# The Peculiar Bursting Nature of CP Pup

M. Veresvarska,1 S. Scaringi,1 S. Hagen,1 D. De Martino,2 C. Done,1 K.
Ilkiewicz, 1,3 C. Knigge, 4 C. Littlefield5
1Centre for Extragalactic Astronomy, Department of Physics, Durham University,
South Road, Durham, DH1 3LE
2INAF-Osservatorio Astronomico di Capodimonte, Salita Moiariello 16, I-80131
Naples, Italy
3Astronomical Observatory, University of Warsaw, Al. Ujazdowskie 4, 00-478
Warszawa, Poland
4School of Physics and Astronomy, University of Southampton, Highfield,
Southampton SO17 1BJ, UK
5Bay Area Environmental Research Institute, Moffett Field, CA 94035, USA
E-mail<EMAIL_ADDRESS>
(Accepted XXX. Received YYY; in original form ZZZ)

###### Abstract

The classical nova CP Puppis has been observed to have particularly puzzling
and peculiar properties. In particular, this classical nova displays
occasional bursts in its long-term ASAS-SN light curve. Here we report on 5
sectors of TESS data displaying 2 of these rapid bursts, lasting $\sim$ 1 day.
Based on the estimated lower energy limits of the bursts we discuss whether
the bursts may be examples of micronovae resulting from localised
thermonuclear explosion. Furthermore, its orbital period remains uncertain,
with several inconsistent periodic signals appearing in spectroscopic and
photometric observations at various wavelengths. Although we cannot
unambiguously unravel the physical origin of the signals, the previously
suggested nature of CP Puppis as a long orbital period system may be a viable
explanation. The recurrence time of the bursts in CP Puppis, together with the
unexplained variable modulations make it a prime candidate for intense
monitoring.

###### keywords:

accretion – accretion discs – novae, cataclysmic variables – individual: CP
Pup

††pubyear: 2023††pagerange: The Peculiar Bursting Nature of CP Pup–The
Peculiar Bursting Nature of CP Pup

##  1 Introduction

CP Puppis (CP Pup hereafter) is a well studied cataclysmic variable that has
undergone a nova explosion in 1942. It is an unusually bright and fast example
of a nova explosion, with the difference in amplitude of apparent magnitude of
$\sim 17$ mag (Payne-Gaposchkin, 1964). It is also a particularly fast nova
with $t_{3}\sim 6.5$ d (Payne-Gaposchkin, 1964), where $t_{3}$ represents the
time it takes for the nova luminosity to decline by 3 magnitudes. It has been
reported that since the nova explosion the brightness levels of the nova have
yet to return to the pre-burst quiescence level (Schaefer & Collazzi, 2010).
Similarities can be drawn between CP Pup and V1500 Cyg (Della Valle & Livio,
1998), which is also uncharacteristically brighter and has also remained
brighter post eruption. CP Pup is also suggested to have a magnetic accretor
(Balman et al., 1995; Orio et al., 2009; Mason et al., 2013).

One of the main peculiarities of CP Pup is its elusive orbital period. There
have been spectroscopic (Bianchini et al., 1985; Duerbeck et al., 1987;
O’Donoghue et al., 1989; White et al., 1993; Bianchini et al., 2012; Mason et
al., 2013) and photometric (Warner, 1985; O’Donoghue et al., 1989; Diaz &
Steiner, 1991; White et al., 1993; Bruch, 2022) periods reported between
$0.06115$ days and $0.06834$ days. A summary of all these is provided in Bruch
(2022). Both spectroscopic and photometric periods reported in literature lie
in a similar range, with most of the spectroscopic periods being shorter.
Furthermore, the spectroscopic periods have been reported to show large
scatter in the folded radial velocity curve (Bianchini et al., 2012). It is
worth noting that the spectroscopic and photometric periods in the literature
are inconsistent with each other (Diaz & Steiner, 1991). Orio et al. (2009)
has further reported 3 periods from XMM-Newton data. However, due to the large
uncertainty, the periods quoted in Orio et al. (2009) are consistent with each
other, as well as consistent with previously reported spectroscopic and
photometric periods (see summary by Bianchini et al. (2012) and Bruch (2022)).

If taken at face value as the orbital period, all the reported values are
short for known classical novae and non-magnetic nova-likes in general ($\sim
3$ hours), making CP Pup one of a handful of novae with a potential period
reported below the period gap (Bianchini et al., 1985; Marelli et al., 2018).
It has thus been suggested that CP Pup is a system whose disc is always hot
and bright, and often exhibits superhumps, persisting in a state similar to SU
UMa type dwarf novae during superoutburst (Warner, 1985; Patterson & Warner,
1998; Patterson et al., 2013).

Further interpretations of the period assigned it to the spin of the magnetic
white dwarf, suggesting a white dwarf slightly out of spin-orbit synchronism
(Warner, 1985; Diaz & Steiner, 1991; Balman et al., 1995). Asynchronous polars
are not the only option for the magnetic interpretation of CP Pup. Orio et al.
(2009) and Mason et al. (2013) consider the hypothesis of CP Pup being an
intermediate polar. However, there is a lack of conclusive evidence for either
interpretation.

The peculiar nature of CP Pup extends to other observations too. The
relatively recent nova explosion, with its characteristics such as peak
absolute magnitude, suggests a high mass white dwarf (WD), due to the smaller
physical accretor size resulting in higher pressure, making it more favourable
to trigger a nova explosion (Prialnik & Kovetz, 1995). However, White et al.
(1993) points out that the reported values of $K_{1}\sim 70-146$ kms-1 require
a WD mass $<0.2M_{\odot}$. This is in contradiction to the nova nature of CP
Pup. Some more recent studies adopting higher WD mass of 1.1 $M_{\odot}$ have
instead obtained better fits to optical spectra with an accretion rate in the
range of $3.3-7.3\times 10^{-10}M_{\odot}yr^{-1}$ (Mason et al., 2013). Also
assuming a high mass WD, $M_{WD}>1.1M_{\odot}$, Orio et al. (2009) derived an
upper limit on mass transfer rate $\lesssim 1.6\times
10^{-10}M_{\odot}yr^{-1}$. However, this estimate assumes a distance of 1600
pc, more than double the distance of 780 $\pm$ 11 pc inferred from the Gaia
parallax (Gaia Collaboration et al., 2016, 2023). Orio et al. (2009) also
quotes a $3\sigma$ upper limit $\lesssim 8\times 10^{-11}M_{\odot}yr^{-1}$ at
850 pc, as determined by Cohen & Rosenthal (1983). Other estimates of the
accretion rate adopting the more precise Gaia parallax measurement have been
made. Selvelli & Gilmozzi (2019) reports a considerably higher value of
$10^{-9.22\pm 0.32}$ $M_{\odot}$yr-1. This is related to the unusually high
extinction of $A_{v}\sim 0.6$ mag, which they adopt.

In this work we report on the analysis of 5 sectors of TESS data of CP Pup
(Section 2). In Section 3 we discuss 2 rapid bursts observed as well as the
phenomenology of the cluster of periods found in the TESS light curves.
Reanalysis of XMM-Newton data of CP Pup is also discussed here and the new WD
mass estimate is derived. In Section 4 we discuss the observed rapid bursts in
the context of the micronova model (Scaringi et al. (2022b, c).

##  2 Observations

The data used in this work was obtained by TESS and is available on the
Mikulski Archive for Space Telescopes
(MAST111https://mast.stsci.edu/portal/Mashup/Clients/Mast/Portal.html). TESS
observed CP Pup over 5 sectors amounting to $\sim$ 5 months of observations at
2 minute cadence. The last sector 61 also contains 20 second cadence data. The
detailed description of the data with the burst characteristics is provided in
Table 1. TESS provides 2 modes for output flux. Here the SAP flux is used, as
opposed to the processed version of PDCSAP. This is to avoid any contamination
of the data through pipeline processing. The flux is then cleaned with a mask
discarding any data points with quality flag $>$ 0, to mitigate presence of
artificial or non-intrinsic features in the data.

The TESS data shows a modulation of the shape of the light curve in the form
of a larger envelope, where the variance and flickering temporarily varies on
timescales of $\lesssim 10$ days (see Fig. 1). The same envelope is present to
a greater extent in the 20 s cadence data of Sector 61 in TESS . The Lomb-
Scargle periodogram of the 20 s data shows a QPO-like feature at $\approx
1500$ cycles$/$day. This is also present in all the neighbouring stars in the
field of CP Pup and hence assumed to not be intrinsic to the CP Pup itself,
but rather a common systematic. Similar behaviour is also observed in other
TESS 20 s observations in Sector 61 (such as TIC 260266504, TIC 289113766 and
TIC 289113764). This points to an instrumental origin of the QPO and is hence
disregarded. As no other new signals are found in the 20 s cadence data of
Sector 61, the 2 minute cadence is used throughout for consistency. The data
is downloaded and cleaned for cosmic rays through the Lightkurve
package222https://docs.lightkurve.org/index.html.

Table 1: Summary of the TESS data of CP Pup, with sector numbers and dates. If a sector contains a burst, the corresponding date in Barycentric Kepler Julian date, duration in days and energy is noted. The details of energy calculation are described in Section 3.1. Sector | Dates | burst dates | Duration | Energy
---|---|---|---|---
| (BTJD-2457000) | (BTJD-2457000) | (d) | (erg)
7 | 1491.6 - 1516.1 | 1491.6 - 1492.2 | $\sim 0.6$ | $1.2\times 10^{38}$
8 | 1518.0 - 1542.0 | $-$ | $-$ | $-$
34 | 2229.0 - 2254.1 | 2229.0 - 2234.4 | $\sim 1$ | $6.3\times 10^{37}$
35 | 2255.7 - 2280.0 | $-$ | $-$ | $-$
61 | 2964.0 - 2988.1 | $-$ | $-$ | $-$

While TESS provides an excellent precision relative photometry, it is better
to leverage data from another observatory to obtain absolute photometry from
TESS data. To convert the TESS data to erg s-1 it is necessary to consider
ground based observations as well as knowledge of its distance. Quasi-
simultaneous ASAS-SN g-band data (Shappee et al., 2014; Kochanek et al., 2017)
is used for this purpose. The band is centered at $475$ nm with $140$ nm
width, with a partial overlap with the TESS band-pass ($600$ \- $1000$ nm).
Although the pass-bands do not exactly overlap, we here assume that any
colour-term variation are minimal. The data is obtained from the ASAS-SN
webpage333https://asas-sn.osu.edu/. The conversion is done through
simultaneous observations for each half sector. Simultaneous data consist of
all ASAS-SN observations within a TESS cadence observation of 2 minutes.
Assuming a linear relation between the 2 bands a direct conversion can be
established in the form of $F_{ASAS-SN}\left[mJy\right]=A\times
F_{TESS}\left[e^{-}s^{-1}\right]+C$. The corresponding values of the
coefficients for each half-sector are specified in Table 2 to account for any
deviations in data due to the gap in the middle of TESS sectors.

Table 2: Summary of the conversion coefficients from TESS flux in $e^{-s}$ to ASAS-SN flux in mJy for CP Pup. As the conversion is done twice per TESS sector, all corresponding coefficients are listed. Sector | Sector half | A $\bigg{(}\frac{mJy}{e^{-}s^{-1}}\bigg{)}$ | C (mJy)
---|---|---|---
7 | 1 | 0.0241 $\pm$ 0.0007 | -1.2 $\pm$ 0.1
| 2 | 0.016 $\pm$ 0.002 | -0.0 $\pm$ 0.3
8 | 1 | 0.011 $\pm$ 0.001 | -0.9 $\pm$ 0.4
| 2 | 0.011 $\pm$ 0.001 | -0.9 $\pm$ 0.4
34 | 1 | 0.0209 $\pm$ 0.0008 | -11.1 $\pm$ 0.5
| 2 | 0.011 $\pm$ 0.001 | -4.6 $\pm$ 0.9
35 | 1 | 0.015 $\pm$ 0.001 | -2.2 $\pm$ 0.4
| 2 | 0.014 $\pm$ 0.003 | -1.8 $\pm$ 0.9
61 | 1 | 0.013 $\pm$ 0.001 | -1.8 $\pm$ 0.4
| 2 | 0.008 $\pm$ 0.002 | -0.5 $\pm$ 0.7

There are 2 cases where the calibration between TESS and ASAS-SN fails. The
first is in Sector 8, where the TESS light curve shows a drop in flux
inconsistent with the ground-based ASAS-SN observations. This is also reported
in Bruch (2022). Another similar case occurs in Sector 35 where the TESS data
shows an uncharacteristic rise in flux, but this appears to be shared within
the entire region surrounding CP Pup and we attribute this to poor background
correction. Consequently both segments of Sectors 8 and 35 where the
calibration to ASAS-SN fails are discarded. The final result is shown in
Figure 1, where the light curve is also corrected for the distance of $780\pm
11$ pc inferred from Gaia DR3 parallax. Despite the high extinction in Gaia
passband, no bolometric correction has been applied. This is to avoid any
effects due to a potential change in the SED between quiescent and bursting
parts of the light-curve, which would affect any bolometric correction. As a
result the obtained luminosity of CP Pup should be treated as a lower limit.
It is however possible to estimate the effect the bolometric correction could
have on the luminosity. The X-ray luminosity in Section 3.4 is estimated to be
$L_{X}~{}\sim 1.1\times 10^{33}$ erg s-1. Similarly, Orio et al. (2009)
estimated the UV luminosity to be $L_{UV}\sim 2\times 10^{33}$ erg s-1.
Alongside with the mean TESS luminosity of $L_{{\it TESS}\ }\sim 1\times
10^{33}$ erg s-1 the total bolometric luminosity is expected to be of the
order of $\sim 4\times 10^{33}$ erg s-1. Therefore the lower limits on the
luminosity are expected to represent about $\sim 4\times$ lower values than
the total bolometric values.

Figure 1: Light curve of CP Pup showing all 5 sectors of TESS data as
described in Table 1 (orange, green, red, black and purple consecutively for
sectors 7, 8, 34, 35 and 61). Overlaid (in blue) is the ASAS-SN light curve
used for calibration. The grey points are data from ASAS-SN for which
calibration failed and hence were excluded. The flux of both TESS and ASAS-SN
light curves is converted to erg s-1 using the Gaia parallax.

Further data used in this work has been obtained by XMM-Newton on 4th of June
2005 and reported in Orio et al. (2009). As described in greater detail in
Orio et al. (2009), the data was obtained with EPIC-pn, MOS-1 and MOS-2 with
$\sim 50$ ks exposure. Grating spectrum with low SNR was also obtained with
RGS and reported in Orio et al. (2009). The spectrum is shown in Orio et al.
(2009) in Figure 3. The reduced data is available online at
heasarc444https://heasarc.gsfc.nasa.gov/cgi-bin/W3Browse/w3browse.pl. However,
the data reanalysed here, is only the data taken with the most sensitive EPIC-
pn instrument.

##  3 Results

In this section we report the bursts found in 2 of the 5 available TESS
sectors in Section 3.1. We further report on the observed variable cluster of
frequencies found in the power spectrum of CP Pup in Section 3.2. The XMM-
Newton spectrum reported in Orio et al. (2009) and our new model fit is
presented in Section 3.4

###  3.1 Bursts

The TESS light curve from all sectors is displayed in Figure 1, where two
bursts are detected in Sector 7 and 34 with the corresponding times in Table
1. The bursts have also been reported in Bruch (2022), however they are
disregarded in the analysis. The recorded bursts in Sectors 7 and 34 all
lasted $\lesssim 1$ day, with the overall amplitude increasing by a factor of
$>2$, from $\sim 1.5\times 10^{33}$ erg s-1 to $>3.0\times 10^{33}$ erg s-1. A
detailed version of the bursts is shown in Figure 2. Only Sector 34 contains
both the immediate pre-burst and post-burst observation. Sector 7 started
during the rise of the burst and contains the post-burst part only. It is
worth noting an additional burst of similar nature in the long term ASAS-SN
light curve at $\sim 3005$ BTJD in Figure 1. Unfortunately there are no TESS
observations around that date. Another interesting characteristic of the
bursts is their shape, showing 2 distinct peaks with $\sim$ 0.35 d separation.

To understand the origin and nature of the bursts, their energy has to be
determined. This is done by integrating the luminosity under the calibrated
light curve after subtracting the baseline luminosity level. We estimate the
baseline luminosity using the pre- and post burst data only. We use these to
compute a running mean and interpolate in between the two with a spline
function. The resulting energy in ergs is reported in Table 1. The most
energetic burst in Sector 7 releases in excess of $1.2\times 10^{38}$ erg. The
energy released by the burst in Sectors 34 is $>6.3\times 10^{37}$ erg. We
note that the aforementioned data gaps, as well as the lack of a bolometric
correction allows us to only provide lower limits to the burst energies.

Figure 2: Zoom in on the bursts shown in Figure 1 and described in Table 1.
The energy of each burst is $1.2\times 10^{38}$ erg for sector 7 and
$6.3\times 10^{37}$ erg for sector 34.

###  3.2 Phenomenology of the period cluster

A peculiar feature of CP Pup is the cluster of periods between $\sim 14$
cycles$/$d and $\sim 15$ cycles$/$d ($\sim 1.5-1.6$ hr). In the 5 TESS sectors
reported here all show this variability. This time-dependent variability is
discussed in detail in Bruch (2022). The persistent signal reported from 4
TESS sectors in Bruch (2022) is at $\sim$ 16.29 cycles$/$d. Independently
redoing the period analysis on all 5 TESS sectors also reveals a single
identical persistent signal in the cluster at 16.29 $\pm$ 0.04 cycles$/$d.

In Figure 3 the cluster is shown right before and after the burst present in
sector 34. There is no apparent connection between the burst and the position
or shift of the cluster frequencies. However, it is evident that there is
considerably more rms power in the cluster before the bursts than after. For
sector 34 the peak power of the cluster is at $\sim 590$ RMS, whereas after
burst it is at $\sim 110$ RMS, excluding the power of the signal at $\sim
16.29$ cycles$/$d at $\sim 200$ RMS. Similarly, sector 7 shows a significant
quasi-periodic variability as well. However lack of data pre-burst does not
allow for similar comparison as in sector 35.

Orio et al. (2009) also reports only one coherent period in the light curve of
CP Pup of 16 $\pm$ 2 cycles$/$d ($\sim$90 minutes). We recover the same signal
from the XMM-Newton data. The signal is marked by the dashed grey line in
Figure 3. Furthermore, this period also corresponds to the only persistent
signal in TESS light curve.

Figure 3: PSD of sectors 34 divided into pre-burst (blue) and post-burst
(orange) parts. Grey dashed line shows the coherent XMM-Newton period at $\sim
16$ cycles$/$d.

###  3.3 The potential nature of the orbital period

Figure 4: Left: model SED of CP Pup for an estimated orbital period of 1.473
hr with the dotted (blue) line representing the WD, dashed (orange) line the
secondary and the dashed-dotted (green) line the accretion disc. The total SED
is shown in solid (red) line. For comparison the measurements from WISE W1 and
W2 are shown (green squares) as well as the 2MASS filters (blue dots) and the
Gaia measurements (orange diamonds). The TESS bandpass is shown in the shaded
(pink) region. Right: Same SED components and data as in the left panel but
for orbital period of $\sim$ 5.6 hr, the longest orbital period included in
the evolutionary tracks in Knigge et al. (2011).

The signal at 1.473 hr is interpreted in literature as the orbital period of
CP Pup. Using the evolutionary tracks of Knigge et al. (2011) we can estimate
the temperature of the donor to be $\sim 2700$ K with radius of $\sim$ 0.131
$R_{\odot}$ for a system with an orbital period of $\sim 1.473$ hr. The
corresponding white dwarf temperature according to Knigge et al. (2011) would
be 11 000 K with radius following from the WD mass in Section 3.4. In such a
case the SED of the system can be roughly approximated by two spheres
radiating as blackbodies with the corresponding radii and temperatures.
Another possible dominant SED component would be the accretion disc. Taking
the X-ray derived accretion rate of $\sim$ 1 $\times$ 10-10 $M_{\odot}$yr-1
(Section 3.4) at face value, a standard Shakura-Sunyaev (Shakura & Sunyaev,
1973) accretion disc from the WD surface to the tidal radius at $\sim$0.31
$R_{\odot}$ is assumed. The resulting SED is hence shown in the left hand side
panel of Figure 4. It is evident that there is a deficit in flux across all
observed wavebands that cannot be corrected by either an increase in accretion
rate or by correcting for extinction ($A_{V}\sim 0.6$ mag (Selvelli &
Gilmozzi, 2019), $A_{H}=$0.107 mag and $A_{J}=$0.173 mag (Cardelli et al.,
1989) and $A_{W1}=$0.203 mag and $A_{W2}=$0.157 mag (Fitzpatrick, 1999;
Indebetouw et al., 2005).

Mason et al. (2013) discusses the possibility that CP Pup may be a long
orbital period CV. Despite no other photometric signal being detected, Mason
et al. (2013) reports a substantial power at low frequencies ($\sim 2.38$
cycles/d). Their H$\beta$ RV curve produces a fit of 9.8 hr. They point out a
large uncertainty of 1 day, making this a lower limit on the potential signal.
The longest orbital period we can use from the evolutionary tracks of Knigge
et al. (2011) at an orbital period of 5.6 hr provides an estimated donor
temperature of $\sim$ 3900 K and radius of $\sim$ 0.63 $R_{\odot}$. Assuming
the same parameters for the WD and accretion disc as previously (with the
slight correction for tidal radius being at $\sim$ 0.19 $R_{\odot}$) the
estimated SED is as shown in the right panel of Figure 4. This is more
consistent to the observed 2MASS and WISE absolute magnitudes of CP Pup,
suggesting a hotter and larger donor than previously anticipated. Since the
evolutionary tracks in Knigge et al. (2011) do not include orbital periods
above 5.6 hr, the right panel in Figure 4 represents an estimate, not a fit,
of the effect a longer orbital period would have on the SED. This also
tentatively suggests a longer orbital period for CP Pup than reported in the
past, but a more intensive radial velocity campaign is required to unravel the
problem of the orbital period of CP Pup.

###  3.4 Revisiting the WD mass and mass accretion rate with X-Rays

The reanalysis of the X-ray spectrum in this section is motivated by the
inconsistent estimates of the WD maximal temperature $T_{\mathrm{max}}$ in the
literature (Orio et al., 2009; Mason et al., 2013), leading to an uncertain WD
mass and unconstrained accretion rate. This is also exacerbated by the pre-
Gaia distances used in past X-ray analyses. More specifically Orio et al.
(2009) infer large values of $T_{\mathrm{max}}$ following from the used models
for non-magnetic systems and a simplified model of magnetic systems. Exploring
the possibility of micronova eruptions present in CP Pup, the model
implemented in this work is adapted specifically for magnetic accreting WD
with an accretion column.

Figure 5 shows the archival XMM-Newton spectrum of CP Pup taken on 4th June
2005, reported in Orio et al. (2009). Orio et al. (2009) fit the RGS spectrum
using mkcflow model (Mushotzky & Szymkowiak, 1988) model which describes the
emission from a cooling multi-temperature plasma and the EPIC-pn spectrum with
a multi-temperature APEC model. However their analysis does not include the
combination with a complex absorber which is instead characteristic of
magnetic CVs (Lopes de Oliveira & Mukai, 2019; Islam & Mukai, 2021) while they
use a single simple absorber.

Orio et al. (2009) also uses APEC multi-temperature model to fit the XMM-
Newton EPIC-pn spectrum. The APEC model consists of 3 temperature components
and obtains a spectrum for a collisionally-ionized diffuse gas. For an IP or
polar we can expect the presence of a magnetically confined accretion flow
onto the magnetic pole (or poles) of the accretor. In this case the impacting
material generates a shock heated plasma, which cools towards the surface of
the WD, and generates the X-ray emission. This is similar to the multi-
temperature APEC model used by Orio et al. (2009).

Testing the possibility that the bursts can be explained by micronova requires
a magnetic accretor. Therefore, we fit the spectrum using xspec version
12.13.0 (Arnaud, 1996), and specifically use the cemekl model (Singh et al.,
1996) which describes the emission from a multi-temperature plasma. The
geometry of an accretion stream onto a WD surface can be approximated as
cylindrical (see Figure 1 in Done & Magdziarz 1999). X-rays emitted on the far
side of the stream (with respect to the observer) encounter a larger effective
column-density to the observer compared to those emitted on the side facing
the observer. To account for this we convolve cemekl with the absorption model
pwab (Done & Magdziarz, 1998), which calculates the emission assuming the non-
uniform absorption profile arising from a cylindrical emitter. We further
convolve the model with relfect (Magdziarz & Zdziarski, 1995), which
calculates the reflected emission from neutral material. The use of the sc
reflect component follows from the presence of the Fe-K$\alpha$ line.
Additionally we add two components to model, the fluorescence lines
Fe-K$\alpha$ ($\sim 6.4$ keV), and O vii ($\sim 0.58$ keV). The strong O vii
line might be due to the wind from the nova shell, as noted in Orio et al.
(2009). Finally, we also include tbabs to account for interstellar absorption
as obtained from the visual extinction from Table 1 in (Selvelli & Gilmozzi,
2019), using the optical extinction to H column density relation (Guver &
Ozel, 2009). The final xspec model is then: tbabs*pwab*reflect*(cemekl +
gaussian(Fe-K$\alpha$) + gaussian(O vii)). The best fit model is shown in
Figure 5, and the corresponding model parameters are shown in Table 3.

Table 3 shows the resulting best fit parameters. In particular,
$T_{\mathrm{max}}=29.3^{+9.4}_{-7.4}$ keV, significantly lower than $\sim 64$
keV for the multi-temperature APEC model and $\sim 78$ keV for the MKCFLOW
model discussed in Orio et al. (2009). Although $T_{\mathrm{max}}$ is loosely
constrained it does not approach the maximum allowed limit of the model. We
attribute the main difference between this fit and that of Orio et al. (2009)
to the inclusion of the pwab component, which represents an approximation for
the accretion column. Our new estimate is consistent with the range determined
by Mason et al. (2013) using Chandra data, which has considerably softer
coverage. We note that although $N_{h,\mathrm{max}}$ is unconstrained its
effect on the overall fit is negligible. This particular behaviour of
$N_{h,\mathrm{max}}$ is also reported in Done & Magdziarz (1998) for BY Cam.
Similarly for reflection, which is loosely constrained but with negligible
effect on the final value of $T_{\mathrm{max}}$. The possibility of fixing
reflection to 1 is also explored. It was found that a fixed value of
reflection changes $T_{\mathrm{max}}$ to
$T_{\mathrm{max}}=29.3^{+10.4}_{-7.1}$ keV, however all variations of the
temperature are well within the errors, making a fixed value of reflection
unnecessary. Done & Magdziarz (1998) explored the possibility of fixing
$\alpha=1$, however found that the best fit is obtained for $\alpha<1$. This
is also the case for CP Pup and is indicative of multiple sources of cooling,
not only Bremsstrahlung, most likely cyclotron.

To infer the WD mass and mass accretion rate it is here assumed that CP Pup is
a magnetic system. This choice was based on partial evidence as presented in
Mason et al. (2013) and other works, but also in order to explore the
micronova nature of the bursts. From the updated X-ray fit we first isolate
the plasma component from the model spectrum, extrapolating it over the energy
range for which the model is defined ($5\times 10^{-3}-60$ keV) and
integrating it. This results in a bolometric flux of $\sim 1.5\times 10^{-11}$
erg s-1cm${-2}$. Using the Gaia inferred distance of $D=780\pm 11$ pc, this
gives $L_{x}\sim 1.1\times 10^{33}$ erg s-1. The WD mass is derived from
$T_{\mathrm{max}}$ using Equation 1, relating the shock temperature to the WD
mass in magnetic cataclysmic variables, in Orio et al. (2009) following from
Wu et al. (2003) and references therein. The resulting mass is $M_{WD}\sim
0.73M_{\odot}$${}^{+0.12}_{-0.11}$ with the corresponding radius $R_{WD}\sim
0.011R_{\odot}$${}^{+0.001}_{-0.002}$, assuming the mean molecular weight from
the abundance in Table 3.

Assuming the inferred $M_{WD}$ is correct, we find a mass accretion rate
$\dot{M}=1-2\times 10^{-10}\,M_{\odot}$ yr-1. The accretion rate is quite low,
however estimating the optical luminosity from Gaia magnitude yields $\sim
8.8\times 10^{32}$ erg s-1. Therefore estimating the accretion rate from the
optical luminosity would lead to similarly low values. As it is likely that
the accretion luminosity is not solely constrained to X-rays it is reasonable
to assume that the accretion rate above is a lower limit.

Another method of estimating the accretion rate is through the optical to NIR
SED. Using the photometric observations of CP Pup available in
vizier111https://vizier.cfa.harvard.edu/vizier/sed/, the peak of the SED is
placed at XMM-Newton optical Monitor at $\sim 543$ nm with 2.26 $\pm$ 0.06
mJy. This translates to a luminosity of 2.53 $\pm$ 0.02 $\times$ 1033 erg s-1
using a distance of 780 $\pm$ 11 pc. Assuming that the luminosity corresponds
to the accretion luminosity $L_{acc}=\frac{GM_{WD}\dot{M}}{2R_{in}}$, the
estimated accretion rate is $\dot{M}\sim$ 6 $\times$ 10-10 \- 6 $\times$ 10-9
$M_{\odot}$yr-1. The range of estimated accretion rate depends on the possible
truncation of the inner accretion disc. The values taken into account here
estimated $R_{in}\sim R_{WD}-10\times R_{WD}$, where $10\times R_{WD}$ is a
value representative of the estimated truncated disc radii in IPs from
(Suleimanov et al., 2019) for systems of similar mass to CP Pup. Whereas this
only represents a very rough estimate of the accretion rate, it corresponds to
the same order of magnitude as the X-ray accretion rate. Integrating the
entire available SED, with a cut-off after $W_{2}$ WISE band, would result in
total luminosity of $\sim$ 4 $\times$ 1034 erg s-1. This would correspond to
$\dot{M}\sim$ 9 $\times$ 10-9 \- 9 $\times$ 10-8 $M_{\odot}$yr-1, a
significantly higher value than previous estimates. The excess may be
accounted for if part of the SED is dominated by the donor star as suggested
in Section 3.3. In such a case the approximation of accretion luminosity
through the SED peak provides a more accurate estimate.

It is important to note that we do not have sufficiently high spectral
coverage to confidently determine the maximum plasma temperature,
$T_{\mathrm{max}}$, which sets the roll-over in the spectrum, and should
therefore also affect the derived luminosity. Hence, the above estimates
should be treated with caution, noting the model dependence. Furthermore, the
accretion rate range above is an lower limit determined by the underlining
assumptions.

Figure 5: Unfolded XMM Epic-PN X-Ray spectrum of CP Pup. The green points show the data, while the solid black line shows the total model. The dotted gaussians show the two additional lines needed to fit the spectrum. These are (from left to right): O vii and Fe-K$\alpha$ Parameter | Units | Value
---|---|---
———– tbabs ———–
$N_{H}$ |  $10^{22}$ cm-2 | 0.16
———– pwab ———–
$N_{h,\mathrm{min}}$ |  $10^{22}$ cm-2 | $1^{+1625}_{-1}\times 10^{-7}$
$N_{h,\mathrm{max}}$ |  $10^{22}$ cm-2 | $7.34^{+\inf}_{-2.32}$
$\beta$ |  | $-0.86^{+0.04}_{-0.09}$
———– reflect ———–
rel_refl |  | $1.06^{+0.94}_{-1.04}$
Redshift |  | 0
abund |  | $0.57^{+0.11}_{-0.11}$
Fe_abund |  | 1
$\cos(i)$ |  | 0.5
———– cemekl ———–
$\alpha$ |  | $0.60^{+0.10}_{-0.04}$
$T_{\mathrm{max}}$ | keV | $29.3^{+9.4}_{-7.4}$
$n_{h}$ | cm-3 | 1
abund |  | $0.57^{+0.11}_{-0.11}$
Redshift |  | 0
switch |  | 1
norm | $10^{-3}$ | $5.9^{+0.9}_{-0.7}$
———– gaussian (Fe-K$\alpha$) ———–
$E$ | keV | 6.4
$EW$ | keV | 0.07
$\sigma$ | keV | $10^{-2}$
norm | $10^{-6}$ | $4.56^{+1.73}_{-1.73}$
———– gaussian (O vii) ———–
$E$ | keV | $0.584^{+0.011}_{-0.010}$
$EW$ | keV | 0.072
$\sigma$ | keV | $10^{-3}$
norm | $10^{-5}$ | $2.2^{+0.4}_{-0.4}$
$\chi^{2}_{\nu}$ | $148.22/159=0.93$ |
Table 3: The best fit parameters to the XMM-Newton EPIC-pn spectrum of CP Pup,
with errors representing $90\%$ confidence limits. Values with no error were
kept frozen during the fitting process, and values with no units are
dimensionless. For completeness we also include the switch parameter that were
used throughout and determine whether the spectrum is computed via
interpolating or pre-computed tables. The abundances used correspond to the
’angr’ Table (https://heasarc.gsfc.nasa.gov/xanadu/xspec/manual/node116.html).

##  4 Discussion and Conclusions

In this section we discuss the interpretation of the nature of the bursts
detected in CP Pup. A connection is drawn in Section 4.1 with micronova, where
the details of the model are discussed as well as its limitations and
application to CP Pup. The recurrence timescale is determined from the model
and compared to that observed from ASAS-SN data. As mentioned in section 3.4,
the interpretation of the nature of the bursts, as well as the accretor mass
and accretion rate are here dependent on the assumption that CP Pup is a
magnetic system.

###  4.1 The CP Pup bursts as micronova eruptions?

The energy and shape of bursts shown by CP Pup resemble at least
phenomenologically the bursts reported in Scaringi et al. (2022c) and Schaefer
et al. (2022). In Scaringi et al. (2022c) the bursts are interpreted as
micronova bursts, somewhat analogous to Type I X-ray bursts in neutron stars.
Micronovae are hypothesised to be the result of a thermonuclear runaway effect
in a relatively small magnetically confined accretion column. The objects so
far reported to show these bursts are TV Col, EI UMa and ASASSN-19bh, with
burst energies ranging from $9\times 10^{37}$ erg to $1.2\times 10^{39}$ erg.
TV Col in particular has been observed at UV wavelengths during one of the
bursts by Szkody & Mateo (1984), displaying high ionisation helium and
nitrogen lines strengthening during the burst as well as clear P-Cygni
profiles suggesting outflow velocities in excess of $3500$km s-1. The bursts
in TV Col and EI UMa have been compared to those in V1223 Sgr (Hameury et al.,
2022). In Hameury et al. (2022) the bursts reported in V1223 Sgr are
interpreted by the magnetic-gating accretion instability model as observed in
other accreting WDs (Scaringi et al. (2017, 2022a); Littlefield et al.
(2022)), but draws no definite conclusion on the nature of bursts in other
IPs. Although magnetic gating may explain some observables for V1233 Sgr, it
may not explain the large outflows observed by Szkody & Mateo (1984) in TV
Col. The recurrent nova V2487 Oph has also been reported to display fast
bursts that appear phenomenologically similar (Schaefer et al. (2022)), but
their interpretation has been argued to be attributed to flares due to
magnetic recconection events within the accretion disc.

If these are indeed related to the micronova events, then we can attempt to
compute the burned/ejected mass during one of these events. The average
conversion rate of Hydrogen during the CNO cycle in a standard nova explosion
(Bode & Evans, 2008; José et al., 2020) is $\approx 10^{16}$ erg g-1. With
this assumption the observed energies convert to ejected masses of $>6.2\times
10^{-12}M_{\odot}$ and $>3.2\times 10^{-12}M_{\odot}$ for sectors 7 and 34
respectively, somewhat lower than those reported for TV Col and EI UMa in
Scaringi et al. (2022c) of $1.8-5.8\times 10^{-11}M_{\odot}$. We note that all
these estimates should be considered lower limits as no colour and bolometric
corrections have been applied. Hence the energies, and consequently ejected
mass would most likely be larger under this interpretation.

From the micronova model (Scaringi et al., 2022b), the recurrence of the
bursts is simply associated to the time it takes for a confined column of
material to reach thermonuclear runaway conditions. In the simplest scenario
the burst recurrence time is given by
$t_{\mathrm{rec}}=\frac{M_{\mathrm{ejected}}}{\dot{M_{\mathrm{acc}}}}$, where
$M_{\mathrm{ejected}}$ represents the ejected mass during the burst and
$\dot{M_{\mathrm{acc}}}$ the accretion rate of the system. The mass accretion
rate estimated from the fit to the XMM-Newton spectra is between $1\times
10^{-10}$ to $2\times 10^{-10}M_{\odot}yr^{-1}$. This assumes a WD mass of
$M_{WD}\sim 0.73M_{\odot}$${}^{+0.12}_{-0.11}$, as estimated in Section 3.4.
Hence the estimated recurrence time would be $t_{\mathrm{rec}}\approx 7-23$
days. For comparison the accretion rates quoted above from literature and the
one determined in this work are all compared in Table 4 alongside with the
corresponding recurrence times. Most of the accretion rates in Table 4 are
also derived from fits to X-ray spectra and in such case the
$T_{\mathrm{max}}$ and appropriate model are given as well. The table also
notes if the model has a constraint on the $T_{\mathrm{max}}$ parameter and
hence the accretion rate. This is judged by the constraint placed on the
$T_{\mathrm{max}}$ by the model boundaries of the parameter.

The long-term ASAS-SN light curve may also constrain the recurrence timescale
of the bursts. Figure 1 reveals an additional burst not observed by TESS after
the ones detected in Sector 7 and 34 around $\sim 3005$ BTJD. Along with the 2
bursts observed by TESS this suggests a recurrence timescale of $\sim 60$
days. Considering the entire long-term ASAS-SN light curve from $\sim 1020$
BTJD to $\sim 3020$ BTJD the recurrence time can be roughly estimated to be
$\sim 30-60$ days. The lower limit of the range is consistent with the
predicted recurrence timescale using the accretion rate from Section 3.4. We
do note that the data gaps could potentially make this an upper limit on the
recurrence time. Beyond the long-term ASAS-SN data it is difficult to estimate
if the recurrence holds.

Given the 1942 classical nova eruption, it is plausible to assume that up
until then the WD in CP Pup had accreted $\approx 10^{-5}M_{\odot}$ during the
preceding centuries, and that this fresh Hydrogen material would be spread
across the entire WD surface. On the other hand, from the ASAS-SN long-term
light curve it appears that fast burst events have been recurring in CP Pup
for at least $\sim$4 years. Assuming these events are indeed micronova
eruptions, the amount of mass used by the micronova events would then
constitute only $\sim$0.005% of the mass required to trigger the next
classical nova. It is of course also possible that micronova events have been
occurring since the 1942 nova explosion. In this case, the consumed mass from
micronovae would represent about 0.1% of the required mass to trigger the next
classical nova. It is unclear how the recurrence timescales of micronovae and
classical nova in the same system relate to each other, but we point out that
the two types of explosions do not necessarily need to be mutually exclusive.
The magnetic confinement criteria required may not be attained for long enough
time to build up enough material to trigger a micronova, allowing lateral
spreading of freshly accreted material. The exact conditions for this remain
unclear and it is possible that the bursts have not always been present in the
system, for example if the conditions for magnetically confined accretion are
fulfilled only temporarily. In such a case one or multiple micronova may be
triggered. Cessation of micronova events could happen if the magnetic
confinement of the accretion column is broken. Whether the accretion column
can stay magnetically confined is thought to depend on the combination of
multiple parameters, such as the strength of the magnetic field, the column
height and footprint area as described in detail in Scaringi et al. (2022b).

Assuming a micronova interpretation, the lower limit of the recurrence deduced
from the long term ASAS-SN light curve $t_{\mathrm{rec}}\approx 30$ days, the
expected ejected mass would convert to $\sim 3-6\times 10^{-11}M_{\odot}$,
consistent with those reported in Scaringi et al. (2022c). This however would
imply that the released energy in the TESS passband underestimates the
bolometric release of energy by a factor of 3 or more. It is however possible
that lower mass transfer rates could yield consistent burst recurrence times
and energies.

All the recurrence timescales in Table 4 are however either shorter or on the
lower boundary of the recurrence timescale deduced from the long term ASAS-SN
light curve. This might be due to the number of necessary conditions for
micronova to occur. For example the magnetic pole can change the size of its
footprint on the surface of the WD before enough material is accreted for a
micronova to be triggered. In such a case the recurrence timescale would
become longer. The recurrence timescales could also be underestimated due to
the WD mass. The same underestimation can be applicable in the recurrence
timescale using accretion rate from other papers, see Table 4. For the mass of
the WD to be better constrained however, a wider spectral coverage is
necessary, such as that of NuSTAR.

Table 4: Summary of accretion rate estimates for CP Pup from literature and the parameters used to determine it. The recurrent timescale range is computed for the ejected masses above and compared to the results in this work. Paper | Data | Model | $T_{\mathrm{max}}$ | Distance | $M_{WD}$ | $\dot{M}_{\mathrm{acc}}$ | $t_{\mathrm{rec}}$ | Constrained
---|---|---|---|---|---|---|---|---
|  |  | keV | pc | $M_{\odot}$ |  $\times 10^{-10}$ $M_{\odot}$yr-1 | days |
Orio et al. (2009) | XMM-Newton | APEC | 64 | 1600 | >1.1 | $\lesssim 1.6$ | 7.3 - 14 | No
Mason et al. (2013) | Chandra | VMCFLOW | $36.5^{+19.2}_{-16.3}$ | 1600 | 0.8${}^{+0.19}_{-0.23}$ | $4^{+3}_{-1}$ | 1.6 $-$ 6.9 | Yes
Selvelli & Gilmozzi (2019) | $-$ | Livio (1992) | $-$ | 795 $\pm$ 13 | 1.16 $\pm$ 0.2 | $6.0^{+6.6}_{-3.1}$ | 0.9 $-$ 7.8 | $-$
This work | XMM-Newton | CEMEKL | $29.3^{+9.4}_{-7.4}$ | 780 $\pm$ 11 |  $0.73$${}^{+0.12}_{-0.11}$ | 1 $-$ 2 | 7 $-$ 23 | Yes

Another important parameter of the micronova model is the fractional area on
which material is being accumulated. From the ejected masses
$M_{\mathrm{ejected}}$ the fractional area $f$ can be derived from Figure 1 in
Scaringi et al. (2022c). For $M_{\mathrm{ejected}}$ and WD mass quoted above
the expected $f$ is $\sim 2-9\times 10^{-7}$. This is particularly low for
triggering micronova and may point to the bombardment scenario suggested for
polars (Frank et al., 2002). In that case, inhomogeneous accretion of dense
blobs penetrate into the WD photosphere and radiate most of their energy in a
form of soft X-rays. In such a scenario in Frank et al. (2002), the accretion
rate is constrained by $\lesssim 1.6\times 10^{-11}$ $M_{\odot}$yr-1. This is
much lower than the estimated accretion rate for CP Pup and argues against the
micronova interpretation. The only way out from this potential problem with
the model is that the WD mass has been underestimated in the X-ray spectral
analysis.

Further photometric, polarimetric and high-time resolution spectroscopic
follow-up observations, of the peculiar behaviour displayed by CP Pup will
provide additional constraints to unravel its true nature. Specifically the
lack of consistent results when it comes to CP Pup, coupled with the lack of a
coherent model to explain its behaviour, may warrant a dedicated observing
campaign.

## Acknowledgements

We would like to thank the anonymous referee for the useful and insightful
comments that have improved this manuscript. 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.
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. MV acknowledges
the support of the Science and Technology Facilities Council (STFC)
studentship ST/W507428/1. SS is supported by STFC grant ST/T000244/1 and
ST/X001075/1. SH acknowledges support from STFC through the studentship
ST/V506643/1. KI acknowledges support from Polish National Science Center
grant 2021/40/C/ST9/00186. DdM acknowledges financial support from ASI-INAF
agreement and INAF Astrofund-2022 grants.

## Data Availability

The TESS data used in the analysis of this work is available on the MAST
webpage https://mast.stsci.edu/portal/Mashup/Clients/Mast/Portal.html. The
ASAS-SN data (Shappee et al., 2014; Kochanek et al., 2017) used for the
calibration of TESS data is available on the ASAS-SN webpage https://asas-
sn.osu.edu/.

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