# ASAS J071404+7004.3 – a close, bright nova-like cataclysmic variable with
gusty winds

K. Inight,1 B. T. Gänsicke,1 D. Blondel,1 D. Boyd${}^{2},$ R. P. Ashley,3 C.
Knigge,6 K. S. Long,7 T. R. Marsh,1 J. McCleery,1 S. Scaringi,4 D. Steeghs,1
J. R. Thorstensen,8 T. Vanmunster,5 P. J. Wheatley1
1Department of Physics, University of Warwick, Coventry, CV4 7AL, UK
2BAA Variable Star Section, West Challow Observatory, OX12 9TX, UKK
3Isaac Newton Group of Telescopes, Apartado de Correos 321, E-38700 Santa Cruz
de La Palma, Spain
4Centre for Extragalactic Astronomy, Department of Physics, Durham University,
DH1 3LE, UK
5CBA Belgium Observatory & CBA Extremadura Observatory, Walhostraat 1a, B-3401
Landen, Belgium
6School of Physics and Astronomy, University of Southampton, Highfield,
Southampton, SO17 1BJ, UK
7Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD,
21218, USA
8Department of Physics and Astronomy, Dartmouth College, Hanover NH 03755, USA
E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

Despite being bright ($V\simeq 11.8$) and nearby ($d=212$ pc) ASAS
J071404+7004.3 has only recently been identified as a nova-like cataclysmic
variable. We present time-resolved optical spectroscopy obtained at the Isaac
Newton and the Hiltner and McGraw-Hill Telescopes, together with Swift X-ray
and ultraviolet observations. We combined these with TESS photometry and find
a period of 3.28 h and a mass transfer rate of $4-9\times
10^{-9}\,\mbox{$\mathrm{M_{\sun}\,yr}^{-1}$}$. Historical photometry shows at
least one low state establishing the system as a VY Scl star. Our high-cadence
spectroscopy also revealed rapidly changing winds emanating from the accretion
disc. We have modelled these using the Monte Carlo python code and shown that
all the emission lines could emanate from the wind – which would explain the
lack of double-peaked lines in such systems. In passing, we discuss the effect
of variability on the position of cataclysmic variables in the Gaia
Hertzsprung-Russell diagram.

###### keywords:

Hertzsprung-Russell and colour-magnitude diagrams – cataclysmic variables –
stars:evolution

††pubyear: 2020††pagerange: ASAS J071404+7004.3 – a close, bright nova-like
cataclysmic variable with gusty winds–A

##  1 Introduction

Cataclysmic variables (CVs) are short-period ($\simeq 1-10\mathrm{h}$) binary
systems consisting of a white dwarf accreting from a main sequence companion
(Warner, 1995). In CVs with non-magnetic (or very weakly magnetic) white
dwarfs, the mass transfer proceeds via an accretion disc. Below a critical
mass transfer rate ($\dot{M}\simeq 1-5\times 10^{-9}$
$\mathrm{M_{\sun}\,yr}^{-1}$) thermal instabilities in accretion discs result
in dwarf novae outbursts, in which the disc cycles between a cool and a hot
state. CVs above the critical $\dot{M}$ appear as nova-likes which have
steady, hot and luminous discs (Meyer & Meyer-Hofmeister, 1984; Shafter et
al., 1986; Cannizzo et al., 1988). In reality, CVs display a range of
behaviour that goes beyond the simple dichotomy described above. Some CVs have
mass transfer rates close to the critical value, and cycle between undergoing
regular outbursts and states where their discs are in a steady hot state
(Buat-Ménard et al., 2001). Even though the accretion discs in nova-likes are
in a (relatively) steady luminous state, they can exhibit small “stunted”
outbursts (Honeycutt et al., 1998). In addition, some nova-likes occasionally
enter a low state during which mass loss from the donor drops significantly,
or even shuts off completely (Honeycutt & Kafka, 2004). During this time they
may exhibit dwarf nova outbursts (King & Cannizzo, 1998).

In summary, CVs provide an ideal environment for the study of accretion
processes under a range of physical conditions, with applications to a range
of other systems including young stellar objects (Lada, 1987) and quasars
(Kembhavi & Narlikar, 1999). Key areas where our understanding of CV accretion
discs is incomplete are the nature of the viscosity driving the mass transfer
through the disc (Shakura & Sunyaev, 1973), the detailed physical structure of
the disc (Godon et al., 2017) and of the interface between the disc and the
white dwarf (the so-called “boundary layer”, see Section 2.5.4 in Warner
2003). In addition, CVs often display signatures of winds (Mauche & Raymond,
1997), and in some cases jets (Coppejans & Knigge, 2020), and both the
structure and the driving mechanisms of these are not very well established.

Discovering new, bright examples of CVs remains important to grow the arsenal
of diagnostic tools that are available to investigate these topics. However,
because of their steady discs, nova-like variables have remained difficult to
identify, in contrast with dwarf novae that stand out in photometric time-
domain surveys because of their large-amplitude variability (Gänsicke, 2005;
Drake et al., 2014; Breedt et al., 2014; Szkody et al., 2020).

Here we report a detailed study of the bright and nearby nova-like ASAS
J071404+7004.3 (hereafter ASAS J0714+7004). ASAS J0714+7004 was identified as
the optical counterpart to the X-ray source 1RXS J071404.0+700413, and a
potential variable star, by Kiraga & Stępień (2013) who suggested it might be
a CV with a putative period of 22.4 d. ASAS J0714+7004 came to our attention
as part of a project to identify white dwarf binary candidates located between
the main sequence and the white dwarf cooling sequence, using a combination of
Gaia and GALEX photometry and astrometry (Inight et al. in prep).

We summarise our observations and provide a qualitative description of the
system in Section 2, analyse these observations to extract the physical
parameters of the system in Section 3 and then discuss the implications
(Section 4) and conclusions (Section 5).

Figure 1:  Spectra of ASAS J071404+7004.3. Top panel: The first spectrum (red
– offset by 0.6 flux units) was obtained on 2020 February 6. The second
spectrum (blue) is the average of the INT observations. The third spectrum
(green – offset by 1.0 flux units) is the average of the MDM observations.
Bottom panel: The average INT spectra for H$\beta$ (left), He i 5876 Å
(centre) and H$\alpha$ (right). The H$\beta$ line shows an emission line
sitting in a broad absorption line, the He i 5876 Å triplet line shows a P
Cygni profile and the H$\alpha$ line is purely in emission although showing
slight traces of a P Cygni profile.

##  2 Observations

###  2.1 Spectroscopy

We obtained an identification spectrum on 2020 February 6 using a 280 mm
Schmidt Cassegrain telescope (SCT) equipped with a Shelyak LISA spectroscope,
giving a spectral resolution of $R\simeq 1000$. The spectrum was integrated
for 30 min before being stopped by clouds which limited the signal-to-noise
ratio. The data were corrected with bias, dark and flat frames, wavelength
calibrated using the integrated ArNe lamp, and corrected for the instrumental
response function and atmospheric losses using the spectrum of a
spectrophotometric standard star taken at the same altitude immediately prior
to the target. Processing was performed using the isis spectral analysis
software111http://www.astrosurf.com/buil/isis-software.html. The spectrum was
then calibrated in absolute flux using a concurrently recorded $V$ magnitude
observation using the method described on the BAA website222
https://britastro.org/sites/default/files/absfluxcalibration.pdf. The SCT
spectrum (Fig. 1) exhibits a blue slope in the continuum with broad and
shallow H$\beta$ and H$\gamma$ absorption lines filled in by single-peaked
Balmer emission lines, and additional emission lines from He i, unambiguously
identifying this system as a CV. With a median magnitude of $V\simeq 11.8$,
ASAS J0714+7004 is among the $\simeq 20$ brightest CVs known (Ritter & Kolb,
2003).

ASAS J0714+7004 was then observed using the Intermediate Dispersion
Spectrograph (IDS) mounted on the 2.54 m Isaac Newton Telescope (INT) on the
island of La Palma, using the R632V grating at a central wavelength of 5720 Å.
This setup covers the wavelength range 4383–6704 Å at a dispersion of 0.89
Å/pixel. A total of 264 120 s exposures were taken over a period of 13 days in
February 2020 (hereafter the February INT data) and a further long sequence on
2020 December 13 (hereafter the December INT data) – see Table 3 for details.

All of the INT spectra were optimally extracted and reduced using the pamela
and molly reduction software (Marsh, 1989). In order to perform the flux
calibration we fitted a spline to the continua of spectrophotometric standard
stars taken on the same night as the observations. We then used their
published flux as a reference, see e.g. Marsh (1990). For the wavelength
calibration, Cu+Ne and Cu+Ar arcs were observed at the start and end of each
observation, and approximately once every hour in between during the longer
sequences. As a final step in wavelength calibration we extracted the sky
spectrum and measured the deviation of the telluric O i emission line at
5577.4 Å from its laboratory wavelength and shifted our science spectra to
align with this feature.

Additional spectra of ASAS J0714+7004 were obtained prior to the SCT and ING
observations333Upon posting the submitted version of this paper to arXiv, John
Thorstensen reminded the other authors that he had observed this star already
in 2013. with the 1.3 m McGraw-Hill and 2.4 m Hiltner telescopes at the MDM
Observatory on Kitt Peak in Arizona; Table 2 gives a journal. These data will
be referred to as the MDM spectra hereafter. On both telescopes we used the
modular spectrograph (modspec), which covered from 4210 to 7560 Å with a
resolution of about 4 Å (full-width at half maximum). To calibrate the
wavelengths, we took spectra of Hg, Ne, and Xe lamps in twilight, and derived
zero-point shifts to the solution using the night-sky background lines. We
flux-calibrated the spectra using observations of standard stars; experience
suggests that the zero point of these calibrations is accurate to $\simeq 30$
per cent. In addition, the continua in modspec data sometimes show unphysical
distortions, which tend to average out.

The average of the 264 INT spectra (Fig. 1) is very similar to the SCT
identification spectrum and the MDM average spectrum, although it has a higher
spectral resolution and signal-to-noise ratio. It contains strong single-
peaked Balmer emission lines embedded within broad and shallow absorption
lines. Overall, the morphology of the spectrum is that of a nova-like, with
the absorption arising from an optically thick disc (Dhillon, 1996; Haug &
Drechsel, 1985; Rosino et al., 1993). Unlike dwarf novae, single-peaked Balmer
lines are not indicative of low inclination (Thorstensen et al., 1991;
Rodríguez-Gil et al., 2007; Neustroev et al., 2011). Strong emission lines in
nova-likes are in fact associated with high inclination (see e.g. Fig. 1 in
Matthews et al. 2015). This is because they are believed to originate in the
disc winds and hence appear strongest relative to the underlying disc emission
at higher inclinations. The He i 5876 Å line exhibits a P Cygni profile.

Table 1:  Stellar parameters of ASAS J0714+7004.

Parameter | Value | References
---|---|---
RA |  $\alpha$ [h:m:s] | 07:14:04.654486$\,\pm\,$ | 0.00001 | 1
Dec |  $\delta$ [d:m:s] | +70:04:18.400755$\,\pm\,$ | 0.00001 | 1
Gaia ERD3 source_id |  | 1109608206832496512 | 1
Parallax |  $\varpi$ [mas] | 4.703$\,\pm\,$ | 0.02 | 1
Distance |  $d$ [pc] | 210.15$\,\pm\,$ | 1.0 | 2
Apparent magnitude |  $G$ [mag] | 11.814$\,\pm\,$ | 0.01 | 1
Absolute magnitude |  $M_{G}$ [mag] | 5.20$\,\pm\,$ | 0.016 | 1,2,4
Proper motion |  $\mu_{\alpha}$ $\mathrm{[mas\,yr^{-1}]}$ | 26.007$\,\pm\,$ | 0.02 | 1
|  $\mu_{\delta}$ $\mathrm{[mas\,yr^{-1}]}$ |  $-$36.923$\,\pm\,$ | 0.025 | 1
Mass transfer rate |  $\dot{M}$ [$10^{-9}\mbox{$\mathrm{M_{\sun}\,yr}^{-1}$}$] | 4-9 | 3
Period |  $P_{\textrm{orb}}$ [h] | 3.2794894$\,\pm\,$ | 0.0000008 | 3
Epoch | HJD [d] | 2456597.8330$\,\pm\,$ | 0.0005 |
Inclination |  $i$ [deg] | 50-70 | 3

1 Gaia Collaboration et al. (2021); 2 Bailer-Jones et al. (2021); 3 this
paper; 4 Not taking account of extinction

###  2.2 Photometry

The long-term light curve of ASAS J0714+7004 obtained by ASAS-SN (Shappee et
al., 2014) displays quasi-periodic brightness variations on time scales of
$\simeq 20$ d with an amplitude of $\simeq 0.5$ mag (panels (a), (b) and (c)
in Fig. 2). In order to investigate the nature of these variations, we began
intensive photometric monitoring of the system using small-aperture telescopes
throughout February and March 2020. However, during this campaign, ASAS
J0714+7004 did not display much variability, remaining at an average magnitude
of $V\simeq 11.85\pm 0.12$ mag, although the quasi-periodic brightening
resumed again in June 2020 (panel (d) in Fig. 2).

A peculiar event captured by the ASAS-SN data was a drop in the minimum
brightness in between the quasi-periodic brightening in April 2015, whereas
the maximum brightness remained constant (panel (b) in Fig. 2). This “low
state” lasted at least two months; unfortunately the ASAS-SN observations were
stopped due to ASAS J0714+7004 moving into day time. When the observations
resumed in August 2015, the system had returned to the normal state.

The TESS satellite (Ricker et al., 2015) observed ASAS J0714+7004 in short-
cadence (2 min exposures) during Sector 20 (2019 December 24 to 2020 January
21) and Sector 26 (2020 June 8 to July 4) and in ultra-short cadence (20 s
exposures) during Sector 40 (2021 June 24 to July 23). The space-based
photometry covers part of the inactive phase (panels (e) and (f) in Fig. 2),
the first brightening following this inactive phase (panels (g) and (h)), and
two further brightenings (Fig. 2, panel (d)). All the TESS data sets reveal
substantial periodic short-term variability.

Figure 2: Ground and space-based photometry of ASAS J0714+7004. The horizontal
bars above the top panel, and within panels (e) and (g), indicate the date
range of the zoom-in views. Note the change in the scale of the magnitude axis
in the different panels. The system exhibits quasi-periodic brightenings most
of the time, but experienced a fainter state (a,b) as well as a phase of
relative inactivity (a,e,f). The high-cadence TESS data reveals short-term
flickering (e,f) and low-amplitude orbital modulations at the orbital period
(g,h).

###  2.3 Swift

ASAS J0714+7004 was observed in January 2021 using the Neil Gehrels Swift
Observatory (Evans et al., 2009). Data were obtained from both the ultraviolet
optical telescope (UVOT) (Roming et al., 2005) and X-ray telescope (XRT)
(Burrows et al., 2005). Ten observations were obtained over a 22 d period with
exposure times $\simeq 10$ min, see Table 4. Six of the ultraviolet
observations used the $U$ filter (mean wavelength 3492 Å and width 866 Å) and
four used the $M2$ filter (mean wavelength 2272.21Å and width 533.85 Å). The
XRT covers the range $0.2-10$ keV and has two modes – Windowed Timing (WT)
with continuous readout (and only one-dimensional imaging) and Photon Counting
(PC) where the exposure length is 2.5 s. We have only used PC data. The XRT
data were processed using the automated pipeline to produce an average
spectrum and a light curve. The UVOT data were saturated in both $U$ and $M2$
filters and therefore a modified approach (Page et al., 2013) was used to
obtain photometry by measuring the “streak” as the CCD is read out.

Table 2:  Journal of MDM observations.

UT (start) | Tel. | HA start | HA end | Exp(s) | $N_{\rm exp}$
---|---|---|---|---|---
2013-11-01 09:39 | M |  $-$02:19 | +00:55 | 600 | 3
2014-01-22 11:05 | H | +04:30 | +05:04 | 480 | 4
2014-01-26 11:06 | H | +04:47 | +05:05 | 240 | 4
2014-01-27 06:22 | H | +00:06 | +00:50 | 480 | 4
2014-10-12 12:37 | H |  $-$00:40 |  $-$00:32 | 480 | 1
2014-10-13 12:49 | H |  $-$00:24 |  $-$00:19 | 300 | 1
2014-10-14 12:55 | H |  $-$00:15 |  $-$00:04 | 300 | 2
2014-10-15 12:10 | H |  $-$00:56 |  $-$00:06 | 300 | 9
2014-10-16 07:17 | H |  $-$05:46 |  $-$00:03 | 240 | 17

Telescopes are: M - McGraw-Hill 1.3 m; H - Hiltner 2.4 m. Hour angles are in
hours and minutes.

Table 3:  The journal of INT observations for the February and December proposals.  Date | JD (start) | N (observations) | Duration
---|---|---|---
|  |  | (minutes)
2020/02/08 | 2458887.583423 | 52 | 117.61
2020/02/09 | 2458888.530952 | 52 | 124.56
2020/02/10 | 2458889.465164 | 54 | 122.25
2020/02/12 | 2458891.600877 | 2 | 4.20
2020/02/13 | 2458892.455516 | 49 | 113.99
2020/02/13 | 2458892.543982 | 44 | 117.47
2020/02/16 | 2458895.603163 | 1 | 2.00
2020/02/17 | 2458896.607681 | 2 | 4.20
2020/02/18 | 2458897.602185 | 2 | 4.20
2020/02/19 | 2458898.575209 | 2 | 4.20
2020/02/20 | 2458899.526425 | 2 | 4.20
2020/02/20 | 2458899.581305 | 2 | 4.20
2020/12/13 | 2459197.430520 | 195 | 522.0
Table 4:  The journal of Swift observations. The central wavelengths of the $U$-band and $M2$-filter are 3492.67 Å and 2272.71 Å, respectively. Date (UTC) | UVOT | XRT PC | UV | UV Mag
---|---|---|---|---
| exposure (s) | exposure (s) | Filter | (Vega)
2020-12-29 23:28:31 | 76.2 | 70.0 | $U$ | 10.947(129)
2021-01-02 17:56:34 | 702.8 | 699.2 | $U$ | 11.040(45
2021-01-04 00:13:36 | 583.5 | 574.7 | $M2$ | 10.510(39)
2021-01-06 08:10:35 | 557.9 | 552.3 | $U$ | 10.785(43)
2021-01-08 09:25:38 | 374.1 | 684.5 | $M2$ | 9.828(32)
| 316.2 |  | $M2$ | 9.680(32)
2021-01-10 11:08:21 | 412.8 | 402.8 | $U$ | 10.409(38)
2021-01-14 08:56:35 | 425.2 | 423.7 | $U$ |
2021-01-16 07:05:34 | 747.3 | 782.5 | $M2$ | 10.627(38)
2021-01-18 08:32:35 | 564.5 | 557.5 | $U$ | 10.897(44)
2021-01-20 06:27:34 | 544.8 | 540.0 | $M2$ | 10.430(38)
Figure 3:  Top panel: Power spectra of the time-resolved radial velocity
variations measured from the $H_{\alpha}$ line (offset in green) and the
Sector 20, 26 and 40 TESS light curves (blue). There is a clear peak at $\sim
3.28$ h. Central panel: Detail from the top plots. Bottom panel: The radial
velocities of the INT data (blue) and the MDM data (brown) phase-folded using
the adopted period of 3.2794894 h, superimposed by a sine-fit to the data.

##  3 Analysis

The morphology of the spectroscopic and photometric data argues against a
dwarf nova nature of ASAS J0714+7004. Although, at a first glance, the quasi-
regular brightenings may seem unusual (see the discussion in Sect. 4.4), we
conclude that this system is best classified as a nova-like variable.

###  3.1 Orbital period

The primary motivation for obtaining the fast cadence INT observations was to
measure the evolution of radial velocities, which we carried out using the
molly software package (Marsh, 2019). We also used molly to process the MDM
spectra. We normalised the continuum of the spectra using a straight line fit,
and then determined the central wavelength of the H$\alpha$ emission line by
convolving each spectrum with a Gaussian template with a standard deviation of
$\sigma=300\,\mbox{$\mathrm{km\,s}^{-1}$}$. A Fourier analysis (Deeming, 1975)
of the resulting INT and MDM radial velocity time-series was then used to
compute a periodogram (Fig. 3 green line in the top panel). The highest peak
in the periodogram identifies the orbital period as $P_{\mathrm{orb}}\simeq
3.28\,$h, albeit the periodogram is plagued by a substantial alias structure
due to the sparse sampling of the radial velocities (Fig. 3 middle panel).

The Fourier transform of the Sector 20,26 and 40 TESS photometry results in a
much less aliased periodogram (blue in Fig. 3) which, however, also contains
some alias structure resulting from the gaps in the middle of the TESS
observations (panel (d) in Fig. 2). We note that our time-series analysis of
the TESS data did not reveal any other periodic signal such as that seen in
the nova-like V341 Ara (which was interpreted as a beat frequency between the
orbital period and the retrograde precession of the accretion disc in that
system (Fig. 8 in Castro Segura et al. 2021).

The combined radial velocity and photometric periodograms overlap only for one
unique period, which we therefore identify as the true orbital period of ASAS
J0714+7004 (middle panel in Fig. 3). In order to improve the precision of the
orbital period, we performed a sinusoidal fit to the INT and MDM radial
velocity data, which spans a much longer temporal baseline, with five distinct
epochs, than the TESS photometry, and hence provides a more accurate value of
the period. We found $P=3.2794894(8)$ h and an epoch of
$\mathrm{HJD(UTC)}=2456597.8330(5)$. We define the epoch as being the blue to
red crossing of the radial velocity. The velocity amplitude of the fit is
$143.5(2.2)\,\mbox{$\mathrm{km\,s}^{-1}$}$ The phase-folded data are shown in
(Fig. 3 bottom panel). A trailed spectrum of the December INT data is shown in
Fig. 10.

The integrated flux from the H$\alpha$ line measured from the December INT
data appears to lag the radial velocity variation by $\simeq 90\degr$ (see
Fig. 4). Assuming that the H$\alpha$ emission originates in the disc, then the
radial velocity variation will track the movement of the disc (and hence the
white dwarf at its centre) around the centre of gravity of the system. The
peak of the integrated flux therefore occurs at the superior conjunction when
the white dwarf is at its furthest point and the donor is at its closest.
Thus, the relative phasing of the integrated line fluxes and radial velocities
argues against the emission originating on the irradiated donor star.

Figure 4:  The radial velocity of the H$\alpha$ core (bottom panel) and the
integrated flux from the H$\alpha$ line (centre panel) for the December INT
data together with the TESS photometry binned and folded at the orbital period
and projected onto the December timeline. Each dataset was fitted with a
sinusoid at the orbital period. The integrated flux is offset by 0.241 cycles
from the radial velocities whilst the Tess photometry is offset by 0.5 cycles.

###  3.2 The mass transfer rate

A key parameter of CVs is their mass transfer rate, $\dot{M}$. This parameter
is difficult to determine in dwarf novae, where the disc effectively buffers
the mass lost from the companion (Hessman et al., 2000), and cyclically
flushes it onto the white dwarf. However nova-likes such as ASAS J0714+7004
can be in a quasi-steady-state, facilitating the measurement of $\dot{M}$.
Here, we estimate $\dot{M}$ in ASAS J0714+7004 by modelling its spectral
energy distribution (SED).

We have assembled the overall SED of ASAS J0714+7004 making use of available
X-ray to infrared photometry (Pan-STARRS, Chambers et al. 2016; 2MASS, Cutri
et al. 2003; WISE, Cutri & et al. 2014; Gaia, Gaia Collaboration 2020; ATLAS,
Tonry et al. 2018; APASS, Henden et al. 2015 and GALEX, Morrissey et al. 2007)
and of our Swift observations (Sect. 2.3).

The SED of ASAS J0714+7004 (Fig. 5) presents a steep blue slope from the
infrared into the ultraviolet, with a significant drop in the spectral energy
density in the X-ray regime. The system is sufficiently bright that the GALEX
observations suffer from saturation effects (Morrissey et al., 2007) – the
Swift UVOT photometry confirms that the GALEX photometry is not reliable, and
therefore it was not included in the fit.

The optical magnitude of ASAS J0714+7004 is $\simeq 12$ mag and this can also
be expected to saturate the Pan-STARRS data (saturation occurs at $12-14$ mag
dependent upon seeing Magnier et al. 2013) and the ATLAS data (Tonry et al.,
2021). The saturated points were therefore ignored when obtaining the best fit
and are not shown (except for GALEX) in Fig. 5.

Figure 5: Multi-wavelength observations of ASAS J0714+7004 (see text for
details). Top panel: Spectral energy density as a function of wavelength. The
disc dominates the SED in the optical; dashed lines show the contribution of
the other components. The boundary layer contribution is only notionally
constrained by the observations. The black solid line shows the optimum fit
whilst the magenta line plots a fit constrained by the estimate of inclination
from Sect. 3.3. It is noticeable that the fits consistently undershoot the
infrared observations. Such an infrared excess has been observed in other
nova-likes (e.g. Hoard et al., 2014), indicating that the adopted model is not
accounting for all components in the system. The excess may be due to a dusty
circumstellar disc (Hoard et al., 2007; Hoard et al., 2009). Bottom panel: A
heat map of the chi-squared of our model fit to the observed SED as a function
of the two free parameters, $i$ and $\dot{M}$. The green dotted line traces
the degeneracy between the orbital inclination and the mass transfer rate,
illustrating that the two parameters cannot be determined independently. The
black line indicates the best fit to the SED, and the magenta line shows the
estimate from Sect. 3.3. The solid purple line traces the $1\sigma$ contour.

In order to model the SED of ASAS J0714+7004, we have to account for the flux
emitted by the individual components of the system, i.e. the disc, the white
dwarf, the donor star and potentially also a hot boundary layer close to the
white dwarf where up to half of the gravitational potential energy is
released. The detailed structure, and even the presence of a boundary layer in
CVs has been subject to decades of discussion (Ferland et al., 1982; Long et
al., 1996; Hertfelder et al., 2013; Nixon & Pringle, 2019), and, on the
observational side, is severely limited by the fact that interstellar
absorption severely hinders access to the extreme ultraviolet (EUV) emission
of CVs444The boundary layer is expected to have temperatures of a few $10^{5}$
K, and hence its emission will peak in the Extreme Ultraviolet. This
wavelength region is heavily absorbed by interstellar neutral hydrogen, and
hence only a handful of nearby CVs have adequate observations, which led to
conflicting results: Long et al. (1996) observed U Gem during outburst and
concluded that the peak spectrum resembled a blackbody with
$\mbox{$T_{\mathrm{eff}}$}\simeq 140\,000$ K. Conversely a joint ROSAT and
EUVE data analysis of the nova-like IX Vel showed no blackbody emission (van
Teeseling et al., 1995).. In the case of ASAS J0714+7004, we note that whereas
a putative boundary layer may contribute in the EUV, it will make only a
minimal contribution when fitting the ultraviolet to infrared emission of the
system (see Fig. 5). Here, we model the SED of ASAS J0714+7004 by the sum of
the disc, the white dwarf and donor and the boundary layer. In nova-likes, the
accretion disc is expected to be the primary source of optical and ultraviolet
flux, in contrast the two stellar components have only relatively small
contributions.

We model the disc as a set of 1000 concentric annuli of constant width between
the inner and outer radii of the disc, $R_{\mathrm{disc,i}}$ and
$R_{\mathrm{disc,o}}$, respectively. To determine the temperature $T(r)$ of
each annulus we assume that the flux per unit surface area at radius $r$ in
the disc is given by (see chapter 5 of Frank et al. 1992):

$F=\frac{3GM_{1}\dot{M}}{8\pi
r^{3}}\left(1-\left(\frac{R_{1}}{r}\right)^{1/2}\right)$ (1)

where $M_{1}$ and $R_{1}$ are the mass and radius, respectively, of the white
dwarf and $\dot{M}$ is the rate of mass transfer. Given that the integrated
flux is a function of the temperature,

$F=\sigma T^{4}$ (2)

we obtain

$T=\left(\frac{3GM_{1}\dot{M}}{8\pi\sigma
r^{3}}\left(1-\left(\frac{R_{1}}{r}\right)^{1/2}\right)\right)^{1/4}$ (3)

We then used the BT-Settl (AGSS2009) spectral models (Allard et al., 2012) to
represent the spectrum for each annulus based on the temperature from Eq. 3,
$\log g=4.0$ and solar metallicity. The shortcomings of using synthetic
stellar spectra to model the emission of an accretion disc have been discussed
at length, (Wade, 1988) however they are better than the alternative which is
to use a blackbody spectrum.

The white dwarf temperature was treated as a function of $\dot{M}$ using Eq. 2
in (Townsley & Gänsicke, 2009):

$T_{1}=17\,000\mathrm{[K]}\left(\frac{\langle\dot{M}\,[\mbox{$\mathrm{M_{\sun}\,yr}^{-1}$}]\rangle}{10^{-10}}\right)^{\frac{1}{4}}\left(\frac{M_{1}\,[\mbox{$\mathrm{M}_{\sun}$}]}{0.9}\right)$
(4)

For the donor star, we adopted a mass, radius, and effective temperature based
upon the evolutionary track from Table 2 by Knigge et al. (2011). The emission
of the white dwarf and the donor were modelled by two blackbodies of the
corresponding temperatures and radii.

The fit shown in Fig. 5 was obtained by fixing
$M_{1}=0.83\,\mbox{$\mathrm{M}_{\sun}$}$ (the average value for CVs from
Zorotovic et al. 2011) and $R_{1}=0.0122\,\mbox{$\mathrm{R}_{\sun}$}$ (using
the mass/radius relationships from Bédard et al. 2020 available at
555https://www.astro.umontreal.ca/~bergeron/CoolingModels/).

We also followed Godon et al. (2017) and assumed that the inner edge of the
accretion disc is located one $R_{1}$ above the white dwarf surface. We
followed Harrop-Allin & Warner (1996) and assumed that the outer radius of the
disc (due to tidal truncation) is $0.7\times R_{\textrm{L1}}$ where
$R_{\textrm{L1}}$ is the distance from the centre of the white dwarf to the
inner Lagrange point given by:

$R_{\textrm{L1}}=a\left(1.0015+q^{0.4056}\right)^{-1}$ (5)

The inclination of the accretion disc and ${\dot{M}}$ were treated as free
parameters.

We assumed extinction based on the galactic formulae in (Howarth, 1983). We
then fitted the data using the system parameters in table 5 and a $320\times
600$ linearly spaced grid of values of $\dot{M}$ from $0.5-20\times
10^{-9}\,\mbox{$\mathrm{M_{\sun}\,yr}^{-1}$}$ , and inclination $10-85\degr$
to produce the heat map in Fig. 5. In order to account for the fact that the
object is intrinsically slightly variable, we adopted for the fit either the
published uncertainties on the photometric data, or a value of 10 per cent,
whichever was larger.

To demonstrate the contribution of a potential boundary layer we also modelled
the effect of adding a circular band with $H=0.1R_{1}$, $R=R_{1}$ and
$T_{\textrm{bl,eff}}=100\,000$ K. This does not significantly affect the fit
as its emission in the optical and near-ultraviolet is far less than that of
the disc (see Fig. 5).

Table 5:  System parameters adopted for fitting the spectral energy distribution.  Parameter | Value
---|---
Mass (white dwarf) |  $M_{1}$ [$\mathrm{M}_{\sun}$] | 0.83
Mass (donor) |  $M_{2}$ [$\mathrm{M}_{\sun}$] | 0.22
Separation |  $a$ [$\mathrm{R}_{\sun}$] | 1.11
Radius (white dwarf) |  $R_{1}$ [$\mathrm{R}_{\sun}$] | 0.0122
Radius (donor) |  $R_{2}$ [$\mathrm{R}_{\sun}$] | 0.30
Inner radius (disc) |  $R_{\mathrm{disc,i}}$ [$\mathrm{R}_{\sun}$] | 0.0244
Outer radius (disc) |  $R_{\mathrm{disc,o}}$ [$\mathrm{R}_{\sun}$] | 0.49
Temperature (donor) |  $T_{\mathrm{2}}$ [K] | 3324

There is considerable degeneracy between the angle of inclination and
$\dot{M}$. This is evident from the dotted green line in Fig. 5 showing the
best fit for a given inclination. The overall best fit is found for
$i=52.9\degr$ (see table 6), consistent with the estimate calculated in
section 3.3 ($65.7\degr\begin{subarray}{c}+8\\\ -13\end{subarray}$ ). The disc
luminosity is better constrained, $3\pm 1\times 10^{33}\,\mathrm{erg\,s^{-1}}$
(see Table 6) .

The white dwarf temperature determined from the SED fit only mildly depends on
the inclination (Table 6) and is consistent with those found in other nova-
likes and VY Scl stars from direct modelling of the white dwarf spectrum: DW
UMa ($\simeq 50\,000$ K, Araujo-Betancor et al. 2003), TT Ari ($\simeq
39\,000$ K, Gänsicke et al. 1999) and MV Lyr ($\simeq 47\,000$ K, Hoard et al.
2004).

Table 6:  Estimated values from fitting the spectral energy distribution.  Parameter | i=52.9 $\deg$ | i=65.7 $\deg$
---|---|---
Mass transfer rate |  $\dot{M}$ [$10^{-9}\mbox{$\mathrm{M_{\sun}\,yr}^{-1}$}$] | 4.75 | 8.13
Temperature |  $T_{\mathrm{1}}$ [K] | 41 158 | 48777
(white dwarf) |  |  |
Disc luminosity |  $L_{\mathrm{disc}}$ [$\mathrm{10^{33}erg\,s^{-1}}$] | 2.59 | 3.66

###  3.3 Orbital inclination

Figure 6:  Top panel: The “Gold Sample” of known CVs within 300pc with
reliable astrometry and astrometry from Inight et al. (2021). ASAS J0714+7004
is marked with a blue square. Bottom panel: Detail showing the area of the HRD
occupied by a representative selection of well-studied nova-likes. Confirmed
SW Sex stars are shown using green dots and candidate SW Sex stars using
yellow. The data is from Gaia EDR3 (Gaia Collaboration et al., 2020) without
adjusting the absolute magnitudes for extinction. The inclination (where
known) is shown in brackets (sources ’SW Sex’, (Groot et al., 2001) ’BH Lyn’,
(Dhillon et al., 1992) ’DW UMa’, (Bíró, 2000) ’Lanning 386’, (Brady et al.,
2008) ’MV Lyr’, (Hoard et al., 2004) ’BG Tri’, (Hernández et al., 2021) ’RW
Tri’, (Smak, 2019) ’UX UMa’, (Suleimanov et al., 2004) ’TT Ari’, (Bruch, 2019)
’UU Aqr’, (Baptista et al., 1994) ’RW Sex’, (Beuermann et al., 1992) ’IW And’,
(Szkody et al., 2013) ’VY Scl’, (Martínez-Pais et al., 2000) ’Z Cam’, (Szkody
& Wade, 1981) ’V3885 Sgr’, (Hartley et al., 2002) ’IX Vel’, (Hartley et al.,
2002)). Lanning 386 is an interesting outlier. It is an eclipsing system
(Kennedy et al., 2017; Kennedy et al., 2016) and more significantly it is most
often, unusually for a nova-like, in its low state.

As discussed in Sect. 3.2, the inclination of ASAS J0714+7004 is not tightly
constrained from the SED fit. Here, we will follow an alternative approach to
estimate the inclination based on the absolute magnitude of the system. Its
location within the Hertzsprung-Russell diagram (HRD) places it firmly among
the nova-likes (Fig. 6), which are characterised by high mass-transfer rates
that keep their accretion discs in the hot state. We therefore adopt the
absolute magnitude-orbital period relation of Paczynski & Schwarzenberg-Czerny
(1980), see also (Warner 1987 )

$M_{\mathrm{}}=5.7\pm 0.13-0.287\pm 0.024\times P_{\mathrm{orb}}$ (6)

which is adjusted for inclination by:

$\Delta M_{\mathrm{v}}=-2.5\times\log\left(\left(1+\frac{3}{2}\cos
i\right)\cos i\right)\pm{0.4}$ (7)

Based on the Gaia $G$-band data in Table 1, the mean absolute $G$-band
magnitude of ASAS J0714+7004 is $5.20\pm{0.016}$. Using Eq. 6 and the orbital
period of 3.28 h gives a predicted absolute magnitude of
$4.76\begin{subarray}{c}+0.21\\\ -0.16\end{subarray}$ and therefore $\Delta
M_{\mathrm{V}}=0.44\pm 0.21$. Solving Eq. 7 for $i$ (and taking account of the
error in the formula of $\pm{0.4}$) yields an inclination of
$65.7\begin{subarray}{c}+8\\\ -13\end{subarray}\degr$. This estimate is also
consistent with the SED fit, the strong Balmer emission lines and the fact
that it is not eclipsing. It is also consistent with Figure 1 in Howell &
Mason (2018).

The Gaia astrometry provides the opportunity to test the absolute magnitude-
orbital period relation: most nova-likes, including ASAS J0714+7004, have
orbital periods of $3-4$ h, and hence their absolute magnitudes should
primarily depend on their inclinations. Inspecting Fig. 6 (lower panel), this
is indeed the case, where the fainter systems have higher inclinations – in
particular the green and yellow coded SW Sex systems are eclipsing and will
therefore inevitably have large inclinations. We caution, however, that some
care needs to be taken when considering cataclysmic variables in the Gaia HRD,
as their exact location is subject to the state(s) that the Gaia sampled. As
such, CVs may move by some amount within the HRD when comparing their DR2 and
EDR3 data (see Appendix A).

###  3.4 Photometric variability

The uninterrupted TESS observations provide the opportunity to investigate the
broad-band variability of ASAS J0714+7004 on timescales of minutes to days.
The Fourier transform (Deeming, 1975) computed from the Sector 40 data that
was obtained in the ultra-short 20 s cadence mode (Fig. 7) shows clearly the
3.28 h orbital period, and a flat distribution ($-1$ powerlaw) in power
density down to $\simeq 6.3$ mHz, after which the power law breaks to $-2$. A
weak signal at 2000 s was detected in the Sector 40 data, but was absent in
the other TESS observations, and may have been related to intermittent quasi-
periodic oscillations (Warner & Woudt, 2002). The power law slopes and break
frequency are typical of those observed in other accreting white dwarfs which
display strong broad-band variability components (see e.g. MV Lyr, Scaringi et
al. 2012; Scaringi 2014; Scaringi et al. 2015). This phenomenology is
generally explained by the so-called fluctuating accretion disc model, which
was originally developed to explain the broad-band noise properties observed
in X-ray binaries and active galactic nuclei (see Lyubarskii 1997; Arévalo &
Uttley 2006; Ingram & Done 2011, 2012; Ingram & van der Klis 2013). In this
model mass transfer variations through the disc couple multiplicatively as
matter is moved inwards through the disc. The observed power spectrum is then
the result of the faster variability driven by mass transfer fluctuations and
emitted in the inner-disc regions (Scaringi et al., 2015). In reality the
variability is generated across a range of disc radii, and is observed in a
specific pass-band given an emissivity profile for the disc. A similar power
spectral density observed in the novalike MV Lyr has been modelled by Scaringi
(2014) using the fluctuating accretion disc model. From the modelling, the
power spectral break reveals that the variability is driven by a geometrically
thick (and optically thin) disc component which extends from the white dwarf
surface up to a few white dwarf radii (Scaringi, 2014). By analogy, the same
disc component, as inferred from the break in the power spectrum shown in Fig.
7 may also be present in ASAS J0714+7004. This would also be consistent with
the inference of the boundary layer discussed in Sect. 3.2, although the
geometrical extent inferred from the broad-band variability would suggest this
component to be substantially extended.

Figure 7: Power density spectrum of ASAS J0714+7004 computed from the Sector
40 observations, which were obtained with at a 20 s cadence. The orbital
period of 3.28 h is clearly detected. The slight peak at $\simeq$2000s does
not appear in the data of the other sectors. It first appears with a period of
1883 s on day 5 and varies between 1594s and 1931s before disappearing after
day 11. It may be related to intermittent quasi-periodic oscillations. The
overall shape of the broad-band noise can be fit with three simple power laws.
The frequency break at $\approx 10^{-3}$ Hz is typical of accretion-induced
variability. The rise in power at $\approx 10^{-2}$ Hz is attributed to
Poisson noise dominating the power spectrum.

###  3.5 X-ray emission

ASAS J0714+7004 was observed by ROSAT during its all sky survey (Voges et al.,
1999). The ROSAT Position Sensitive Proportional Counter (PSPC) observation,
covering 0.1–2.4 keV, detected a count rate of $0.17\pm
0.023\,\mathrm{cts\,s^{-1}}$. We adopted the flux based on a fit of a stellar
model to the PSPC data ($\mathrm{Flux2}=1.92\times
10^{-15}\,\mathrm{W\,m^{-2}}$) to illustrate the ROSAT detection in Fig. 5.
The stellar model (Voges et al., 1999) is based on an empirical conversion
between count-rates and fluxes following (Schmitt et al., 1995) that was
originally developed to obtain flux values for stars.

Our Swift XRT spectrum consists of a set of counts binned by energy (see Fig.
8) in the range 0.3 to 6 keV. To obtain an estimate for the luminosity we
fitted the data using xspec. We have used a model with two components – one
with a temperature over 16 keV and the other with a temperature of
$0.83\pm{0.15}$ keV. This model was then used to extrapolate upwards into
harder X-rays and downwards into the EUV (see Table 7). Absorption is based on
assuming an interstellar hydrogen column density of $N_{\mathrm{H}}=3.2\times
10^{19}\,\mathrm{cm^{-2}}$ ; in practice this has a minimal effect.

Table 7:  Flux and luminosity based on Swift XRT observations. | Range | Flux | X-ray luminosity
---|---|---|---
keV | Å | ($\mathrm{10^{-12}erg\,s^{-1}\,cm^{-2}}$) | ($\mathrm{10^{30}erg^{-1}}$)
Observed | 0.3–6.0 | 2.1–41 | 1.67 | 9.07
Extrapolated | 0.3–10 | 1.2–41 | 2.33 | 12.65
Extrapolated | 0.1–10 | 1.2–124 | 2.57 | 13.95
Extrapolated | 0.01–10 | 1.2–1242 | 2.66 | 14.44

The Swift pipeline spectrum is shown in Fig. 8 and qualitatively resembles
those of BZ Cam and MV Lyr (Figures 2 and 3 in Balman et al. 2014). The X-ray
luminosity of ASAS J0714+7004 of $1.17\pm{0.27}\times
10^{31}\,\mathrm{erg\,s^{-1}}$ is lower than that of some dwarf novae
(Byckling et al., 2010), although an order of magnitude greater than that of
the average CV (Reis et al., 2013). The optically-thin spectrum and inferred
temperatures are similar to other non-magnetic CVs (Baskill et al., 2005;
Mauche & Mukai, 2002).

Figure 8:  Spectrum of the Swift XRT observations (crosses). We assumed
interstellar absorption fixed at $N_{\mathrm{H}}=3.2\times
10^{19}\,\mathrm{cm^{-2}}$ and solar abundances. The spectrum was then fitted
using the sum of two APEC components with temperatures of $0.8$ keV and $20$
keV respectively, and the sum is shown as the grey line, binned to the
observed energy channels.

We inspected the average count rates for each of our Swift XRT observations ,
and we note that there is mild ($1\,\sigma$) evidence for variability across
the 22 d spanned by the data.

Figure 9:  Phase resolved spectra centred on H$\alpha$ (top), H$\beta$
(middle) and He i 5876 Å(bottom) from the December INT data. The individual
exposures were binned such that each sub-panel in the sequence covers $\simeq
8$ min. The top number in each panel is the orbital phase and the lower number
is HJD$-$2 459 190. The red line is the time averaged spectrum of the December
sample. Note the stochastic variability in the emission lines.

Figure 10:  (left) The trailed spectrum of the H$\alpha$ line in the December
sample and (right) the same data phase-folded using the fitted period and
epoch. The “S" shape is clearly visible. Note also the stochastic peaks in the
left plot that are less evident in the right hand plot due to averaging and
orbital binning. We inspected the trailed spectra of the earlier observations
which provided partial orbital coverage, and did not find repeating orbital
behaviour.

###  3.6 Spectral variability and detection of a wind outflow

The high cadence of our INT spectroscopy reveals rapid variations in the
morphology of the emission lines. This is evident in the Hα trailed spectrum
(see Fig. 10) and is particularly striking in the variability of a P Cygni-
like feature in He I $\lambda$5876 (see bottom panel in Fig. 9). These
variations appear to be stochastic with no relation to orbital phase or indeed
variations in other lines (see Fig. 9).

Figure 11:  Top panel: the python model fit (blue) to the observed average spectrum (red). An inclination of 62$\degr$ was found to best reproduce the observed data. Bottom panel: enlargements of the H$\beta$, He I $\lambda$5876 and H$\alpha$ lines.  Table 8:  Adopted values used for fitting the python model. See Matthews et al. (2015) for more details of the model and parameter definitions.  Parameter | Value
---|---
Wind fill factor |  | 0.008
Wind mass transfer rate |  $\dot{M}$ [$10^{-10}M_{\sun}yr^{-1}$] | 5.0
Mass (white dwarf) |  $M_{1}$ [$\mathrm{M}_{\sun}$] | 0.83
Mass (donor) |  $M_{2}$ [$\mathrm{M}_{\sun}$] | 0.22
Radius (white dwarf) |  $R_{1}$ [$\mathrm{R}_{\sun}$] | 0.0122
Radius (donor) |  $R_{2}$ [$\mathrm{R}_{\sun}$] | 0.30
Inner radius (disc) |  $R_{\mathrm{disc,i}}$ [$\mathrm{R}_{\sun}$] | 0.0244
Outer radius (disc) |  $R_{\mathrm{disc,o}}$ [$\mathrm{R}_{\sun}$] | 0.49
Temperature (white dwarf) |  $T_{\mathrm{1}}$ [K] | 41158
Temperature (donor) |  $T_{\mathrm{2}}$ [K] | 3324
Orbital period |  $P_{\mathrm{orb}}\,$[hr] | $3.2794$
Wind inner radius |  $r_{\mathrm{min}}\,$[$R_{1}$] | $4$
Wind outer radius |  $r_{\mathrm{max}}\,$[$R_{1}$] | $12$
Inner opening angle | $\theta_{\mathrm{min}}\,[^{\circ}]$ | $20.0^{\circ}$
Outer opening angle | $\theta_{\mathrm{max}}\,[^{\circ}]$ | $65.0^{\circ}$
Angle exponent | $\gamma$ | $1$
Terminal velocity | $v_{\infty}\,[v_{\mathrm{esc}}]$ | $3$
Wind Acceleration length | R${}_{\mathrm{v}}\,[10^{11}$cm] | 1.0
Acceleration exponent | $\alpha$ | $1.2$

##  4 Discussion

###  4.1 Winds

The optical P Cygni lines are clear evidence of the presence of winds (Kafka &
Honeycutt, 2004). Our INT spectroscopy demonstrates variability within the
wind structure on the time scale of minutes – probably due to clumping. The
lack of a correlation of the wind features with the orbital phase precludes a
stable and localised origin within the disc. The blue absorption wing of the
He I $\lambda$5876 P Cygni line is unusually strong compared to other well-
studied nova-likes (e.g. Kafka & Honeycutt 2004; Kafka et al. 2009) and this
may be due to a fortuitous inclination. Proga et al. (1998) finds that slower
dense winds are confined to within 45°of the orbital plane. The radial
velocity component of the winds will therefore be greater in high inclination
systems such as ASAS J0714+7004. We did not detect the concentration of winds
at the superior conjunction of the donor as reported by Honeycutt et al.
(2013) in BZ Cam.

To explore whether a wind outflow is plausible in the case of ASAS J0714+7004,
we have calculated simulated spectra for bi-conical outflow using the Monte
Carlo radiative transfer code python (Long & Knigge, 2002), which has been
used in several previous attempts to model winds in disc-dominated cataclysmic
variables at ultraviolet (Noebauer et al., 2010) and optical wavelengths
(Matthews et al., 2015). python starts with a parameterised description of a
biconical outflow, in this case one originally formulated by Shlosman &
Vitello (1993). It then solves for the ionisation structure of the outflow
given the nature of the white dwarf and disc, and then simulates the emergent
spectrum as a function of the inclination angle666python is a collaborative
open-source project available at github.com/agnwinds/python.. The Shlosman &
Vitello (1993) models have a large number of parameters. For the modelling
discussed here, we explored a small fraction of the allowable parameter space,
beginning with models that were similar to those described by Matthews et al.
(2015). For the models we used parameter values consistent with those derived
earlier (Table 5), and varied various of the parameters in Table 8. The
external radiation sources were the white dwarf and the disc; solar abundances
were assumed; H and He were treated as multi-level macro-atoms, while the
metals were treated in the so-called two-level approximation.

A comparison of the simulated to the observed spectra for the best model we
calculated is shown in Fig. 11. The line strengths and line shapes of the H
and He lines are approximately correct. He I $\lambda$5876 shows evidence of a
P Cygni like feature. Despite being produced from a wind that arises from the
disc at distances between 4 and 12 white dwarf radii, the lines in the
simulated spectra and in the observed INT spectra are single peaked. The
reason that the lines are single peaked in the model spectra (at the
resolution of the INT spectra) is likely to be due to two effects. First,
there is emission from fairly far out in the wind which reduces the azimuthal
velocity since the flow is designed to conserve angular momentum. In this
model, the fact that the P Cygni absorption line is $\simeq
600\,\mathrm{km\,s^{-1}}$ implies that the interaction with the wind occurs at
a considerable distance ($\sim 5-6\,\mbox{$\mathrm{R}_{\sun}$}$ i.e. $>5$
times the binary separation) from the system. Second, the blue edge of a
double peaked line tends to be absorbed by material further out in the flow.
With somewhat different parameters one often sees a double H$\alpha$ profile
with the blue wing suppressed.

Although the comparison of model spectra to the observations suggests that
this might be a plausible wind model for ASAS J0714+7004, and although the
model is not that different from those that Matthews et al. (2015) used to
model the spectra of UX UMa and RW Tri, there are a number of critiques of the
model that might be made. First, while the model spectra resemble the observed
spectra, the model spectra are quite sensitive to the inclination angle. The
nature of the simulated spectra change fairly rapidly with inclination angle
near $62\degr$ because this is very close to the outer opening angle of the
conical flow; at lower angles the P Cygni lines grow more prominent; at higher
angles the emission lines grow in strength. Changing the outer opening angle
to a somewhat lower value, would allow one to produce similar spectra at
similarly lower inclination angles. Second, the wind model we use has a mass
accretion rate of $10^{-9}\,\mbox{$\mathrm{M_{\sun}\,yr}^{-1}$}$, which
compares with the mass accretion rate we have derived at high inclination of
$8.9\times 10^{-9}\,\mbox{$\mathrm{M_{\sun}\,yr}^{-1}$}$; this is probably the
maximum value that is plausible. Furthermore we have assumed that the wind is
clumped, and only occupies three per cent of the volume. If one preferred a
lower mass loss rate, one could have simply assumed that the wind occupied an
even smaller fraction of the volume, and obtained a fairly similar spectrum,
since the ionisation parameter at any particular point depends on the density
of the filled fraction of the wind. Finally, we note in passing, that there
are several He i lines that appear double peaked in our models, but which are
single peaked in the time-averaged spectra. Despite these difficulties, we
would suggest that the models indicate that ASAS J0714+7004 has a relatively
wide-angle outflow with properties that are similar to those in the model, and
similar to those in other nova-like systems.

###  4.2 Disc structure

The extensive INT spectroscopy never shows evidence of the double-peaked
emission lines we would expect from a high inclination disc (Horne & Marsh,
1986). We considered the possibility that ASAS J0714+7004 was an SW Sex star
(Dhillon et al., 2013) as this would explain the lack of double-peaked lines
even though it did not have the other SW Sex defining characteristics (deep
continuum eclipses and emission lines that do not share the orbital motion of
the white dwarf and exhibit transient absorption features).

However, in contrast to a large fraction of the confirmed SW Sex stars, ASAS
J0714+7004 is not eclipsing and the trailed spectrum shows no evidence of a
second peak. A wide flared disc, as proposed to explain the lack of double-
peaked lines in SW Sex stars (Dhillon et al., 2013), would require a flare
angle of over 20° which is geometrically unlikely. The alternative explanation
is that the single-peaked emission lines originate in the wind, as discussed
above.

###  4.3 Mass transfer rate

The widely accepted thermal disc instability model (Meyer & Meyer-Hofmeister,
1981; Osaki, 1996) predicts that CVs with a mass transfer rate ($\dot{M}$)
below a critical value ($\dot{M}_{\mathrm{crit}}$) will exhibit outbursts, and
such systems are known as dwarf novae. Nova-likes are CVs with $\dot{M}$ above
this critical value and are expected to have discs in a steady high-luminosity
state. The actual value for $\dot{M}_{\mathrm{crit}}$ depends on the size of
the disc, and hence largely on the orbital period, and is somewhat model
dependent (Shafter et al., 1986; Knigge et al., 2011). The estimated range of
the mass transfer rate for ASAS J0714+7004 is typical for that found among
nova-likes (e.g. table 4 from Hewitt et al. 2020), and places it above
$\dot{M}_{\mathrm{crit}}$.

Figure 12: Light curves of several typical dwarf nova and nova-like sub-
classes. Dates are MJD$-$2 450 000. The panels on the right are enlargements
of the grey areas in the left panels. The dwarf nova Z Cam shows classic
“standstills" of $\simeq 100-200$ d during which there are no outbursts whilst
exhibiting large-amplitude outbursts between standstills; standstills are
normally terminated by a decline to quiescence. This contrasts with the IW And
class (of which V513 Cas is an example) in which characteristic outbursts
occur during standstills. VY Scl stars show relatively deep ($2-5$ mag) low
states which can last a few weeks to about a year, and during which small
outbursts occur. (Observations from the AAVSO International Database,
https://www.aavso.org).

We have used a simple blackbody model for a boundary layer as, in the optical,
it has little effect on the fitted value of $\dot{M}$. However this is not the
whole story. Puebla et al. (2007) concluded that the standard model needed to
be augmented with a modification to the temperature profile of the inner disc,
either by increasing the temperature close to the white dwarf or by the
addition of a wind or corona component (Meyer & Meyer-Hofmeister, 1994).
Balman et al. (2014) analysed three nova-likes and confirmed this conclusion.
These authors suggested that the boundary layer consists of either optically
thin regions traversed by Advection-Dominated Accretion Flows or X-ray corona
regions close to the white dwarf but above the disc plane. This then leads to
speculation as to why optically thin regions of this sort form in nova-likes
but not in dwarf novae. Nixon & Pringle (2019) suggests that the cause is a
magnetically controlled zone that builds up due to the persistently ionised
disc; dwarf nova outbursts do not last long enough for the field to build up.

In this context we note that the X-ray spectrum of ASAS J0714+7004 is less
than one percent of the optical flux which is typical of nova-likes (Figure 3
in Mukai 2017) and clearly shows that the expected 50 per cent of
gravitational energy released at the white dwarf boundary does not escape in
the form of X-rays. ASAS J0714+7004 is therefore another example of where the
standard model fails to account for this energy.

###  4.4 Quasi-periodic outbursts

The photometry in Fig. 2 displays quasi-periodic variability of $\simeq 1$ mag
on time scales of $\simeq 10-20$ d, reminiscent of the “stunted” outbursts
seen in some other nova-likes (Honeycutt et al., 1998). We investigated the
possibility of non-orbital periodic variability and found no evidence for this
despite a number of apparently regular structures. The TESS power spectrum
(Fig. 7) clearly confirms the absence of any other periodicities. This quasi-
periodic variability is distinct from lower level and rapid “flickering"
(Bruch, 1992) which is evident in the TESS light curves (Fig. 2, particularly
panels (e) and (f)). The cause of the stunted outbursts remains unclear. One
possibility is stochastic changes in the mass transfer rate whilst another is
a form of disc instability similar to the cause of dwarf nova outbursts that
involves only a small fraction of the disc. Honeycutt (2001) strongly favoured
the model of dwarf nova outbursts arising from a disc instability – the
apparent stunting of the outbursts being caused by dilution from an additional
source of illumination within the system. Honeycutt (2001) also argued that
the presence of stunted outbursts is not correlated with $\dot{M}$ which
suggests that they are not due to those nova-likes displaying them being close
to $\dot{M}_{\mathrm{crit}}$. Inspecting the long-term light curve of ASAS
J0714+7004 (top panel in Fig. 2) shows that the (dis)appearance of the quasi-
periodic variability does indeed not correlate strongly with the average
brightness of the system.

Robertson et al. (2018) propose a contrary view based on observational
evidence of UU Aqr that about half of the stunted outbursts coincided with
increased hot spot emissions and hence higher $\dot{M}$. Subsequently Schlegel
& Honeycutt (2019) studied an eclipsing nova-like (AC Cnc) and found that the
stunted outbursts were eclipsed by the disc suggesting that either the
outbursts originated on the facing hemisphere of the donor or were due to
changes in the disc.

Closer investigation of an outburst observed at two minute cadence by TESS
(see Fig. 2, panel (g)) shows a gradual rise and fall typical of an “inside
out” dwarf nova. The Swift ultraviolet observations show a similar outburst
with AAVSO optical observations also peaking at the same time. Given that the
bulk of the disc must be ionised to sustain the high $\dot{M}$ we speculate
that there is an outer part of the disc that occasionally becomes sufficiently
cool for H to recombine, and that this outburst starts at the inner edge of
this area and subsequently spreads outwards. As such, we conclude that the
quasi-periodic variability seen in ASAS J0714+7004 is likely related to the
standard thermal disc instability acting in a small outer part of the disc,
leaving the inner disc in a steady hot state which providies the “background”
luminosity that Honeycutt (2001) required to dilute the amplitude of the
outbursts.

###  4.5 Nova-likes and dwarf novae – secular changes in mass transfer rate

The observed temporary change to a low state in 2015
($\mathrm{HJD}-2450000=7150$) is interesting as it leads to a potential
explanation for the under-representation of dwarf novae with orbital periods
between $\simeq 3-4$ hours – just above the so-called “period gap" (Shafter et
al., 1986).

Knigge et al. (2011) showed that populations of dwarf novae and nova-likes
coexist at these orbital periods and so the direct association between orbital
period and mass transfer rate breaks down in this region. Our observation
supports the idea that CVs can change back and forth between dwarf novae and
nova-likes more frequently (a few years) – particularly when mass transfer
rates are close to $\dot{M}_{\mathrm{crit}}$. Nova-likes are brighter (by
$\simeq 5$ mag) than dwarf novae and hence there will be a large selection
effect leading to the under-representation of dwarf novae at such periods.

In the light of our observations it is interesting to compare ASAS J0714+7004
with known sub-classes of nova-likes (see Fig. 12).

Z Cam stars (e.g. Schreiber et al. 2002) are dwarf novae that experience
occasional standstills of constant magnitude of $1-1.5$ mag below maximum
usually following an outburst. The remainder of the time Z Cam stars exhibit
frequent dwarf novae commencing $\simeq 1$ mag below the standstill level and
peaking above it. The bottom right light curve in Fig. 12 shows ASAS
J0714+7004 has a high-level standstill but the amplitude of the outbursts at
other times is far lower than Z Cam.

The IW And stars (Hameury & Lasota, 2014) form a sub-category of Z Cam (see
Fig. 12) and show continual outbursts (not unlike dwarf novae) during the high
state (i.e. during the “standstill"). Again this is not consistent with our
light curve except for our low state in 2015.

The VY Scl stars (Schmidtobreick et al., 2018) are nova-likes that show
occasional $\sim 3-6$ mag low states similar to ASAS J0714+7004 in 2015 and
typically have orbital periods of $3-4$ h. Whilst the cause of the low state
is not well understood it is clear that the previously hot steady-state disc
will continue to deposit its mass onto the white dwarf until it is exhausted –
potentially causing dwarf nova outbursts which we do see (panel (b) in Fig.
2). Based on these observations we conclude that ASAS J0714+7004 is a VY Scl
star.

##  5 Conclusions

Our study of ASAS J0714+7004 has shown it to be a bright and nearby nova-like
CV. It exhibits the low states characteristic of the VY Scl sub-class. We have
obtained a reliable estimate of the orbital period and estimated the mass
transfer rate and inclination (see Fig. 1). We have identified strong wind
signatures that are highly variable – coming and going in a few minutes in
veritable gusts, and uncorrelated with orbital phase. Because of its
brightness and proximity, this system is an excellent target for future
detailed studies of its accretion disc and wind outflow.

##  6 Acknowledgements

We thank Professor David Buckley for completing a detailed review and helpful
comments.

We gratefully acknowledge the photometric observations of ASAS J0714+7004
obtained in 2020 and 2021 by Tamás Tordai (Hungary), Stephen Brincat (Malta),
Kenneth Menzies (USA), Charles Galdies (Malta), Sjoerd Dufoer (Belgium),
Thanasis Papadimitriou (Greece), Richard Sabo (USA), David Boyd (UK), Molly
Wakeling (USA) and Steve Boemer (USA) which we retrieved from the AAVSO
International Database. Based on observations made with the Isaac Newton
Telescope operated on the island of La Palma by the Isaac Newton Group of
Telescopes in the Spanish Observatorio del Roque de los Muchachos of the
Instituto de Astrofísica de Canarias. The IDS spectroscopy was obtained as
part of P9/C48/N1 and N1/C48/P9. This paper includes data collected with the
TESS mission, obtained from the MAST data archive at the Space Telescope
Science Institute (STScI). Funding for the TESS mission is provided by the
NASA Explorer Program. STScI is operated by the Association of Universities
for Research in Astronomy, Inc., under NASA contract NAS 5–26555. We greatly
appreciate the use of the spectres python code to rebin spectra (Carnall,
2017). This work made use of data supplied by the UK Swift Science Data Centre
at the University of Leicester. We acknowledge with thanks the processing of
the UVOT data by Paul Kuin (Mullard Space Science Laboratory, University
College London). We acknowledge that this work made use of the Tych0-2
catalog, This work has made use of data from the European Space Agency (ESA)
mission  Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia Data
Processing and Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. This publication makes use
of data products from the Two Micron All Sky Survey, which is a joint project
of the University of Massachusetts and the Infrared Processing and Analysis
Center/California Institute of Technology, funded by the National Aeronautics
and Space Administration and the National Science Foundation. This publication
makes use of data products from the Wide-field Infrared Survey Explorer, which
is a joint project of the University of California, Los Angeles, and the Jet
Propulsion Laboratory/California Institute of Technology, funded by the
National Aeronautics and Space Administration. The Pan-STARRS1 Surveys (PS1)
and the PS1 public science archive have been made possible through
contributions by the Institute for Astronomy, the University of Hawaii, the
Pan-STARRS Project Office, the Max-Planck Society and its participating
institutes, the Max Planck Institute for Astronomy, Heidelberg and the Max
Planck Institute for Extraterrestrial Physics, Garching, The Johns Hopkins
University, Durham University, the University of Edinburgh, the Queen’s
University Belfast, the Harvard-Smithsonian Center for Astrophysics, the Las
Cumbres Observatory Global Telescope Network Incorporated, the National
Central University of Taiwan, the Space Telescope Science Institute, the
National Aeronautics and Space Administration under Grant No. NNX08AR22G
issued through the Planetary Science Division of the NASA Science Mission
Directorate, the National Science Foundation Grant No. AST-1238877, the
University of Maryland, Eotvos Lorand University (ELTE), the Los Alamos
National Laboratory, and the Gordon and Betty Moore Foundation. We gratefully
acknowledge NASA’s support for construction, operation, and science analysis
for the GALEX mission, developed in cooperation with the Centre National
d’Etudes Spatiales of France and the Korean Ministry of Science and
Technology. Support for the ATLAS survey was provided by NASA grant NN12AR55G
under the guidance of Lindley Johnson and Kelly Fast. We have made use of the
ROSAT Data Archive of the Max-Planck-Institut für extraterrestrische Physik
(MPE) at Garching, Germany. This research has made use of the XSPEC tool
obtained from the High Energy Astrophysics Science Archive Research Center
(HEASARC), a service of the Astrophysics Science Division at NASA/GSFC and of
the Smithsonian Astrophysical Observatory’s High Energy Astrophysics
Division." BTG and TRM were Leverhulme Research Fellowships and the UK STFC
grant ST/T000406/1.

##  7 Data availability

The INT spectroscopy will be available via the data archives of the ING after
the expiry of the default proprietary period. All the time-series photometry
is available via the archives of the AAVSO and TESS mission.

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##  Appendix A From Gaia DR2 to EDR3: improved precision and overall changes
due to variability

Gaia EDR3 is a complete re-release of the astrophysical and photometric data
and inevitably the absolute magnitudes and colours of any source are expected
to differ somewhat with respect to DR2. For the set of nova-likes shown in
Fig. 6, the uncertainties of the $G$, $G_{\mathrm{BP}}$, $G_{\mathrm{RP}}$
magnitudes and the parallaxes were 0.0101, 0.0355, 0.0308 and 0.0102
respectively in EDR3 compared to 0.0144, 0.0423, 0.0346 and 0.0216 in DR2. Fig
13 shows the corresponding changes in the Gaia HRD as black arrows. In most
cases, these changes are quite small but one system (MV Lyr) stands out as its
absolute magnitude decreased by $\simeq 1$. The reason is an important
reminder that nova-likes are variable stars and Gaia EDR3 essentially computes
a weighted average of measurements taken between 25 July 2014 and 28 May 2017
(whereas DR2 data collection stopped on 23 May 2016). Fig 14 shows that MV Lyr
was in a prolonged low state for about 2/3 of the DR2 observations, whereas
EDR3 sampled the system for a substantially longer period in the high state –
explaining the large change in $G$, $G_{\mathrm{BP}}$, and $G_{\mathrm{RP}}$.

Figure 13:  Selected nova-likes in the Gaia EDR3 HRD showing the changes
(black arrows) in absolute magnitude and colour between Gaia DR2 and EDR3.
Whilst some minor changes are due to refinement in Gaia astrometry and
photometry the major changes are due to transitory changes in the observed
light curve. MV Lyr is a good example of this where it was mostly in a low
state during the 162 observations contained in DR2 but in a high state during
the remaining 107 observations in EDR3) see Fig. 14)  Figure 14:  Long-term
light curve of MV Lyr, obtained from the AAVSO data base. The arrows indicate
the date range covered by the observations contained in Gaia DR2 and EDR3,
illustrating why the absolute magnitude of this system dropped significantly
in EDR3 compared to DR2.