4.376464 | 78.1${}_{-3.0}^{+3.7}$ | 33.0${}_{-1.6}^{+3.4}$ | 59.5${}_{-1.2}^{+3.4}$
H2 | 4.410159 | 454.0${}_{-2.0}^{+2.1}$ | 292.0${}_{-1.3}^{+4.9}$ | 117.0${}_{-1.2}^{+2.2}$
H2 | 4.416979 | 142.0${}_{-1.8}^{+2.8}$ | 71.2${}_{-2.3}^{+3.7}$ | 41.4${}_{-1.0}^{+1.4}$
H2 | 5.053536 | 972.0${}_{-4.8}^{+7.1}$ | 644.0${}_{-5.8}^{+5.7}$ | 289.0${}_{-1.9}^{+4.5}$
Hu$\delta$ | 5.128669 | 148.0${}_{-3.6}^{+6.7}$ | 120.0${}_{-3.5}^{+5.3}$ | 183.0${}_{-2.7}^{+2.1}$

  * Notes: This list comprises all emission lines detected in band3; see also Tables 2 and 5 for those in band1 and band2, respectively. The reported fluxes are obtained from single Gaussian fits, without dust- and aperture-corrections.