11institutetext: Nicolaus Copernicus Astronomical Center, Polish Academy of
Sciences, Rabiańska 8, 87-100 Toruń, 11email<EMAIL_ADDRESS>22institutetext: Institute of Astronomy, Nicolaus Copernicus University,
Grudziadzka 5, 87-100 Toruń, Poland 33institutetext: Max-Planck-Institut für
Radioastronomie, Auf dem Hügel 69, 53-121 Bonn, Germany

# A search for cool molecular gas
in GK Persei and other classical novae

Tomek Kamiński 11 Helena J. Mazurek 22 Karl M. Menten 33 Romuald Tylenda 11

Detecting molecular line emission from classical nova remnants has the
potential of revealing information on the composition of the ejecta, in
particular, it can deliver accurate isotopic ratios in the matter processed by
a thermonuclear runaway. We conducted searches toward more than 100 classical
novae for emission in lines of the CO or HCN molecules using single-dish
telescopes and interferometric arrays at millimeter and submillimeter
wavelengths. The survey demonstrates that classical novae, young or old, are
not strong sources of molecular emission at submillimeter and millimeter
wavelengths. Additionally, we mapped CO emission around Nova Persei 1901 (GK
Per), earlier claimed to be circumstellar in origin. Our measurements indicate
that the observed emission is from the interstellar medium. Although no
molecular emission at millimeter and submillimeter wavelengths has been found
in classical novae, it is still likely that some will be detected with high-
sensitivity interferometers such as ALMA.

###### Key Words.:

surveys – stars: novae, cataclysmic variables – circumstellar matter – stars:
individual: GK Per, V1148 Sgr

##  1 Introduction

Classical novae are close binary systems where a white dwarf accretes material
from a companion, typically a late-type dwarf. Accretion via a disk
accumulates hydrogen in a degenerate layer on the surface of the white dwarf
(Warner 1995; Bode & Evans 2008). At some point, the pressure and temperature
at the bottom of the accreted envelope become high enough for hydrogen burning
to start. This initiates the thermonuclear runaway process that is directly
responsible for classical novae outbursts at luminosities of 104–105 L☉. The
energetic explosion on the surface of a white dwarf ejects the accumulated
material, and, according to some models, may even remove some material forming
the outer layers of the white dwarf. A typical mass lost during a single nova
explosion is 10-7–10-3 M☉. A few percent of this material was processed by the
thermonuclear burning.

Although studied for decades, classical novae are still not fully understood.
In particular, the physics of thermonuclear runaway and dust formation
associated with novae remain unclear. It is very difficult to model the
thermonuclear runaway, mainly because (1) we do not know the mechanism that
mixes the accreted envelope and the outermost shells of the white dwarf; (2)
we do not fully understand the process of material ejection; and (3) nuclear
reaction rates are uncertain (Hix 2001). In consequence, it is not known how
effective classical novae are in enhancing the ejected material with products
of nuclear burning. Knowing the composition of classical novae ejecta is of
great importance because the material eventually reaches the interstellar
medium, enhancing it with the nuclear ashes. Classical novae are currently
thought to be the major source of 13C, 15N, and 17O in the Galaxy (Li et al.
2016; Romano & Matteucci 2003) and are likely to significantly contribute to
the Galactic content of 7Li, 19F, and 26Al (Gehrz et al. 1998; Izzo et al.
2015).

There is thus a great need for robust measurements of the composition of
classical novae ejecta. The so-called presolar “nova grains”, should they
indeed originate in classical novae outburst, provide only a limited and
biased information on the ejecta composition (Amari et al. 2001; Iliadis et
al. 2018). To date, only a handful of observations has been made whose results
permit direct tests of theoretical predictions for the isotopic yields of nova
thermonuclear runaway. There is thus a great need for observational data from
which elemental isotopic abundances could be inferred (e.g. Black 2005) and
compared to models. Observational astronomy often relies on molecular spectra
to constrain the isotopic composition of objects.

Molecules are expected to form in ejecta of classical novae (Pontefract &
Rawlings 2004; Derdzinski et al. 2017). Observations collected in the last
decades corroborate this notion. First, it is clear from observations that
30–40% of classical novae produce dust 10–100 days after their outbursts (Bode
& Evans 2008). Dust cannot be produced without the presence of parent
molecules or nucleation seeds, whose formation requires simple diatomic
species. Furthermore, dust provides protection for the molecules from the
Galactic ultraviolet (UV) field and from the harsh radiation of the nova
explosion. Second, molecules are observed in spectra of classical novae at
optical and near-infrared (NIR) wavelengths, including bands of C2, CN, and CO
(e.g., Wilson & Merrill 1935; Nagashima et al. 2014; Banerjee et al. 2016;
Fujii et al. 2021, and references therein). At these wavelengths, the isotopic
shift between different isotopologs is typically too small to be resolved,
owing to generally modest spectral resolutions and high intrinsic line widths
(500–5000 km s-1). The spectral confusion hampers stringent observational
tests for nova nucleosynthesis. A straightforward method of measuring isotopic
ratios, successfully applied to the envelopes of evolved stars and the
interstellar medium, are ground-based observations of molecular rotational
transitions at millimeter (mm) and submillimeter (submm) wavelengths.

There have been attempts to detect mm-wave molecular emission in classical
novae, which almost all (see Sect. 3) turned out to be unsuccessful or
resulted in detections in ambient material unrelated to the nova phenomenon
(for a review see Sect. 13 in Bode & Evans 2008). Past attempts resulted in
upper limits on CO emission obtained with the Swedish-ESO Submillimeter
Telescope (SEST) (Weight et al. 1993; Nielbock & Schmidtobreick 2003), the
IRAM 30 m telescope (Shore & Braine 1992), and with the University of Texas
Millimeter Wave Observatory (Albinson & Evans 1989). The reports from before
$\sim$1995 are however unreliable and often contradictory. The outflow
velocities in novae are of 500–5000 km s-1 and are thus expected to have very
broad lines. which require receivers with a broad bandwidth, and with flat and
stable spectroscopic baselines, which in the past were difficult to attain.
Most of the published negative results were obtained with instruments which
did not fulfill these requirements. Some claims of nonexistent emission have
been recently debunked. For instance, a lack of detection of CO(1–0) was
reported for Nova 1670 (also known as CK Vul) by Weight et al. Observed in
2014–2015 with modern receiver systems, CK Vul showed clear emission lines of
CO for $J_{\rm up}$ $<$7 (Kamiński et al. 2015, 2017). However, CK Vul turned
out to not be a classical nova but a red nova, that is a stellar merger
product, although with similarly broad lines as expected for classical novae,
that is $\approx$300 km s-1. Another example is $\eta$ Carinae, which is not a
nova but a luminous blue variable with broad lines. $\eta$ Car was long
claimed to be free of molecular mm-wave emission (e.g. Cox & Bronfman 1995).
However, when observed with APEX/FLASH in 2011, it showed a plethora of
molecular features (Loinard et al. 2012). Another example is the first
observation of rotational lines of SiO and CO in the remnant of SN 1987A
(Kamenetzky et al. 2013) using the ALMA interferometer. The lines of the
supernova remnant are very broad, FWHM$\approx$2200 km s-1, but were recovered
within the ALMA correlator bandwidth, even though the lines were broader than
the covered spectral band. For classical novae, spectral lines can be as broad
as a few GHz in the mm range, and one needs a receiver system with a large
instantaneous bandwidth to be even able to recognize a line and to distinguish
its presence from an elevated continuum level.

Encouraged by the discoveries in CK Vul, $\eta$ Car, and SNR 1987A, we decided
to revise the view on the presence of molecular emission in classical novae by
undertaking an observation campaign aimed to detect molecules at mm
wavelengths using modern instruments, including single-dish telescopes and
interferometers. The search for molecular emission in the common isotopologs
was meant as the first step towards providing isotopic ratios in thermonuclear
ashes of classical novae. In Sect. 2.1 and 2.2, we present observations with
APEX and IRAM 30 m telescopes for a large sample of classical novae and
related objects. In Sect. 3, we revise the nature of molecular emission near
nova remnant GK Per based on mapping observations from the IRAM telescope. In
Sect. 4, we present our attempts to detect molecular emission with the SMA and
ALMA interferometers. Finally, in Sect. 5 we summarize the observational
material and conclude our findings.

##  2 Single-dish observations

In 2014–2016, we conducted surveys of old classical nova with the APEX and
IRAM telescopes, targeting emission lines of CO and HCN in novae on both
hemispheres. The idea was to observe many sources at a sensitivity level
comparable to that at which emission lines would be easily detectable in CK
Vul. This typically required 5–20 min of integration on source. We searched
for lowest rotational transitions of CO and HCN, because these are usually the
strongest lines in cool circumstellar111By circumstellar we mean in this paper
material that is directly originating from stellar outflows. envelopes.
Particularly, emission of HCN might be expected to be strong if the gas is
enhanced in CNO-processed material, as expected for classical novae.

The sources were selected from the Central Bureau for Astronomical Telegrams
(CBAT) List of Novae in the Milky
Way222http://www.cbat.eps.harvard.edu/nova_list.html, which includes many
historical novae. Individual targets were chosen for observations based on
their availability during periods assigned to us at the observatories. The
lists of sources observed with the APEX and the IRAM telescopes are given in
Tables 2 and 2, respectively. There, we classified the observed objects using
the SIMBAD astronomical database333http://simbad.u-strasbg.fr/simbad/ of
Centre de Donneés Astronomiques de Strasbourg (CDS). According to this
classification, as of August 2020, eleven of the observed sources are not
novae despite being listed as such in the CBAT catalog. Twelve other objects
are dwarf novae. We decided to include them in our survey summary
nevertheless. Technical details on the observations are given below.

Table 1: Objects observed with APEX.

Object | ID | CDS type | Observation date | RA (J2000) | Dec (J2000)
---|---|---|---|---|---
NOVA Cir 1914 | AICIR | Nova | 13-Jul-2016 | 14:49:31.2 | –68:51:35.9
V* AP Cru | APCRU | Nova | 09-Jun-2015 | 12:31:20.5 | –64:26:25.2
V* AT Sgr | ATSGR | Nova | 14-Jul-2016 | 18:03:30.8 | –26:28:28.5
V* BD Pav | BDPAV | CataclyV* | 10-Jun-2015 | 18:43:11.9 | –57:30:44.9
V* CG CMaaa$a$Observation reported in Kamiński et al. (2015).bb$b$ Observed only in CO(3–2), CO(3–2) OSB setups. | CG-CMA | DwarfNova | 24-Aug-2014 | 07:04:05.0 | –23:45:34.6
NOVA Vel 1905aa$a$Observation reported in Kamiński et al. (2015). | CN-VEL | Nova | 24-Aug-2014, 09-Jun-2015 | 11:02:38.5 | –54:23:09.5
V* CP Pup | CPPUP | Nova | 08-Jun-2015 | 08:11:46.1 | –35:21:04.9
NOVA Vel 1940aa$a$Observation reported in Kamiński et al. (2015). | CQ-VEL | Nova | 26-Aug-2014 | 08:58:50.9 | –53:20:17.8
V* DY Pupaa$a$Observation reported in Kamiński et al. (2015). | DY-PUP | Nova | 25-Aug-2014, 08-Jun-2015 | 08:13:48.5 | –26:33:56.5
NOVA Sgr 1926 | FMSGR | Nova | 13-Jul-2016 | 18:17:18.1 | –23:38:27.0
NOVA Mus 1983aa$a$Observation reported in Kamiński et al. (2015). | GQ-MUS | Nova | 28-Aug-2014 | 11:52:02.3 | –67:12:20.2
V* GR Sgr | GRSGR | Nova | 13-Jul-2016 | 18:22:58.5 | –25:34:47.3
V* GU Musaa$a$Observation reported in Kamiński et al. (2015).bb$b$ Observed only in CO(3–2), CO(3–2) OSB setups. | GU-MUS | HMXB | 24-Aug-2014 | 11:26:26.6 | –68:40:32.3
NOVA Sgr 1900 | HSSGR | Nova | 13-Jul-2016 | 18:28:03.4 | –21:34:24.7
NOVA Nor 1893 | ILNOR | Nova | 07-Jul-2016 | 15:29:23.1 | –50:35:00.7
NOVA Nor 1920 | IMNOR | Nova | 07-Jul-2016 | 15:39:26.4 | –52:19:17.9
V* LZ Musaa$a$Observation reported in Kamiński et al. (2015). | LZ-MUS | Nova | 24-Aug-2014 | 11:56:09.2 | –65:34:20.1
NOVA Cen 1931 | MTCEN | Nova | 30-Jun-2016 | 11:44:00.8 | –60:33:39.5
V* RR Picaa$a$Observation reported in Kamiński et al. (2015). | RR-PIC | Nova | 25-Aug-2014 | 06:35:36.0 | –62:38:24.3
V* RR Tel | RRTEL | Symbiotic* | 10-Jun-2015 | 20:04:18.5 | –55:43:33.2
NOVA Car 1895 | RSCAR | Nova | 30-Jun-2016 | 11:08:06.6 | –61:56:04.6
NOVA Pyx 1890aa$a$Observation reported in Kamiński et al. (2015). | T-PYX | Nova | 24-Aug-2014 | 09:04:41.5 | –32:22:47.5
V* TV Crvaa$a$Observation reported in Kamiński et al. (2015). | TV-CRV | DwarfNova | 24-Aug-2014 | 12:20:24.1 | –18:27:02.0
V* U Sco | USCO | CataclyV* | 10-Jun-2015 | 16:22:30.8 | –17:52:43.2
NOVA Sgr 1905 | V1015SGR | Nova | 06-Jul-2016 | 18:09:02.0 | –32:28:32.0
NOVA Sgr 1899 | V1016SGR | Nova | 13-Jul-2016 | 18:19:57.6 | –25:11:14.6
V* V1017 Sgr | V1017SGR | Nova | 13-Jul-2016 | 18:32:04.4 | –29:23:12.5
NOVA Cen 2007aa$a$Observation reported in Kamiński et al. (2015). | V1065CEN | Nova | 24-Aug-2014 | 11:43:10.3 | –58:04:04.3
V* V1148 Sgr | V1148SGR | Nova | 04-Apr-2015 | 18:09:05.8 | –25:59:08.0
NOVA Sgr 1945 a | V1149SGR | Nova | 13-Jul-2016 | 18:18:30.4 | –28:17:17.0
V* V1151 Sgr | V1151SGR | Nova | 14-Jul-2016 | 18:25:23.7 | –20:11:59.3
NOVA Sgr 1928 | V1583SGR | Nova | 14-Jul-2016 | 18:15:26.3 | –23:23:18.0
V* V351 Caraa$a$Observation reported in Kamiński et al. (2015). | V351-CAR | Mira | 28-Aug-2014 | 10:45:19.1 | –72:03:56.0
V* V359 Cenaa$a$Observation reported in Kamiński et al. (2015). | V359-CEN | DwarfNova | 24-Aug-2014 | 11:58:15.3 | –41:46:08.4
NOVA Sgr 1927 | V363SGR | Nova | 11-Jun-2015, 01-Jul-2016 | 19:11:16.3 | –29:50:00.0
V* V365 Car | V365CAR | Nova | 09-Jun-2015, 30-Jun-2016 | 11:03:16.7 | –58:27:24.9
NOVA Sco 1901aa$a$Observation reported in Kamiński et al. (2015). | V382SCO | Nova | 13-Jul-2016 | 17:51:56.1 | –35:25:05.4
NOVA Vel 1999 | V382-VEL | Nova | 28-Aug-2014, 08-Sep-2014 | 10:44:48.3 | –52:25:30.7
NOVA Sgr 1893 | V384SCO | Nova | 14-Jul-2016 | 18:01:43.1 | –35:39:27.8
V* V522 Sgr | V522SGR | Nova | 11-Jun-2015, 01-Jul-2016 | 18:48:00.4 | –25:22:21.9
NOVA Oph 1940 | V553OPH | Nova | 01-Jul-2016 | 17:42:53.5 | –24:51:26.2
NOVA Pup 2007 baa$a$Observation reported in Kamiński et al. (2015). | V598-PUP | Nova | 25-Aug-2014 | 07:05:42.5 | –38:14:39.4
NOVA Sco 1944 | V696SCO | Nova | 13-Jul-2016 | 17:53:11.5 | –35:50:14.4
NOVA Sco 1941 | V697SCO | Nova | 13-Jul-2016 | 17:51:21.8 | –37:24:55.2
NOVA Sco 1922 | V707SCO | Nova | 06-Jul-2016 | 17:48:26.3 | –36:37:54.9
NOVA Sco 1906 | V711SCO | Nova | 14-Jul-2016 | 17:54:06.1 | –34:21:15.5
V* V729 Sco | V729SCO | V* | 10-Jun-2015, 01-Jul-2016 | 17:22:02.6 | –32:05:48.8
NOVA Sgr 1936 | V732SGR | Nova | 13-Jul-2016 | 17:56:07.5 | –27:22:16.1
V* V733 Sco | V733SCO | $\rm Candidate{Mi*}$ | 06-Jul-2016 | 17:39:42.8 | –35:52:38.4
NOVA Sgr 1933 | V737SGR | Nova | 14-Jul-2016 | 18:07:08.6 | –28:44:52.3
V* V745 Sco | V745SCO | LPV* | 13-Jul-2016 | 17:55:22.2 | –33:14:58.5
NOVA Sgr 1937 | V787SGR | Nova | 14-Jul-2016 | 18:00:02.2 | –30:30:31.0
V* V794 Oph | V794OPH | Nova | 01-Jul-2016 | 17:38:49.2 | –22:50:48.9
V* V840 Oph | V840OPH | Nova | 10-Jun-2015 | 16:54:43.9 | –29:37:26.8
NOVA Sgr 1941 | V909SGR | Nova | 13-Jul-2016 | 18:25:52.3 | –35:01:27.0
V* V941 Sgr | V941SGR | Mira | 13-Jul-2016 | 18:34:43.4 | –29:34:49.1
NOVA Sgr 1910 | V999SGR | Nova | 13-Jul-2016 | 18:00:05.5 | –27:33:14.0
V* VX Foraa$a$Observation reported in Kamiński et al. (2015). | VX-FOR | DwarfNova | 24-Aug-2014 | 03:26:45.7 | –34:26:25.2
V* WX Cetaa$a$Observation reported in Kamiński et al. (2015). | WX-CET | DwarfNova | 24-Aug-2014 | 01:17:04.1 | –17:56:23.0

444Objects were observed in the CO(3–2), CO(3–2) OSB, CO(4–3), CO(4–3) OSB
setups.  Table 2: Objects observed with IRAM.

###  2.1 APEX observations and data reduction

The fifty-nine objects listed in Table 2 were observed with the Atacama
Pathfinder Experiment (APEX) 12-m submillimeter telescope (Güsten et al.
2006). Forty-five of these objects are classified as classical novae. The
observations took place in 24–28 July 2014, 4 April 2015, 8–11 June 2015, 10
July 2015, 30 July 2016, and 1, 6, 13 and 14 July 2016. For frequencies
between 278 and 492 GHz, we used the FLASH+ receiver (Klein et al. 2014),
which operated simultaneously in two atmospheric bands at about 345 GHz and
460 GHz. The version of FLASH+ separates the two heterodyne sidebands in each
of its two modules, which cover the 345 and 460 GHz bands, giving four
spectra, each 4 GHz wide. The image-band rejection factors used are all above
10 dB and at most frequencies the rejection is higher than 20 dB. In all APEX
observations, we used the eXtended Fast Fourier Transform Spectrometer (XFFTS;
Klein et al. 2012), which typically provided us with resolutions better than
0.05 km s-1. To decrease the noise to levels appropriate for the expected
broad lines, we rebinned the spectra to a much lower resolution, typically
near 30 km s-1. In all observations, we used wobbler switching with the
maximal azimuthal wobbler throw of 300″ (typically throws of 60–100″ were
used), which resulted in flat baselines.

Overall, with APEX we covered four spectral ranges characterized in Table 3.
The table includes the nominal FWHM beam size and the main-beam efficiency
($\eta_{mb}$). In the two out of four units of FLASH+, we obtained spectra
centered on the CO(3–2) and (4–3) lines with rest frequencies of 345.79598990
and 461.0407682 GHz, respectively. Also covered were the H13CO(4–3),
SiS(19–18), and SiO(8–7), transition which are readily apparent in the
spectrum of one of our the sources we observed to check the telescope’s
pointing performance, that is IRC+10216 (cf. Fig. 2). We also registered the
upper side band spectra (which we refer to as ‘CO(3–2) OSB’ and ‘CO(4–3)
OSB’). We present all results in the antenna ($T_{A}^{*}$) temperature scale.
Velocities are expressed in the local standard of rest frame (LSR).

The APEX spectra were reduced in the CLASS package within
GILDAS555https://www.iram.fr/IRAMFR/GILDAS. The data reduction involved 0th
order baseline subtraction and spectra averaging. Any faulty channels at
spectral edges or caused by instrumental effects were blanked or interpolated
between two nearest non-blanked channels. Sample APEX spectra are displayed in
Figs. 2–2 where we marked the corresponding 3$\sigma$ noise levels.

APEX data for seventeen sources have been already reported in Kamiński et al.
(2015). They are included here for completeness and consistency, and were re-
processed using the procedures presented here.

Table 3: Spectral setups used in the survey. Telescope | Setup name | Range (GHz) | Beam (″) | $\eta_{\rm mb}$
---|---|---|---|---
APEX | CO(3–2) OSB | 331.8–335.8 | 18.7 | 0.70
APEX | CO(3–2) | 343.8–348.8 | 18.0 | 0.69
APEX | CO(4–3) | 459.0–463.0 | 13.5 | 0.64
APEX | CO(4–3) OSB | 471.0–475.0 | 13.2 | 0.63
IRAM | HCN(1–0) | 83.1–91.0 | 27.8 | 0.85
IRAM | HCN(1–0) USB | 98.8–106.6 | 24.3 | 0.84
IRAM | CO(1–0) LSB | 94.1–101.9 | 25.5 | 0.84
IRAM | CO(1–0) USB | 109.8–117.5 | 21.3 | 0.83

###  2.2 IRAM observations and data reduction

Observations with the IRAM 30 m telescope and the EMIR dual-sideband receiver
(Carter et al. 2012) were obtained in 6–11 August 2015. Fifty-eight objects
observed with IRAM are listed in Table 2 and include 49 classical novae.
Spectra for each linear polarization covered the same spectral range. The Fast
Fourier Transform Spectrometer (FTS, Klein et al. 2012) and the WILMA
correlator were used in parallel as backends. The FTS was used as the primary
spectrometer, while WILMA data were used for reference since it covered a
smaller part of the band (7.5 GHz) than FTS (16 GHz). WILMA spectra, as free
from a platforming effect, were analyzed carefully, too. All IRAM observations
were obtained using wobbler switching with OFF positions located azimuthaly
110″ away from the central ON position.

The spectral setups used are listed in Table 3. All objects were observed with
the setup whose lower side band was centered on the HCN(1–0) transition at
88.63160230 GHz; it serendipitously covered also the 1–0 transitions of H13CN
and HNC (cf. Fig. 2). For six objects, additional spectra were obtained in the
12CO(1–0) transition at 115.2712018 GHz; the corresponding lower sideband
spectra covered also the 13CO(1–0) transition at 110.20135430 GHz.

All data reduction procedures were performed using CLASS and involved a low-
order baseline subtraction and averaging. The FTS spectra required applying a
deplatforming routine. While searching for broad nova lines, we used Hanning
smoothing to reduce the resolution to about 30 km s-1. Spectra of all objects
observed in the 12CO(1–0) transition and a sample of spectra with the HCN(1–0)
setup are displayed in Figs. 2–2 together with marked 3$\sigma$ noise levels
calculated for the entire spectrum. Owing to the changing atmospheric
transparency with wavelength, the noise levels change over the spectrum,
especially in the CO(1–0) setup (cf. Fig. 2).

###  2.3 Single-dish survey results: no detection of nova lines

To search for broad lines expected from classical novae outflows, we analyzed
the spectra at a reduced spectral resolution of about 30 km s-1. Except for CK
Vul, known to have broad molecular emission from earlier studies, none of our
survey targets shows broad features that can be associated with a nova system.

The measured rms noise levels for APEX spectra are listed in Table 2.3. The
rms noise levels for the CO(3–2) spectra are relatively low, of 1–6 mK. The
observations of CO(4–3) are much less constraining due to a lower atmosphere
transparency at the respective frequencies. The IRAM observations in the
CO(1–0) line for a handful of objects are comparatively deep to the best APEX
spectra, with global rms noise levels listed in Table 6. The noise levels are
several times smaller near the rest frequency of the CO(1–0). The upper limits
on HCN emission can be derived from the rms values listed in Table 2.3.

The southern part of the survey constitutes the largest sample of novae and
related objects that were checked for the presence of rotational lines of CO.
The noise levels achieved in our CO survey are one order of magnitude lower
than in earlier observations of classical novae (e.g. Shore & Braine 1992).
The IRAM survey of northern sources in the HCN(1–0) transition is the first
one where lines other than those of CO were targeted for a large sample of
classical novae.

Table 4: Noise levels in APEX observations. 666Root-mean-square values are
given per velocity bins of 30 km s-1 for CO(3–2) and CO(4–3), 31 km s-1 for
CO(3–2) OSB, and 29 km s-1 for CO(4–3) OSB. Table 5: Noise levels in IRAM
spectra near the HCN(1–0) line. 777The noise rms values per 30 km s-1 are
given for both backends used (they differ in covarage and baseline quality).
The LSB values correspond to the spectra which covered the HCN(1–0)
transition. Table 6: Noise levels in IRAM spectra near the CO(1–0) line.

| rms (mK)
---|---
| WILMA | FTS
Object | LSB | USB | LSB | USB |  CO10aa$a$The rms value measured in 114.27–116.27 GHz, near the rest frequency of the CO(1–0) transition and ignoring the ISM features.
V* BC Cas | 2.21 | 3.14 | 2.38 | 16.06 | 6.22
V* CP Lac | 2.29 | 3.47 | 2.22 | 10.39 | 6.87
V* DI Lac | 2.13 | 3.12 | 2.27 | 11.10 | 7.56
V* V630 Cas | 1.62 | 2.28 | 1.57 | 6.23 | 4.71
V* V723 Cas | 2.13 | 2.88 | 2.16 | 8.81 | 6.68
NOVA Per 1887 | 2.11 | 2.96 | 2.52 | 10.15 | 7.09
888The $T_{A}^{*}$ noise rms values per 30 km s-1 are given for both backends
used (they differ in covarage and baseline quality). The USB values correspond
to the entire spectra which covered the CO(1–0) transition.

Our upper limits on the molecular emission are not easily convertible to the
limits on the total molecular mass because no direct information about the
flow velocity, envelope size, and gas excitation temperature are known for
individual objects. Here, we make an attempt to make such a conversion for a
model nova representative of our survey. By referring to Gaia distances to
nearly half of our sources in Schaefer (2018), we find that our typical
classical nova is at a distance of 2 kpc. We further assume a typical flow
velocity of $\Delta V$=1000 km s-1. Then, with an age of our typical nova
remnant of about 50 yr, the shell has a size of 0.1 pc. In the optically thin
limit, the mass of molecular material can be calculated as M_mol=8 πk
μmHZ(Tex) ν2exp(Eu/Tex) ηmbh c3guAulWa ( θbeamθsource)^2, where W=∫T_A^*
d≈3×rms×2×ΔV is the integrated intensity within the line profile, assuming a
rectangular feature with a peak intensity equal to three times the measured
noise level. $k$, $h$, and $c$ are the Boltzmann and Planck constants, and the
light speed, respectively; $A_{ul}$, $g_{u}$, $E_{u}$, and $\nu$ are Einstein
coefficient, upper level statistical weight, energy (in K) of the upper level,
and frequency of the considered transition; $Z$ is the partition function at
the excitation temperature $T_{\rm ex}$; the angular beam size is designated
as $\theta_{\rm beam}$ while the source angular size is $\theta_{\rm source}$;
$\mu$ is the molecular weight and $m_{\rm H}$ is the mass of hydrogen atom;
$a$ is the source projected area; both $a$ and $\theta_{\rm source}$ depend on
the distance. We considered excitation temperatures of 10, 75, and 150 K to
calculate rough upper limits on the total masses of CO from the CO(1–0) and
(3–2) transitions and of HCN based on the HCN(1–0) spectra. The results are
presented in Table 7.

Table 7: Upper limits on molecular mass in a model classical nova.

Observed |  3$\times$rms | $T_{\rm ex}$ | Mass limit
---|---|---|---
transition |  ($T_{A}^{*}$, mK) | (K) |  (M☉)
CO(3–2) | 6 | 10 | $M_{\rm CO}<9.0\times 10^{-5}$
|  | 75 | $M_{\rm CO}<3.7\times 10^{-5}$
|  | 100 | $M_{\rm CO}<5.8\times 10^{-5}$
CO(1–0) | 18 | 10 | $M_{\rm CO}<3.0\times 10^{-4}$
|  | 75 | $M_{\rm CO}<1.4\times 10^{-3}$
|  | 100 | $M_{\rm CO}<2.6\times 10^{-3}$
HCN(1–0) | 9 | 10 | $M_{\rm HCN}<1.0\times 10^{-6}$
|  | 75 | $M_{\rm HCN}<5.3\times 10^{-6}$
|  | 100 | $M_{\rm HCN}<1.0\times 10^{-5}$
999The 3$\times$rms values are based on Tables 2.3–6.

The most constraining were the sensitive observations in CO(3–2), which
indicate that the CO mass must be lower than 10-5 M☉. For a CO fractional
abundance relative to H I+H2 of 10-4 (cf. Pontefract & Rawlings 2004), one
obtains upper mass limits on molecular matter on the order of 0.1 M☉. This
value is not that constraining, given that a typical classical nova is thought
to eject a total mass of $\lesssim$10-3 M☉. However, the CO fractional
abundance in a gas heavily enhanced in CNO burning product does not have to be
close to that known from other circumstellar environments. The fractional
abundance of HCN in classical novae is even more uncertain (cf. Schöier et al.
2013), but the upper limits we get are comparable or worse in constraining the
molecular mass than the CO(3–2) data. Molecular gas in classical novae is not
easily observable with current single-dish telescopes.

Our survey can also be compared to observations of the ro-vibrational band of
CO at NIR wavelengths. As an example, we choose the case of V5668 Sgr whose
prominent CO first-overtone emission was observed about a week after maximum
light and was analyzed in Banerjee et al. (2016). The emission is thought to
arise from a source of the size of 42 mas, in an outflow with and excitation
temperature of about 4000 K and an expansion velocity $\Delta V$=530 km s-1.
At the CO column density of 2$\times$1019 cm-2 and assuming LTE conditions,
the corresponding mm-wave emission in CO(1–0) observed by IRAM would have a
peak antenna temperature of 0.5 $\mu$K. Analogously, the CO(3–2) line emission
observed with APEX would reach 43 $\mu$K ($T_{A}^{*}$). These extremely low
values are a consequence of the high temperature and small source size,
illustrating how ambitious it is to detect molecular emission from a nova
shortly after the outburst. Let us however assume that the CO gas observed at
NIR wavelengths survives for the next 50 years, decreases its gas temperature
to 30 K, and expands into a source of a size of
2$\aas@@fstack{\prime\prime}$75\. The size was calculated for free gas
expansion at the distance of 2 kpc (cf. Banerjee et al. 2016). For such a
source, the predicted peak antenna temperatures are of 0.1 and 0.3 K, well
within the reach of our survey. Admittedly, V5668 Sgr is the most optimistic
case with a particularly high molecular abundance and we have arbitrarily
chosen the excitation temperature at which mm and submm lines are strong.
Given the large number of observed novae, we would expect however to see a few
objects with emission above our detection limits. That we did not detect any
such source is possibly due to the decay of CO in later phases of ejecta
expansion and much smaller filling factors (or a smaller source size) than we
assumed. Indeed, chemical models of CO evolution of Pontefract & Rawlings
(2004) predict high CO abundances only till $\approx$10 days after the
eruption, after which they drop by at least two orders of magnitude. With some
caution, results of our survey give qualitative support to such behavior.

###  2.4 Red-nova connection

The survey can be interpreted also in the context of red novae research. Nova
1670 or CK Vul had long been considered a classical nova until in 2014 strong
mm-wave molecular emission was discovered in the object (Kamiński et al.
2015). This and other observational characteristics of the object led to the
interpretation that instead of being a classical nova, CK Vul is a remnant of
a stellar merger (Kamiński et al. 2018b, 2021) and its 1670–72 outburst can be
classified as belonging to the group of (luminous) red novae (Kamiński et al.
2018a; Pastorello et al. 2019). Our survey can then be treated as a search for
CK Vul analogs among historical novae, some of which could have been
misidentified, especially if the nova event was not followed up
spectroscopically sufficiently long. Since this experiment gave a negative
result, objects similar to CK Vul must be very rare in the Galaxy. Indeed,
luminous red novae outbursts are expected to occur only twice a decade (Howitt
et al. 2020a). More systematic searches of red novae among historical objects
classified as classical novae can improve our statistics of stellar mergers
(cf. Kochanek et al. 2014; Howitt et al. 2020b).

Among the targets observed with APEX is V1148 Sgr, a nova from 1943 reported
by Mayall (1949) and whose observations suggest that it could have been a red
nova (Kimeswenger 2007). The exact position of the nova is uncertain. Although
all the disputed positions of the remnant were well within the APEX beam, we
do not detect any strong circumstellar emission in the CO transitions. This
strongly advocates against a red nova classification of V1148 Sgr.

###  2.5 Observations of interstellar clouds

Figure 1: The IRAM spectrum of BC Cas with ISM features of CO(1–0). The top
axis gives LSR velocity with respect to the rest frequency of the CO
transition. The dashed lines indicate the local 3$\sigma$ noise level.

Data gathered for our entire sample of targets resulted in no detection of
broad molecular emission that would be expected from classical novae shells.
However, in the spectra of ten targets we observed narrow features which most
likely arise in the foreground or background interstellar medium (ISM). This
is not surprising as these objects are located close to the Galactic plane.
The spectra are shown in Figs. 2–2. Some narrow spectral features appear in
“absorption” when observed in one of the off-source wobbler positions. Below,
we summarize the detected ISM features.

  * •

In the direction of BC Cas, we detected 12CO(1–0) features near –3 and –42 km
s-1 (Fig. 2), the stronger of which is also detected in the isotopolog line of
13CO(1–0) at $\approx$–42 km s-1(Fig. 2). The main line has the highest peak
intensity among all observed ISM features in our survey and is optically thick
based on the intensity ratio of the two observed CO isotopologs. This is in
accord with BC Cas being a highly-reddened nova ($A_{V}$=3.7 mag) seen towards
a star-forming region (Kato & Kojiguchi 2020).

  * •

The spectra of CP Pup, RS Car, and AT Sgr show narrow emission and absorption
features of CO(3–2) which combine into a P-Cyg type profiles (Fig. 2). This
profile shape is a consequence of a small velocity difference between the ISM
cloud velocity at the ON and OFF wobbler positions. The lines have a FWHM of
0.7–1.1 km s-1, typical for the translucent ISM (cf Snow & McCall 2006).
Emission is observed at 8.8 km s-1 in CP Pup, at –11.8 km s-1 in RS Car, and
at 12.0 km s-1 in AT Sgr. The lack of the corresponding CO(4–3) lines at our
sensitivities (implying a line ratio of $\gtrsim$2) is consistent with origin
in a diffuse or translucent cloud.

  * •

An absorption line centered at $\approx$–38 km s-1 and with a FWHM of 2 km s-1
is visible in the CO(3–2) spectrum of the recurrent nova IM Nor. A similar
CO(3–2) feature is seen in the spectrum of V729 Sco at 19 km s-1. The
responsible ISM clouds are located in one of the wobbler OFF positions and do
not necessarily cross the direction towards the novae.

  * •

In the spectrum of a reddened nova V365 Car (Zwitter & Munari 1996), we found
emission and absorption combining into an inverse P-Cyg profile. This again is
a consequence of the ISM cloud extending between the ON and OFF wobbler
positions with a small velocity gradient. Both emission and absorption
components appear double, but the absorption component (OFF) is stronger. The
emission component is centered near –25 km s-1. No CO(4–3) emission is seen.

  * •

The spectrum of V1148 Sgr shows two emission components, one centered at 21 km
s-1 and with a FWHM of 3 km s-1, and another narrow feature combining
absorption and emission in 6–9 km s-1. The emission lines are the weakest
features detected in our survey.

  * •

Two targets, AP Cru and V732 Sgr, show ISM features in both the CO(3–2) and
CO(4–3) transitions (Fig. 2). While the line intensity ratio in AP Cru is not
unusual for cool low-density ISM, the emission towards V732 Sgr is almost
equal in both transitions, indicating a warmer interstellar environment. The
case of V732 Sgr is discussed in more detail below.

Figure 2: APEX spectra with ISM features of CO(3–2). The top axes give LSR
velocity with respect to the rest frequency of the corresponding CO
transition. Dashed lines indicate local 3$\sigma$ noise levels. The rms values
are provided for the baseline outside the ISM features and for the current
binning.

Figure 3: Spectra with ISM features detected on CO(3–2) (left) and CO(4–3)
(right) in AP Cru (top) and V732 Sgr (bottom). The top axes give LSR velocity
with respect to the rest frequency of the corresponding CO transition. Dashed
lines indicate local 3$\sigma$ noise levels.

####  2.5.1 Emission towards V732 Sgr

Spectra of V732 Sgr centered at the CO(3–2) and (4–3) transitions are shown in
the two bottom panels of Fig. 2. Both spectra display a broad (FWHM$\approx$
12 km s-1) line centered near 62 km s-1 and several very narrow lines that
appear as emission or absorption between 0 and 22 km s-1. Additionally, the
spectrum centered on the CO(3–2) transition shows a marginally significant
feature near 111 km s-1. Given the relatively broad profile of the CO lines
near 62 km s-1 and a relatively high intensity of the CO(4–3) transition, this
component is reminiscent of lines observed from envelopes of asymptotic-giant-
branch (AGB) stars. However, CO emission towards V732 Sgr is known since the
early 1990s, when it was mapped with SEST in the CO(1–0) transition, to be of
interstellar origin (Weight et al. 1993; Albinson et al. 1994). The maps show
that the molecular cloud is extended (10′$\times$3.5′) and massive ($10^{4}$
M☉), and thus it cannot be directly associated with the nova. Interestingly,
however, $\approx$100 days after V732 Sgr erupted in 1936, a weak light echo
was visible for about 30 days (Swope 1940). The echoing medium was most likely
associated with the giant interstellar cloud observed with SEST with APEX. The
presence of light echo make V732 Sgr very similar to the nearby Nova Persei
1901 (GK Per), which we discuss in more detail in Sect. 3. All the novae for
which ISM emission has been detected (including NQ Vul; Nielbock &
Schmidtobreick 2003) are candidates to have been associated with a light echo
(some could have been easily overlooked). Except GK Per and V732 Sgr, none of
our sources with ISM detection was targeted for an echo search (van den Bergh
1977; Schaefer 1988). The light echo can arise only in clouds relatively
close, up to tens of parsecs (Schaefer 1988).

##  3 CO maps around GK Per

###  3.1 Motivation

GK Per is a unique cataclysmic variable whose distance of 442 pc (Schaefer
2018) makes it the second-closest nova. The central binary consists of a
magnetic white dwarf of 0.9–1.0 M☉ (Wada et al. 2018; Álvarez-Hernández et al.
2021) and a K1 IV subgiant (Álvarez-Hernández et al. 2021). It erupted as a
classical nova in 1901 and, as mentioned, is one of two known classical novae
to had been associated with light echos. The echoes reached a size of tens of
arcseconds (Ritchey 1902). Since the middle of the last century, GK Per has
experienced dwarf-nova outbursts every 1–3 yr (Šimon 2002). Nova ejecta was
first observed in 1916 and remains visible to this day (Liimets et al. 2012;
Shara et al. 2012). The Firework Nebula, as it is often referred to, has a
size of nearly 100″ and is considered the longest-lived remnant of a classical
nova event. The highest gas velocities found in the remnant are of 2800 km s-1
but typical values are of about 1100 km s-1.

GK Per is the only classical nova that has been claimed to have molecular gas
in its immediate proximity long after its outburst101010Evans (1997) claimed
the presence of H2 emission in DQ Her some 56 years after its eruption and
based on imaging observations. The emission identification has never been
confirmed spectroscopically, though.. The nature of its observed CO emission
has been somewhat controversial. Observations of the CO $J=2-1$ rotational
line with the James Clerk Maxwell Telescope (JCMT) acquired in 1993 by Scott
et al. (1994) revealed elongated symmetrical structures with two main
“concentrations” connected by a bridge of emission $\approx$1′ north of the
object. The concentrations are located $3^{\prime}$ east and west off GK Per.
However, the emission region was not fully mapped; in particular, the
concentrations are seen near the edges of the undersampled map, which puts
into question the extent of the emission. Observed lines are narrow, with a
FWHM of a few km s-1, atypical for circumstellar envelopes. No large velocity
gradients within the cloud were found. The molecular gas was proposed to be
associated with a filamentary and dusty bipolar cloud of a size of 17′ from
which dust emission was first observed with the Infrared Astronomical
Satellite (IRAS). This filamentary cloud was interpreted by Bode et al.
(1987), Seaquist et al. (1989), and Dougherty et al. (1996) as a bipolar
outflow or a planetary nebula (PN) from the nova system. Their argument was
based mainly on the symmetry of the cloud’s morphology relative to GK Per. An
atomic component of the cloud was also mapped in the 21 cm line of neutral
hydrogen (H I) by Seaquist et al. (1989). In order to explain the very large
extent and the very high mass of 3.8 M⊙ derived for the cloud, both highly
atypical for classical novae, Dougherty et al. invoked a rather elaborate
scenario where mass loss involving the secondary had occurred prior to the
Nova 1901 event. They considered the binary to be then in the second Roche-
lobe overflow phase, in which the mass transfer from the secondary, then of a
mass of 1.3–1.5 M☉, was sufficient to convert the white dwarf into a “born
again” AGB star. Knowing the composition of the nebula would be of great value
in order to verify this hypothesis. If the CO emission observed with the JCMT
indeed represents such circumstellar gas, it could provide means to measure
relevant isotopic ratios.

On the other hand, the CO emission near GK Per may be purely interstellar. In
1985 Hessman (1989) mapped in the CO(1–0) line the large-scale structure that
was identified in IRAS maps. With the 4.9-m telescope of the Millimeter Wave
Observatory, he mapped a much larger area than was covered with JCMT, roughly
of 40′$\times$50′ and at a resolution of 2$\aas@@fstack{\prime}$6\. He only
found a weak correspondence between the CO emission and the IRAS emission,
which would be atypical for any circumstellar material. Hessman concluded that
most of the CO material belonged to an interstellar cirrus cloud that is
unrelated to the nova. This interpretation is not favored in the recent
literature on GK Per.

Figure 4: WISE images of the field around GK Per. Top: A wide-angle view
(1°$\times$1°) with blue, green, and red colors representing signal in $W2$,
$W3$, and $W4$ filters, respectively. It shows the large cirrus IS cloud
reported earlier in IRAS observations. Middle: Closeup view with blue, green,
and red showing $W1$, $W2$, and $W4$ images, respectively. In yellow, a
bipolar structure is seen near the position of GK Per. Bottom: $W2$ image of
the bipolar structure. Its shape was arbitrarily outlined with a dashed line.
The northern edge is not well-defined. The field shown
(17$\aas@@fstack{\prime}$0$\times$14$\aas@@fstack{\prime}$5) corresponds to
the white boxes drawn in two upper panels. All panels show intensity in linear
scale.

Most recently, the vicinity of GK Per was mapped with the Wide-field Infrared
Survey Explorer (WISE Wright et al. 2010) at a higher sensitivity and with
better angular resolutions than with IRAS. WISE maps in $W3$ and $W4$ bands,
at 12 and 22 $\mu$m, respectively, readily show the same extended dusty cloud
that was identified in IRAS data (Fig. 4). It has the appearance of an
interstellar cirrus. However, at shorter wavelengths, that is at 3.4 $\mu$m
($W1$) and 4.6 $\mu$m ($W2$), WISE maps display a smaller bubble-like, bipolar
structure centered on GK Per. Although weak in both bands, it is most readily
present in the $W2$ image. The longer axis of the nebula is of about 15′ and
is oriented at a position angle of $\approx$–30°. The structure appears
irregular, with the southeastern lobe being bigger and of a width of about 7′.
The same bipolar structure was observed earlier in the H$\alpha$ filter
(covering the H$\alpha$ line and the [N II] doublet) and in the [O III]
$\lambda$5007 line by Tweedy (1995)111111Their Fig. 1 has an incorrectly
assigned orientation. The maps have east up and north right. and Bode et al.
(2004). The observations imply that the optical emission is dominated by line
emission of recombining low-density gas. Since the $W1$ and $W2$ bands also
cover several emission lines, we cannot dismiss the possibility that the
bipolar WISE structure is seen through IR emission lines. However, the
southern part of the bipolar structure was seen in the light echoes in
1901–1902 and thus must contain dust grains. We assume that the mid-IR
emission of this structure arises as continuum emission of relatively warm
dust, warmer than the larger ISM cirrus seen in $W3$ and $W4$. Tweedy excluded
the possibility that the observed feature represents a genuine PN, but later
authors (e.g., Bode et al. 2004; Harvey et al. 2016; Dougherty et al. 1996;
Álvarez-Hernández et al. 2021) challenged this conclusion. This nebula is
smaller and much less massive than the dusty cloud seen in the IRAS images,
but both structures have often been confused in the literature and named a
planetary nebula interchangeably. Hereafter, we refer to the smaller structure
as a planetary or bipolar nebula, and we call the larger-scale emission a
cirrus.

Based solely on the morphology of the smaller nebula, there is little doubt
that it is a signature of mass-loss activity in GK Per before the Nova 1901
event. The longer axis of the structure aligns well with the larger
interstellar cirrus, perhaps causing the confusion in the literature. The
molecular gas mapped earlier with the JCMT and reported in Scott et al.
(1994), would be filling the bipolar lobes of the planetary nebula, with the
concentration located closely to the main axis of the bipolar structure. This
strongly suggests circumstellar origin of the CO gas, against the
interpretation of Hessman (1989). With the presence of a PN now firmly
established with the WISE maps, we decided to revisit the claim of the
presence of circumstellar molecular gas around GK Per by obtaining new deeper
maps of CO emission near GK Per.

###  3.2 Mapping observations

Mapping observations around GK Per were executed in 9, 10, 12 April 2021 and
on 28 May 2021 with the IRAM 30 m telescope. The data were acquired within the
IRAM program 155-20 (PI: H. Mazurek). We used the EMIR E90 and E230 receivers
simultaneously and centered them near the CO $J=1-0$ and $2-1$, transitions,
which have rest frequencies of 115.2712 and 230.5380 GHz, respectively. The
FTS and WILMA backends were used as spectrometers with the FTS having a native
resolution of 0.5 and 0.25 km s-1 for the CO(1–0) and (2–1) spectra,
respectively. Although effectively the spectral ranges of 113.522–117.570 GHz
and 228.240–236.016 GHz were covered, no lines other than the targeted 12CO
transitions were detected. Regular pointing observations were obtained at the
beginning and during mapping runs. Focus calibration on a planet was secured
before each observing session.

The maps were obtained in the on-the-fly (OTF) mode, with a reference OFF
position at (RA, Dec) offsets of (700″, 350″) from the position of GK Per. The
reference position was checked earlier to be free of any molecular emission.
An area within a field of view of
1128$\aas@@fstack{\prime\prime}$8$\times$987$\aas@@fstack{\prime\prime}$0 was
mapped in two smaller parts, each with multiple coverages obtained in
orthogonal directions of the beam scanning. Relative to the position of GK
Per, the map is composed of two rectangular regions at (RA, Dec) offsets of
(610″ to 10″, –410″ to 210″) and (80″ to –450″, –150″ to 500″) and with some
overlap between the two areas near the position of the nova. The map sampling
was set to 8$\aas@@fstack{\prime\prime}$5, which is optimal for the beam size
of the CO(1–0) observations (21$\aas@@fstack{\prime\prime}$3). Owing to
several coverages of each part of the map at a slightly different offset (of a
few arcsec), the grid of observed positions is effectively denser, allowing
also to obtain a map of CO(2–1) that, with the beam size of
10$\aas@@fstack{\prime\prime}$7, is only slightly undersampled. The data were
calibrated every 8 min and registered in $T_{A}^{*}$ units. The conversion of
the CO(2–1) data to the main beam temperature units requires $\eta_{\rm
mb}$=0.64 (cf. Table 3). All data were reduced in CLASS, including baseline
subtraction, spectra averaging and rebinning, and map gridding. The final rms
noise in the CO(1–0) cube is of 0.22 K ($T_{A}^{*}$) at 0.45km s-1 binning;
for the CO(2–1) cube, it is 0.41 K per 0.45 km s-1.

###  3.3 Results of CO mapping

At the position of GK Per, no broad emission lines are found that could
originate in the nova system. A narrow and weak CO(1–0) line is seen at this
position and originates in the interstellar cloud.

Figure 5: Total-intensity emission of CO(1–0) near GK Per. The cross marks the
position of the nova. Contours are at 3, 6, 9 times the rms noise level. Left:
Color map of the CO emission. Right: The same CO contours are plotted as red
solid lines over the $W2$ image of the field (cf. Fig. 4). The CO emission
does not follow the dust emission seen in the $W2$ image.

Both maps show interstellar clouds emission over most of the mapped area and
in the 1.5–8.0 km s-1 range. Here we focus on the better sampled CO(1–0) map,
as the CO(2–1) map looks qualitatively the same. The total intensity CO(1–0)
map is shown in Fig. 5. It displays a ridge of molecular emission spreading
over most of the mapped area. The ridge appears to be broken near the position
of the nova. The orientation of the emission at RA offsets from 500″ to the
west-most edge of the map is close to horizontal. At larger offsets towards
the east, the ridge turns southeast. This southeastern part is also the
brightest fragment of the mapped filament. It very likely continues beyond the
mapped region. Most of the emission belongs to the same kinematic structure
within 1.3–4.5 km s-1. However, the channel maps in Fig. 6 show weak emission
near offsets (500″, 50″) at higher velocities, 5.0–7.7 km s-1, thus revealing
another overlapping cloud. The main ridge has a small velocity gradient, with
velocity increasing from east to west. Only near 3.6 km s-1, very weak
emission is seen close to the position of the nova.

Our CO maps cover an area that is larger than the map of Scott et al. (1994)
but smaller than that of Hessman (1989). The sampling, resolution, and
sensitivities of our maps are far superior to the earlier studies and show the
region in a greater detail. It is clear that the incomplete coverage of the
field in the maps of Scott et al. (1994) lead them to believe that they saw a
point-symmetric structure originating from GK Per. Our maps reveal that the
emission they found is only a fragment of a longer ridge that most likely is
much larger than the field covered by our map and is part of the large-scale
complex mapped by Hessman.

The ratio of the main-beam intensities of the CO(1–0) to (2–1) emission
integrated over the entire line profile is between around 1.2 and 1.5 over
most of the ridge surface, with the lower values representing the brightest
southeastern part. The varying ratio reflects temperature changes coupled with
variations in the line optical depth. Without measurements in one extra
transition, neither can be unambiguously determined. Under the assumption of
the local-thermodynamic equilibrium, however, we calculate that the typical
excitation temperature is of about 12 K. Near the region with peak emission,
southeast of GK Per, the CO column densities are about 1.0$\times$1016 cm-2
and optical depths are about 3 or lower. In the brightest patch west of the
star, the column densities are about three times lower. If the emission arises
in a diffuse medium of low density, CO may be subthermally excited and a non-
LTE approach would necessary for column density determinations. More
transitions of CO or extra measurements in 13CO would constrain the physical
conditions better. We did not manage to measure 13CO emission, but Hessman
found a 12CO to 13CO emission ratio of 15 south of the area mapped by us,
where the emission is stronger and more optically thick. This ratio is thus
consistent with the interstellar nature of the emitting medium.

Figure 6: Left: Channel maps of CO(1–0). The map LSR velocity is given in the
upper left corner of each panel. GK Per is located at the offset of (0, 0).
The maps show two clouds at 1.8–4.95 km s-1 and 5.4–7.65 km s-1. Right:
Representative spectra. Top panel compares profiles of CO(1–0) at two
positions east and west of GK Per as indicated in the legend. Hardly any shift
is seen, unlike ones expected for any circumstellar material. Bottom panel:
Line profiles of the two lowest CO transitions are compared for the eastern
position where emission is the strongest.

###  3.4 Interpretation

The mapped CO cloud has every characteristic of an interstellar translucent
cloud (cf. Snow & McCall 2006). The mass of the molecular material seen in our
maps can be estimated using the X-factor method. We use the $X_{\rm CO}$ value
of 2.0$\times$1020 derived for translucent clouds (Magnani & Onello 1995).
Within 5$\sigma$ contours of the total intensity maps of CO(1–0), which
includes the western and eastern parts of the ridge, we measured an average
intensity of 4.9 K km s-1 (in $T_{\rm mb}$) and an area of 4.1$\cdot$104
arcsec2 or 1.8$\cdot$1036 cm2 (at 442 pc). This yields 1.8$\cdot$1057
particles of hydrogen or a mass of 3 M☉. As we conservatively included only
emission above 5$\sigma$, this is only a lower limit on the mass. The entire
complex mapped by Hessman must be several times more massive still. Such a
massive complex could not originate from the low mass cataclysmic binary GK
Per.

The narrow lines that we observe in the field of GK Per are very typical for
the low-density interstellar medium and would be highly unusual for any
circumstellar envelope, practically ruling out the latter interpretation. As
shown in Fig. 6, the peak velocities almost do not change over the different
parts of the filament. For a circumstellar outflow, this would only be
possible if the filamentary cloud were flowing out of the system in exactly
the sky plane, which is highly unlikely (as even then a non-zero opening angle
would lead to some spread in velocity). The kinematics of the molecular region
is fully consistent with a quiescent interstellar cloud.

The distribution of CO emission does not correlate strongly with that of the
dust emission seen in $W3$ and $W4$ maps, or with the emission of H I
(Seaquist et al. 1989). A strong correlation between the three phases is not
expected for the low-density ISM (e.g., Planck Collaboration et al. 2011;
Reach et al. 2015). For instance, for the translucent cloud responsible for
the light echoes of V838 Mon no such strong correlation is found either (see
discussion in Kamiński et al. 2011).

What is then the relative location of the interstellar molecular cloud
relative to GK Per? Already Hessman (1989) noticed that the harsh radiation of
the nova remnant, including the dwarf novae outbursts, may photodissociate the
local interstellar clouds out to a radius of 0.5 pc or 3′. Were the gap in the
filament that appears to be centered on GK Per not coincidental but caused by
the nova radiation, the interstellar filament must be physically close to the
nova system. This would be also consistent with the 1901 light echo arising in
the dusty counterpart of the CO cloud. Since forward scattering is dominating
the echo generation, the scattering cloud is likely located slightly closer to
us than GK Per. As noted by earlier authors, observing an evolved object such
as GK Per close to an interstellar cloud is rather unexpected and should be
rare. On the other hand, a few such cases are known, including the cepheid RS
Pup (Kervella et al. 2012) or V732 Sgr (Sect. 2.5.1).

The systemic heliocentric velocity of GK Per is 32 (28–41) km s-1 (Álvarez-
Hernández et al. 2021, and references therein) or 24 km s-1 in the LSR frame.
This velocity is quite different from that of the observed CO emission in the
field (1.5–8 km s-1). If these velocities are used as distance indicators, the
two objects could be considered as located far apart along the line of sight.
On the other hand, GK Per may be moving through the local interstellar medium
with a velocity that does not follow closely the overall Galactic rotation
pattern. If physically related with the cloud, GK Per would have a relative
radial velocity of $\approx$15 km s-1 (plus unknown relative tangential
motion) which is supersonic and sufficient to produce a shock. Based on the
proper motion of the nova, Bode et al. (2004) found a space velocity (proper
and radial combined) of 45 km s-1. As was noticed by Bode et al. (2004), the
bipolar nebula bears some signatures of interactions with the interstellar
medium, but we find no other evidence of strong interaction with the ISM.

The bipolar planetary nebula seen in WISE maps and in optical emission lines
is most likely circumstellar in origin. As we have shown, it does not possess
a molecular component, which is common among genuine PNe of various types and
ages (e.g. Bublitz et al. 2019; Guzman-Ramirez et al. 2018). The material
filling the bipolar shell may be interacting with the inner nova ejecta, as
suggest in earlier studies (e.g., Harvey et al. 2016). In particular, the jet
identified within the ejecta is directed near the plane of the equatorial
waist of the WISE structure.

In conclusion, near GK Per we find only CO emission associated with the
surrounding ISM. Elucidating the nature of the bipolar (planetary) nebula
associated with GK Per is very important for understanding the full evolution
of this rather bizarre cataclysmic system. It is thus very unfortunate that it
cannot be studied through mm-wave molecular lines.

##  4 Interferometric search of molecules in novae

It appears that all single-dish observations of classical novae have resulted
in non-detections of molecular gas. This may be a consequence of insufficient
sensitivities or problematic baselines. Millimeter-wave interferometric arrays
are usually more sensitive and offer flat spectroscopic baselines, and thus
provide the most optimal way to test classical novae for the presence of
molecules (Sect. 1). Here, we present a few experiments targeting classical
novae that were conducted with modern interferometric millimeter arrays.

A dedicated search for CO and HCN emission in four classical novae, NQ Vul, LW
Ser, V2676 Oph, and V3662 Oph, was performed with the Submillimeter Array
(SMA). We used the SWARM correlator (Primiani et al. 2016) which effectively
covered four spectral ranges 224.5–232.5, 240.5–248.5, 336.0–344.0, and
352.0–360.0 GHz. The spectra thus cover the CO(2–1) transition at 230.538 GHz
and the HCN(4–3) transition at 354.505 GHz. Each source was observed
interchangeably with two nearby quasars, whose observations were later used
for phase calibration. In each observing run, a bright quasar or a planet were
observed for bandpass and absolute-flux calibration. The SMA data were
calibrated in MIR121212https://lweb.cfa.harvard.edu/rtdc/SMAdata/process/mir/
and further processed (including data flagging and imaging) in CASA. We
typically observed with 7–8 SMA antennas and at baselines from 9 to 77 m. The
details and results of these observations are listed in Table 8.

None of the novae observed with the SMA was detected in molecular lines. One
object, V3662 Oph or Nova Oph 2017, was observed 2 months after its discovery
in outburst and about a month after detection of its near-IR first-overtone
band of CO (Joshi & Banerjee 2017). In SMA data, only free-free continuum was
detected in the 224–249 GHz window for this nova (Kamiński & Gehrz 2017). We
achieved the best sensitivity, of 9.6 mJy, in observations of V2676 Oph in the
224–249 GHz spectra. This observation is several times deeper than our best
APEX observations in CO(3–2) at an rms noise of 2 mK or about 82 mJy.

Table 8: Novae observed with the SMA.

Object | Alter. name | RA (J2000) | Dec (J2000) | Obs. dates |  BeamCO |  BeamHCN |  rmsCO |  rmsHCN
---|---|---|---|---|---|---|---|---
|  |  |  | 2017 | (″) | (″) | (mJy/beam)
NQ Vul | Nova Vul 1976 | 19:29:14.75 | +20:27:59.63 | 19,20 Jun; 5,9 Jul |  3.7$\times$3.0 |  3.4$\times$2.3 | 13.9 | 15.2
LW Ser | Nova Ser 1978 | 17:51:50.89 | –14:43:50.60 | 5 Jul |  4.0$\times$3.0 |  2.6$\times$2.1 | 14.7 | 36.7
V2676 Oph | Nova Oph 2012 | 17:26:07.08 | –25:51:45.40 | 19 Jun |  4.5$\times$3.1 |  2.9$\times$2.0 |  9.6 | 27.1
V3662 Oph | Nova Oph 2017 | 17:39:46.08 | –24:57:55.50 | 20 Jun |  4.4$\times$3.1 |  2.8$\times$2.0 | 13.2 | 31.9

131313Beam sizes are given as FWHM at natural weighting of visibilities. The
rms values, rmsCO and rmsHCN, are specified for a spectral bin of 30 km s-1
near the rest frequencies of CO(2–1) and HCN(4–3), respectively.

ALMA offers the greatest sensitivity among all currently available mm-wave
instruments, but, to our knowledge, no program dedicated to molecule detection
in classical novae has ever been conducted with the array. In the ALMA
archive, we find only two classical novae observed so far. Observations
towards T Pyx (L. Schmidtobreick) were obtained with the ALMA Compact Array
(ACA), that is, with a smaller subset of the ALMA array composed of 7 m
antennas. The observations were made in bands 4, 6, and 7, which covered
several transitions commonly observed in circumstellar media, but our
inspection of the available data show that no line was detected. For instance,
near the SiO(3–2) line at 130.3 GHz, we get an rms of 4.0 mJy per beam of
12$\aas@@fstack{\prime\prime}$3$\times$7$\aas@@fstack{\prime\prime}$1 and per
30 km s-1. This is comparable to the sensitivity of our experiments at the SMA
(given the smaller SMA beams).

The main array of ALMA observed V5668 Sgr (Nova Sgr 2015 b) in 2017 (PI: M.
Diaz) in band 6. Emission in continuum and in H30 $\alpha$ near 231.9 GHz was
reported by Diaz et al. (2018). Although their data cover CO(2–1), no
molecular features were found. In a spectrum extracted for the entire
continuum source of a diameter of about 0$\aas@@fstack{\prime\prime}$5, the
3$\sigma$ upper limit on the CO emission is 0.54 mJy. This is possibly the
most stringent upper limit on CO emission in a classical nova to date.

The experiments to detect molecules in classical novae with mm-wave
interferometers conducted for young and old novae have given negative results
so far, but the number of observed sources is still small. Future similar
experiments are encouraged.

##  5 Summary and conclusions

Our single-dish survey of the 94 old classical novae and several dwarf novae
demonstrates that circumstellar envelopes of post-novae systems are poor in
molecular gas. Radiation of the classical novae outburst and the interstellar
radiation field must effectively destroy most molecules. Dust shielding, even
if provided shortly after the eruption, must be insufficient in the later
expansion phases of the nova ejecta.

With new maps of CO emission near GK Per, we debunked the long-standing
conviction that the remnant of Nova 1901 is associated with a circumstellar
molecular cloud. We find only interstellar emission in the direction of the
nova. It is possible that the dwarf nova flashes of GK Per destroy molecules
in the local ISM, reinforcing our conclusion of a deadly effect of novae on
molecular material.

The lack of molecular matter in classical and dwarf novae remnants show also
the remarkable difference with the newly recognized group of eruptive
variables of red novae, which are rich in molecules decades and centuries
after their eruptions. While red nova outbursts are often practically
indistinguishable from classical novae, the amount of molecular material
produced by these two classes of objects is drastically different and may be
used for new identifications of red nova in the future.

The upper limits we derive on molecular emission in classical novae are
relatively shallow and do not indicate that nova shells are completely devoid
of molecular matter. There is a chance that sensitive interferometric
observations will be able to detect this emission, especially given the
superior sensitivity of ALMA. A few experiments with the SMA and ALMA arrays
have resulted in negative results so far, but the sample of novae at different
times since the last eruption should be enlarged, especially for novae
actively forming dust. There is still hope for measuring accurate isotopic
ratios in classical novae using mm-wave data.

###### Acknowledgements.

We thank the referee, Nye Evans, for constructive comments on the manuscript.
T.K. and H.M. acknowledge funding from grant no 2018/30/E/ST9/00398 from the
Polish National Science Center. R.T. acknowledge a support from grant
2017/27/B/ST9/01128 financed by the Polish National Science Center. This
research has made use of the SIMBAD database, operated at CDS, Strasbourg,
France. This research made use of hips2fits, https://alasky.u-strasbg.fr/hips-
image-services/hips2fits a service provided by CDS. This paper makes use of
ALMA data ADS/JAO.ALMA#2019.2.00172.S and ADS/JAO.ALMA#2016.1.00682.S. ALMA is
a partnership of ESO (representing its member states), NSF (USA) and NINS
(Japan), together with NRC (Canada), MOST and ASIAA (Taiwan), and KASI
(Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA
Observatory is operated by ESO, AUI/NRAO and NAOJ. The National Radio
Astronomy Observatory is a facility of the National Science Foundation
operated under cooperative agreement by Associated Universities, Inc. Based on
observations with the APEX telescope under programme ID 094.F-9503(A). APEX is
a collaboration between the Max-Planck-Institut für Radioastronomie, the
European Southern Observatory, and the Onsala Observatory. Part of this work
is based on observations carried out with the IRAM 30m telescope. IRAM is
supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).

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##  Appendix A Sample APEX and IRAM spectra

Figure 7: Sample spectra covering the CO(3–2) transition at a resolution of 30
km s-1. The top left plot shows a reference spectrum of IRC+10216 with its
most prominent lines labelled. The velocity scale is given with respect to
CO(3–2). The green dashed lines mark 3$\sigma$ noise levels.

Figure 8: Same as Fig. 2 but for the CO(3–2) OSB setup. The velocity scale is
with respect to 333.795769 GHz.

Figure 9: Same as Fig. 2 but for spectra covering the CO(4–3) transition.
Velocities are with respect to the rest frequency of CO(4–3).

Figure 10: Same as Fig. 2 but for the CO(4–3) OSB setup. The velocity scale is
given relative to 473.041715 GHz.

Figure 11: Sample IRAM/FTS spectra covering the HCN(1–0) transition and at a
resolution of 30 km s-1. The top left plot shows a reference spectrum of CK
Vul with its most prominent lines labelled. The top axes give LSR velocity
with respect to the rest frequency of HCN(1–0).

Figure 12: Same as Fig. 2 but for the upper side band. The top axes give LSR
velocity with respect to 101.130514 GHz.

Figure 13: Spectra from IRAM/FTS lower side band covering the CO(1–0)
transition and at a velocity resolution of 30 km s-1. The velocity scale is
given relative to the rest frequency of CO(1–0).

Figure 14: Same as Fig. 2 but for the lower side band of the same setup. The
velocity scale is relative to 96.4125766 GHz.