Journal "Peremennye Zvezdy"
Peremennye Zvezdy (Variable Stars) 30, No. 4, 2010
Received 29 September; accepted 5 October.
Article in PDF
Sternberg Astronomical Institute, University Ave., 13,
Moscow, 199992 Russia
- Astronomical Institute of the Slovak Academy of Sciences,
Tatranska Lomnica, 05960 Slovakia
- Sternwarte Sonneberg, Sternwartestrasse, 32, Sonneberg, D-96515
- Special Astrophysical Observatory of the RAS, Nizhnij Arkhyz,
Karachai-Cherkesia, 369167 Russia
V445 Pup was a peculiar nova having no hydrogen spectral lines
in the outburst. The spectrum contained strong emission lines of carbon,
oxygen, calcium, sodium and iron.
We have performed digital processing of photographic images of the
V445 Pup progenitor using astronomical plate archives. The brightness
of the progenitor in the B band was 14.3 mag. It was
found to be a periodic variable star, its most probable period being
0d.650654±0d.000011. The light curve shape
the progenitor was a common-envelope binary having a spot
on the surface and variable surface brightness. The spectral energy
distribution of the progenitor between 0.44 and 2.2 µm was similar
to that of an A0V type star.
After the explosion in 2001, the dust was formed in the ejecta,
and the star became a strong infrared source. This resulted in the star's
fading below 20m in the V band. Our CCD BVR observations
acquired between 2003 and 2009 suggest that the dust absorption minimum
finished in 2004, and the remnant reappeared at the level of 18m.5
The dust dispersed but a star-like object was absent in frames
taken in the K band with the VLT adaptive optics. Only expanding ejecta of
the explosion were seen in these frames till March 2007.
No reddened A0V type star reappeared in the spectral energy distribution.
The explosion of V445 Pup in 2000 was a helium flash on the surface of
CO-type white dwarf. Taking into account the results of modern dynamic
calculations, we discuss the possibility of white-dwarf core detonation
triggered by the helium flash and the observational evidence for it.
Additionally, the common envelope of the system was lost in the explosion.
Destructions in the system and mass loss from its components exclude the
future SN Ia scenario for V445 Pup.
The outburst of V445 Pup was discovered on 30 December 2000
by Kanatsu (Kato & Kanatsu 2000). The earliest observation of
V445 Pup in the outburst
dated 19 November 2000 was found in ASAS archive. At that time,
the brightness of the star was 8m
.8. The brightness
maximum of 8m
.46 in the V
band was reached on 29 November 2000.
The first spectroscopic observations in the outburst by Wagner et al. (2001)
showed that the Balmer emission and He I lines typical for classical novae
were not present in the spectrum of V445 Pup. The spectra were
dominated by emission lines arising from Fe II, Ca II, C II, Na I,
O I. Line widths corresponded to an expansion velocity
of about 1000 km s-1
. The ejecta produced during the outburst
allow us to consider V445 Pup as a nova.
The nature of classical novae
is known to be a thermonuclear explosion of hydrogen on the surface
of a white dwarf in a semidetached binary system. Hydrogen
accumulates on the surface of the white dwarf due to accretion
from a donor, usually a red dwarf. As a result of hydrogen
explosion, strong Balmer lines are observed in the spectra of
classical novae. Ashok & Banerjee (2003) suggested that V445 Pup is
the unique helium nova predicted theoretically by
Kato et al (1989) and Iben & Tutukov (1994) who considered the case of a
degenerate white dwarf accreting helium from a helium-reach donor.
Note that a subclass of classical novae called helium-nitrogen (He/N)
novae was introduced by Williams (1992). These novae have spectra
with strong Balmer lines, and they also have He I and He II lines.
CNO elements seen in the spectra were mixed by accumulated hydrogen
envelope from the surface of the white dwarf through a dredge-up
mechanism (see e.g. Glasner & Livne 2002).
The case of helium nova suggests that the donor is a nucleus of an
evolved star that previously lost its hydrogen envelope due to accretion.
In the outburst of V445 Pup, the decay of brightness by 1m
continued gradually for 164 days and followed by a small rebrightening between
12 May and 21 June 2001. The last observation of V445 Pup in the
outburst was registered on 11 July 2001 at visual magnitude 11.5.
Then the star faded rapidly and was not seen in August 2001.
On 4 October 2001, no object brighter than V
.5 was found by Henden et al. (2001)
at the position of V445 Pup. They remark: "The object is
evidently shrouded in a thick and dense carbon dust shell, in view of
the apparent over-abundance of carbon in ejecta previously observed in
infrared and optical spectra". Lynch et al. (2001) detected
the infrared radiation in the 3–14 µm range just 1 month after the
object had been discovered. The spectrum revealed
smooth and featureless continuum, which they treated as a thermal
emission of dust with the temperatures ranged between 280 and 1300 K.
They suggest that this dust was a product of previous outbursts, at
least in part.
The detailed spectral evolution of V445 Pup in the outburst was studied by
Iijima & Nakanishi (2008). They acquired both high- and medium-resolution
spectra for the
optical wavelengths 3900–7000 Å. They confirmed the absence of hydrogen
lines and noted unusually strong emissions of carbon ions. Some
metal lines had P Cyg-type profiles with absorption components blue-shifted
roughly by –500 km s-1
; this velocity was assumed to be the
outflow velocity. The cited authors measured large radial velocity of
V445 Pup, +224±8 km s-1
, which suggested that the object
belonged to the old disk population. The distance was estimated using the
interstellar NaI D1
absorption lines to be
3.5 < d
< 6.5 kpc; the reddening is E(B–V)
Lynch et al. (2004) reported that in January, 2004 the object had faded to
fainter than J
, so that they could not take its spectrum
in the visible range. In the infrared, they detected only two
He I lines at 1.0830 and 2.0581 µm, both showing doubled profiles
due to bipolar outflow. The very red continuum was detected only
at λ ≥ 1.5 µm.
It was produced by emission of hot dust.
Woudt et al. (2009) published the results of post-outburst JHK
adaptive optics imaging in the K
band, and optical-range spectroscopy
of V445 Pup. They discovered an expanding and narrowly confined bipolar
shell, the outflow characterized by large velocity of
6720±650 km s-1
. Some knots were moving with larger
8450±570 km s-1
. They derived an expansion parallax distance
of 8.2±0.5 kpc. They noted that the expansion velocity measured
by Iijima & Nakanishi (2008) with the high resolution spectra in outburst
was only 500 km s-1
. Such a big difference may be due to strong
collimation of bipolar ejection located just in the plane of the sky and
inclined to this plane only by
5°.8–3°.7 (Woudt et al. 2009). The authors assume that the small
inclination angle may confirm the presence of an orthogonal dust structure
closely aligned to the line of sight and causing the strong extinction
observed after the outburst.
In their spatially resolved optical spectrum obtained with VLT in 2006
January in the 4465–7634 Å range, only the emission lines of [O I],
[O II], [O III] and He I were seen, but not the continuum.
The presence of a bright progenitor of V445 Pup having a visual magnitude
of 13.1 was first noted by Platais et al. (2001). Its absolute proper motion
was small, µ = 0".008±0".004. With the distance derived by
Woudt et al. (2009), the luminosity of the progenitor proved to be very large,
= 4.34±0.36 which is consistent with the
absolute bolometric magnitude value Mbol
Woudt et al. (2009) note
that the derived luminosity suggests that V445 Pup probably contains
a massive white dwarf accreting at high rate from a helium star companion.
But they did not exclude that the companion was also a bright star.
Liller (2001) reminded of three hydrogen-deficient cataclysmic variables,
CR Boo, CP Eri, and V803 Cen, all of them being hot blue objects showing
no hydrogen, but revealing He I emission lines.
The absolute magnitude of –6.1 is unprecedented high
for a cataclysmic variable, making us to think about the nature of
The observations of the light curve in the outburst and light curve
modeling by Kato et al (2008) reveal that the CO white dwarf in V445 Pup
is very massive and close to the Chandrasekhar mass limit
≥ 1.35 MSun
a half of the accreted matter remained on the
white dwarf after the outburst. Therefore, V445 Pup was considered
a strong candidate for a type Ia supernova progenitor. P. Woudt and
D. Steeghs called V445 Pup a "ticking stellar time bomb"
in the ESO Science Press Release 0943. Taking into account the observations
with adaptive optics by Woudt et al. (2009) which show only
spatially resolved products of eruption but no stellar
component, it is hard to maintain the concept that the mass of the system
increases. The question is what mass of the components was left after the
The scenario for V445 Pup may be quite different. Recent dynamic 3D simulations
by Guillochon et al. (2010) discovered a new mechanism for the detonation of
a core of a sub-Chandrasekhar CO white dwarf (with a mass lower than
in the system with a pure He white dwarf or a He/CO hybrid secondary.
Fink et al. (2010) found that secondary core detonations were triggered
for all of the simulated models ranging in core mass from 0.810 up to
with corresponding helium shell masses from 0.126 down to
. In that paper, the double detonation scenario
a potential explanation for type Ia supernovae. But the destruction
of the CO white dwarf means that V445 Pup, after its outburst in 2000,
will not be a type Ia supernova progenitor. It is of
great interest if the narrowly confined bipolar cones observed by
Woudt et al. (2009) are debris of the detonated white dwarf. The overabundant
carbon in the outburst will also be an evidence for CO white dwarf
Fortunately, there are many photographic images of V445 Pup in the world
astronomical plate collections suitable for resolving the puzzle of the
progenitor. Woudt et al. (2009) verified the plate archives at the
Harvard-Smithsonian Center for Astrophysics (USA) and found no prior outbursts
The progenitor of V445 Pup was identified on many plates at approximately
constant brightness (from visual comparison with surrounding stars).
We found many plates in archives of the Sternberg Astronomical Institute (SAI)
of the Moscow University (Russia) and in archives of the Sonneberg Observatory
(Germany). Two of us (V.P.G. and S.Yu.Sh.) performed eye estimates of
V445 Pup independently, V.P.G. for Moscow plates, and S.Yu.Sh. for
Sonneberg plates. The two sets showed similar behavior and
marginal variability. But unexpectedly, the preliminary frequency
analysis revealed the same periodicity in both sets with the period
.650653, coinciding to the 6th significant digit. Both
light curves were of low quality. Therefore we decided to digitize the
images of V445 Pup and to perform digital processing.
In the Moscow SAI plate collection, we found 51 plates with images of V445 Pup
taken with the SAI Crimean Station 40-cm f/4 astrograph and dated
between 15 November 1969 and 4 November 1989. AGFA ASTRO and ORWO ZU-2
photographic plates, produced in the former GDR and
having high sensitivity in blue light, were used, the exposure times
were 45 minutes. The geographic position of the SAI Crimean Station
+44°43´42". The declination of V445 Pup is
about –26°. This means that the highest altitude of the star
above the horizon is 19°. Observations were limited to a 3-hour visibility
time around this point. Photographic plates were mostly centered at
τ CMa, they cover an area of 10°×10°. The region
of about 20´×20´
centered at V445 Pup was digitized for each plate using the SAI
CREO EverSmart Supreme scanner. CREO scanner frames are in the TIFF format.
We found 56 measurable images of V445 Pup on the plates of
the Sonneberg Observatory
collection dated between 19 March 1984 and 17 January 1991.
These plates were taken with the 40-cm f/4 astrograph having optics basically
similar to that of the SAI Crimean Station astrograph. Also, plates of basically
the same type produced in the former GDR were used, so all our
photographic material is very uniform. The Sonneberg Observatory has the
geographic position 0h
+50°22´39", it is located
about 5° to the north by latitude compared to the SAI Crimean Station.
Thus, the star rises only to
14° for this geographic point, and its visibility time is less
than that for the Crimean Station. Sky images are evidently affected with
variable atmospheric extinction across the plate field. In these plates,
the star is located near the center. The typical exposure time for
these plates is 20 minutes. Images of V445 Pup were digitized using the
Fuji FinePix F10 CCD camera and an ordinary biconvex lens. Frames
made with this camera are in the JPEG format.
After several experiments, we found that this method of digitizing gave the
quality of measurements near the field center as good as that of a
scanner. To increase the S/N ratio, we co-added several subsequent frames
in night series. For this purpose, frames were put together by matching two
stars, with the needed field rotation. The combined exposures
of co-added frames were between 40 and 60 minutes.
The total number measurements for the star, including co-added ones, is 31.
Additionally, we measured all the Internet-accessible Digital Sky Survey
images of V445 Pup in B,R,
bands and used 2MASS JHK
magnitudes to study the spectral energy distribution of the progenitor.
All the frames were processed in the Windows BITMAP format. Extraction
of images was made using the aperture method with star-profile
correction; the WinPG software developed by V.P.G. was utilized.
Special software was written by V.P.G. to approximate the
characteristic curves with an n
th-degree polynomial, with graphical output.
Practically, the approximations with n
= 1 or 2 were optimal.
The total number of comparison stars used to build a characteristic
curve was 23; a few stars with the largest deviations were
eliminated from calculations, and the characteristic curves were
re-calculated in such cases. The r.m.s. deviation of comparison
stars from the polynomial fit was formally taken for the uncertainty
of V445 Pup measurement.
The finding chart of V445 Pup and comparison stars. This is a copy of a
digitized image obtained with the UKST Schmidt telescope on 4 April 1980
on IIIaJ emulsion with a GG 395 filter. The progenitor is indicated as
"var". V magnitudes, colour indices of marked stars, and
(in units of thousandths of a magnitude) are given in Table 1.
Carrying out oure photographic measurements, we used the CCD
BVRCIC standard sequence
in the vicinity of V445 Pup published by A. Henden for VSNET.
We present the finding chart of the progenitor in Fig. 1.
The standard stars chosen by us
are also marked in this Figure. We give Henden's magnitudes and their
uncertainties for the chosen comparison stars in Table 1 because they
are no longer accessible at the VSNET address.
The Moscow archive observations of the V445 Pup progenitor are presented
in Table 2; the Sonneberg ones, in Table 3; and those from digitized
sky surveys are collected in Table 4.
We performed our observations of the V445 Pup remnant
between 31 March 2003 and 20 October 2009. These observations were acquired
in the Special Astrophysical Observatory (SAO), with the 1-m Zeiss reflector
and CCD UBVRCIC photometer equipped
with an EEV 42-40 CCD chip. The geographic
position of SAO is 2h45m46s
+43°39´12". The highest
altitude of the star over the horizon is 20°. This object
is difficult for observations and needs good sky transparency and
seeing. Some constructions of the 6-m telescope dome located to the
south of the 1-m reflector humper observations of objects with
such a southern declination. Additionally, a
part of our observations were obtained with the SAI Crimean Station's 60-cm
reflector and UBVRJIJ photometer with
the Princeton Instruments
VersArray CCD. Both devices are cooled with liquid nitrogen to a temperature
stabilized at –130°C, allowing to record signals from very
faint astrophysical objects. The frames were reduced in the FITS format.
The extraction of images was made using the same aperture method with
star-profile correction, the WinFITS software by V.P.G. was
utilized. Our CCD observations are presented in Table 5.
The light curve of V445 Pup in the B
band excluding the outburst
is shown in Fig. 2.
The variability of the progenitor is evident, and its full amplitude exceeds
the mean uncertainty of the observations more than thrice.
The pre- and post-outburst light curve of V445 Pup in the B band.
The start and the end of the 2000–2001 outburst are indicated with a
double vertical line.
The frequency analysis of the progenitor observations was performed
using two independent methods: (1) the discrete Fourier transform for
arbitrarily distributed time series (Deeming 1975), and (2) the phase dispersion
minimization (PDM) method (Lafler & Kinman 1965).
Implementations of these methods
are provided by the EFFECT code developed by V.P.G. We analysed the combined
time series including Moscow and Sonneberg photographic observations.
The periodograms are shown in Fig. 3 a-c.
The panels (a) and (b) of this figure present the amplitude spectrum and the
spectral window for this series.
We estimate significance levels for peaks of the amplitude
spectrum using the empirical method suggested by Terebizh (1992). This
method is based on a statistical analysis of simulated chaotic series
generated by mixing the original series. In the chaotic series,
each Julian date gets an accidental magnitude chosen from the same
original series and, as a result, the chaotic series includes the same
magnitudes and times. When we compute the amplitude spectrum of the
chaotic series, we make more than a million of accidental light-curve
with arbitrary periods and estimate their amplitudes. The software provides
the analysis of the cumulative probability distribution function for amplitudes
in the spectrum of the chaotic series. The amplitude levels corresponding
to cumulative probabilities of 90, 99, 99.9 and 99.99 percent for the
chaotic series are plotted in Fig. 3a as straight lines.
Periodograms of the V445 Pup progenitor.
(a): The Deeming amplitude spectrum
in the (10–0.3)-day period range. The parameter is the half-amplitude.
(b): The spectral window of the same time series in the (1000–0.3)-day period
range. The parameter is the half-amplitude. (c): The Lafler–Kinman periodogram.
The parameter is θ-1, θ being the normalized
sum of squared magnitude differences between each two subsequent points
of the phased light curve calculated with a trial period.
The presence of strong peaks in the amplitude spectrum of the original series
exceeding 99.99 percent amplitude level of the chaotic series means
that the probability of casual appearanse of these peaks is less than 0.01
percent. The progenitor of V445 Pup was evidently a periodic variable star.
The multiplicity of peaks means that we have multiple solution
for periodicity with the Moscow and Sonneberg series.
The spectral window (Fig. 3b) demonstrates the periodicity in time
discontinuities in our series amounting to the sidereal day
.997262) and of
, where n
= 2, 3, 4....
The amplitudes of these peaks decrease when the period decreases because of
increasing phase window. The phase window for Psd
0.2. Thus, for each peak in window spectrum, we have a pair of symmetrically
located alias peaks in the amplitude spectrum.
The light curves corresponding to this pair of peaks have reverse
phase count, so they look as mirror-reflection ones.
The list of periods and frequencies of aliases is given in Table 6.
One can see that the dominating peaks are sidereal-day-related. The
formula of corresponding interdependence is given for each peak in the
last column, `Remark', of Table 6. For f0
, we chose the
lowest-frequency wave with the highest amplitude. Peaks of a lunar month
(29.5 days) and
of about a year (363 days) in the spectrum of the window are also present,
which are responsible for combs of small peaks located around
sidereal-day-related peaks in the star spectrum.
The light curves plotted for the periods determined with
Deeming (1975) method and presented in Table 6. The elements used
to calculate phases are given below each curve.
The light curves corresponding to all alias periods are given in Fig. 4.
These are single-wave curves. The light curves are approximately sinusoidal,
the full amplitude of the sine wave is about 0m
the Deeming method
reveals the highest-amplitude light curves for two periods,
.00009 and 2d
.00011 with equal half amplitudes
.22. The scatter of all the light curves
reveals essential intrinsic variability. A few points do evidently
contradict the sinusoidal solution. We verified these points and confirmed
their Julian dates and magnitudes. These light curves may represent the
case of reflection effect on the surface of a secondary star due to heating of
a part of its surface by the X-ray or short-wavelength radiation coming
from the primary star.
However, no X-ray source was associated with V445 Pup before its outburst.
In principle, such light curves may arise due to a large hot spot on the surface
of a star. FK Com-type stars may be examples of a rotating star
having a hot spot on the surface. These stars are considered to be close
binary systems with a common envelope.
The light curves plotted for double periods
determined with the Deeming (1975) method and given in Table 6. The elements used
to calculate phases are given below each curve.
Double-wave light curves for periods found by the Deeming method presented
in Table 6 were also calculated and are shown in Fig. 5. Double-wave light curves
are exhibited by W UMa-type binaries, these are also stars with
common envelopes. In most observational cases of photometry without
additional spectroscopic information, like radial velocity curves or
double lines in the spectra, we can distinguish between single-wave and double-wave
orbital periods only by differing alternate minima.
The difference of minima depths appears due to difference of surface
brightness of the components. Our observations do not show alternate minima
of different depth. Unfortunately, the accuracy of
photographic observations is insufficient to make a reliable choice between
single- and double-wave curves. However, taking into account that this system
contains a massive accreting CO white dwarf (Kato et al. 2008), we think that
a W UMa light curve is not an acceptable solution because components
of such a system should have different brightness.
The light curve for the best single-wave period found with
Lafler & Kinman (1965) method.
The Lafler–Kinman (L–K) method reveals periods not exceeding one day
as preferable: 0d
their double-wave aliases 1d
.359423±0.000048. The 0.650654-day period
determined with the L–K method differs essentially from that
determined by the Deeming method because the light curve plotted with
the PDM solution shows a local detail at the phases between 0.9 and
0.1 that looks like a shallow eclipse (Fig. 6). However, additional
photographic material is needed to verify if this detail is real.
Certainly, a detail like a shallow eclipse cannot be revealed
in a light curve with the Deeming method.
The double-wave light curves found by the L–K method seem more irregular
and asymmetric. Additionally, the light curve with the period of
.000011 is the most symmetric one and has the lowest
dispersion, so we choose it as the best solution for the present time.
Our investigation shows that one can find a single final solution
for the orbital period of V445 Pup only using
observations taken at different geographic longitudes, thus
increasing the observational phase window of the sidereal-day period,
. Fortunately, there are
enough plates in the world plate collections to get such a solution.
It is known that Harvard plates were taken at the station located
in the southern hemisphere, and this series would have the widest window.
We hope that the
digitally reduced Harvard plates and Japanese pre-outburst photographic
observations, along with our data, will provide a true final solution.
We have remeasured B, R
frames of Digital
Sky Surveys for the V445 Pup
progenitor using Henden's photometric comparison-star sequence.
The infrared JHK
observations of the progenitor were taken from the
Additionally, we compared these observations to our optical
observations of the
remnant based on the same comparison-star sequence.
The brightness of the remnant was found to be variable
in the magnitude ranges of 19.3–20.8 in the B
in the V
band; 17.5–19.4 in the RC
band; and 17.7–20.0 in the IC
In Fig. 7, we show the spectral energy distributions (SEDs) of the progenitor
remnant for three different dates: December, 2003, in minimum light after
the outburst when its brightness was V
≈ 20.1 (r1);
January, 2005, in the peak of rebrightening at V
≈ 18.6 (r2); and in
November, 2008, after the rebrightening at V
≈ 19.2 (r3). To plot
the distributions for the remnant, we used JHK
by Woudt et al. (2009) from their Table 2, the closest to our dates.
All distributions were dereddened;
those for V445 Pup were dereddened with the colour excess of E(B-V)
(Iijima & Nakanishi 2008), based on calibration of the interstellar
Na I doublet equivalent width measured in high-resolution spectra.
The spectral energy distributions determined from photometry. The upper solid
curve, marked p, is the energy distribution of the V445 Pup progenitor,
the bars of the B point correspond to the variability amplitude.
The JHK data are from the 2MASS survey.
Interstellar extinction corresponding to E(B - V) = 0.51 was
taken into account.
The gray curve marked A0 is the energy distribution of an A0V type star,
plotted for comparison. The three lower curves are energy distributions of the
V445 Pup remnant measured at different time: (r1) December 2003;
(r2) January 2005; (r3) November 2008. JHK observations are from Woudt
et al. (2009).
We compared our pre-outburst SED with that previously published by Ashok
& Banergee (2003) who find consistency of the SED with an accretion disk
model. In the JHK
bands, they used the same 2MASS data,
but in the optical B
bands, they used the USNO A2.0 catalog. The USNO A2.0 R
well with our photometry based on Henden's comparison stars.
But the USNO B
value is a half-magnitude brighter than the mean
magnitude from our archive observations and is outside the variability
range. In Fig. 7 by Ashok & Banergee (2003), the B
located above the straight line fitting the other points. It is clear that
it is this deviation in the B
band that caused the choice if accretion
disk model. Additionally, Ashok & Banergee (2003) used a smaller
value of 0m
.25, derived from the extinction maps by
Neckel et al. (1980) and Wooden's (1970) UBV
photometry of stars in the
5° × 5° field around V445 Pup. Woudt et al. (2009) reanalyzed the
problem of interstellar extinction based on published photometry of open
clusters and ascertain that E(B-V)
= 0.51 can
be the lower limit of Galactic reddening, and the actual value can be
between 0.51 and 0.68. They also include circumstellar reddening in
their estimates of the star's absolute magnitude and luminosity.
Our SED of the V445 Pup progenitor fits that for an A0V
star well (A0 in Fig. 7). If we exclude a possible small excess visible in the
infrared bands, the spectral type may by somewhat earlier.
Probably, this excess is the
radiation of dust detected by Lynch et al. (2001) in the outburst.
But this was not the dust of the surrounding thick disk,
assumed to explain the obscuration of the whole system after the outburst.
No essential infrared excess is seen in the SED of the progenitor
to cause additional circumstellar light absorption and reddening.
It is notable that the energy distribution of the progenitor is that
of a single star: we do not see
any contribution of a cool donor or hot radiation from an accretion disk.
Moreover, Kato et al. (2008) demonstrated that the progenitor was too bright
to be an accretion disk at the known distance (3 < d
< 6.5 kpc)
known Galactic interstellar absorption (AV
even for the largest
reasonable mass accretion rate (10-6MSun
They note that "it is very unlikely that an accretion disk mainly
contributes to the pre-outburst luminosity". It is evident from our
photographic photometry that there were no strong emission lines in
the SED of the progenitor which might disturb the plain continuum in such
a way as it is observed in the SED of the remnant. But we could not establish any
contribution of Balmer absorption or discontinuety of an A0 star due
to absence of U
-band observations of the progenitor.
We determined the mean B
magnitude for V445 Pup
to be 14m
Taking into account the A0V energy distribution and the Galactic
reddening of E(B–V)
= 0.51 (Iijima & Nakanishi 2008),
we find (B–V)0
which corresponds to the star's reddened values
.51 and mean V
With the interstellar reddening of 0m
and the distance of 5 kpc, we have
.28. With the bolometric
correction of –0.25 determined for an A0 type star by Straizys (1982),
we have log L/LSun
= 2.54. This value is essentially lower than
that of 4.34 derived by Woudt et al. (2009). Even with their distance of
8.2 kpc, we find MV
These luminosity estimates are still too bright to fit models of an
accretion disk with a high mass transfer rate.
We think that the luminosity of the progenitor was overestimated
by Woudt et al. (2009), mostly due to inferring additional
circumstellar absorption by an equatorial disk/torus. V445 Pup is an
old-population object from its large radial velocity, and its
location in the CM diagram is about 1m
.6 above the
RR Lyrae horizontal branch gap of the globular cluster population.
On the other hand, the energy distributions of the V445 Pup remnant are
composite. In the H
bands, this is mostly continuum
from cool dust.
Woudt et al. (2009) noted two emissions of He I in the 0.9–2.5 µm
spectrum taken in January 2004; one of them, at 1.0820 µm, is
very strong. In their spectrum, the continuum is present only beginning
with 1.5 µm.
Woudt et al. (2009) wrote that V445 Pup was highly reddened in the
bands more than
8 years after the outburst. However, Fig. 2 in their paper shows that
all the radiation in K
band does not belong to an A0-type star highly
reddened by the ejected dust envelope. It belongs to the dust envelope itself
which consists of two jet-like elongated lobes. No reddened stellar object
(such as an A0-type star) is located between the lobes. The
bands contained mostly the thermal component of
= 250 K, and
the intensity of this thermal component gradually declined. The J
contained the He I λ2.0581 µm recombination line with an
equivalent width of approximately –180 Å. The behaviour
of the flux in the J
band was different, it did not decline. The
figure shows some brightening of the remnant in the J
simultaneously with the decline in other infrared bands, possibly due to
appearance and strengthening of He I emission.
In the optical and near-infrared bands, the continuum of the thermal dust
component is not found, and we see the recombination lines.
In the 4500–7600 Å spectrum taken with VLT (Woudt et al., 2009),
four strong emission lines: [O III] 4959 Å, He I 5875 Å,
He I 7065 Å, and
the blend [O II] 7319/7330 Å containing four components are present.
The He I and [O III] lines are spatially resolved.
Therefore, our BVRCIC
magnitudes and colours do not reflect temperature information, they
are formed only by line radiation accidentally hitting the photometric
bands. The [O III] line is at opposite edges of the B
passbands, and the other lines are in the RC
show brightening of all emission lines in the time range from
JD 2453004 (r1 in Fig. 7) to JD 2453386 (r2) and subsequent weakening
from JD 2453386 (r2) to JD 4254777 (r3).
The changes in the SED due to the outburst in 2000 are radical.
We do not see such a blue and bright star as before the
outburst. Most investigators, from Henden et al. (2001) to Woudt et al. (2009),
write that the binary
is obscured by the carbon dust shell ejected in the outburst.
However, we do not see such a bright and blue but reddened star
either in the deep IR spectra or in the energy distributions.
Does any radiation come from under the shell? But there is the ionized gas,
ions of [O II] and [O III]. What is then the source of ionization and
The light curve of V445 Pup in the
V band including the 2000 outburst. 1: The primary
light maximum; 2: rebrightening before the deep fading; 3: rebrightening
in the low state. The gray curve is a dip in the light curve due to carbon dust
obscuration. The dashed curve is a typical decay without dust obscuration.
The bottom light curve is the post-outburst one in the RC band.
The light curve of V445 Pup in the optical V
band presented in Fig. 8
looks like a typical light curve of a dust nova having a prolonged
dip due to the dust obscuration. Good examples
of contemporary dust novae are V1419 Aql, V705 Cas, and V475 Sct. The light
curves of these stars can be found in the AAVSO database and in
our database at
For V445 Pup, such a dip
continued approximately since JD ~ 2452200 till JD ~ 2453500.
One may assume that, as in other dust novae, the end of this dip and
the subsequent star rebrightening are associated with the dispersion
of dust and improvement of visibility of the binary system's remnant.
The dispersion of dust leads to temporal strengthening of emission lines in
the expanding gaseous envelope. But after that, the tendency of weakening
for light from the gaseous envelope continued due to its expansion.
This process is typical of classical dust novae in the late stages of
There is a consensus that SNe Ia result from the explosion of a carbon-oxygen
(CO) white dwarf in a binary system (Meng & Yang, 2010). The cited paper
contains a good review of evolutionary ways leading to the SN Ia explosion,
with corresponding references to the literature.
The first theoretical studies of helium flashes on the surface of a CO-type white
dwarf were performed in the context of SN Ia scenarios. Such dwarfs were
members of common-envelope binary systems containing a low-mass red
giant with a helium core and mass transfer (Hachisu et al., 1989).
the mass transfer rate, helium burning on the surface of a white dwarf may be
stable or unstable, with flashes. As a result of helium
burning on the surface of a CO white dwarf, its mass gradually increases,
exceeds the Chandrasekhar limit for the white dwarf mass, and therefore
the CO dwarf explodes as a thermonuclear runaway.
Hachisu et al. (1989) assume that the progenitor may be observed as a
symbiotic star (if the donor is a subgiant) or as an A–F giant
(if the donor is a red giant). In the last case, the mass-receiving component
(mass gainer) has the photospheric radius comparable with that of the donor.
Kato et al. (1989) discuss models of a helium nova outburst in a compact binary
with direct accretion of helium from a helium star companion which
fills its inner critical Roche lobe.
In those models, helium shell burning ignites when
the density and the temperature in the helium envelope reach critical values.
The envelope expands and forms dense stellar wind when it exceeds the
critical Roche lobe. When the helium burning ends, the star returns to the
same point in the luminosity–temperature (L–T
) diagram which was
occupied by the progenitor. Helium flashes on the CO white dwarf surface were
not considered as a direct cause of SN explosion.
The appearance of a real helium nova, V445 Pup, showed that these
ideas were simplified: dust formation and asymmetry of ejecta
were not predicted. Dust forming and dissipation were not
considered in the paper by Kato et al. (2008) written after the V445 Pup outburst,
they considered only free-free-emission-dominated light curves based
solely on the optically thick wind theory.
Recent dynamic calculations by Guillochon et al. (2010) reveal
the mechanism of helium ignition on the surface of a white dwarf.
With the high accretion rates on the white dwarf surface that
range from 10-5
to 10-3 MSun
the hot helium torus forms, and the accretion stream fails to impact
the white dwarf surface. The mass of the torus accumulated during
such an event varies between 0.05 and 0.14 MSun
Due to velocity difference between stream and torus,
large-scale Kelvin-Helmholtz instabilities arise along the interface
between the two regions which "eventually grow dense knots of material
that periodically strike the surface of the primary, adiabatically
compressing the underlying helium torus". The temperature of compressed
material increases above a critical temperature (2·10-9
and then triple-α reactions lead to detonation of the primary's
helium envelope. Moreover, Guillochon et al. (2010) show that shock waves
of the detonated envelope tend to concentrate in the focal points within
the CO core of the primary. The CO core of a white dwarf may detonate itself.
This kind of detonation of the CO core does not lead to a SN Ia, but these
transient events resemble dim type I SNe.
The cited authors used smoothed particle hydrodynamics (SPH) simulations in their
calculations. In such an event, the stability of the system depends on
the structure of primary and secondary components. However, the particle
approximation cannot fully resolve the stream and the helium shell
because they have small masses as compared to the masses of the two
components of the binary. To increase the resolution, the authors used a hybrid
approach that combined both SPH- and grid-based methods. Grid-based
models predict that the donor will survive in such an explosion.
Principally, Guillochon et al. (2010) presented models of merging of a helium
star and a CO white dwarf, i.e. a dynamically unstable system in the last
orbits before merging. In such an event, the orbital period of the system
may be unstable due to angular momentum losses. Theoretical and observational
studies of the orbital period stability for a long time interval before
the merging event are needed to compare the V445 Pup explosion with the merging
Fink et al. (2010) use dynamic calculations to investigate whether an
assumed surface helium detonation is capable of triggering a subsequent
detonation of the CO white dwarf core. Their
calculations were performed for a range of masses of white
dwarf cores and masses of helium shells on the white dwarfs' surface.
The authors find that secondary core detonations are triggered for all
of the simulated models, ranging in core mass from 0.810 up to 1.385
with corresponding shell masses from 0.126 down to
The result of the calculations was the following: "as soon as a detonation
triggers in a helium shell covering a CO-type white dwarf, a subsequent core
detonation is virtually inevitable". For the case of V445 Pup explosion
which began with a helium flash on the surface of the massive CO white
dwarf, the white dwarf was destroyed partly or totally. The infrared
bipolar nebulosity discovered using adaptive optics may be the
debris of the destroyed white dwarf. But V445 Pup
is not a dim SN Ia. Having an absolute magnitude MV
-5.9 and -7.0 in the light maximum, it did not reach the absolute magnitude
level ranging between -13 and -14 for ultra-faint supernovae
(Smith et al. 2009), or so called "supernova impostors".
What happened in the 2000 explosion event of V445 Pup? Our and other studies
established that the progenitor was a hot binary system without a dense
dust circumstellar disk. We do not see absorption or emission of
such a disk in the energy distribution of the progenitor. The assumption
of a cataclysmic binary with the accretion disk is not probable because of high
progenitor's luminosity and an orbital light curve unusual for cataclysmic
variables. Indeed, a highly inclined cataclysmic variable has
a hump in its light curves caused by visibility of a hot spot on
the edge of the accretion disk; the hump is visible only during a half
of the orbital period. We assumed that the
progenitor was a common envelope binary with a luminous helium donor.
The periodic variability suggests a single star with a hot spot on
its surface. We know sdB + M type binaries with bright spots on the
surfaces of their cool components due to irradiation of ultraviolet flux
sdB star. However, it is not our case because we do not see a cool companion:
the companion was hot. The hypothesis of irradiation of X-rays in a
binary with a compact companion is rejected due to absence of X-ray
source associated with the progenitor.
Many observers detected dust formation during the outburst
of V445 Pup and the strong dust
absorption after the outburst. The observations with adaptive optics
reveal highly collimated bipolar dust and gas ejection located just in the
plane of the sky. The small inclination angle of ejection may confirm
the presence of an orthogonal dust structure closely aligned to the line
of sight and causing the strong extinction observed after the outburst.
In the case of merging components, such an orthogonal structure might be
a dust disk formed due to angular momentum conservation, but the merging
hypothesis needs orbital period instability of the progenitor. Otherwise,
this structure should be radially expanding in the equatorial plane.
Our optical photometry reflecting changes of fluxes in He I, [O II], and
[O III] emission lines shows that the dust minimum finished in 2004,
leading to strengthening of these lines. Later on, emission line fluxes
continued to decrease. This means that the ionizing and exciting radiation
comes from outside the place of the explosion, from the dust surroundings.
However, no stellar object was seen in the latest adaptive optics image taken
in March, 2007.
The absence of a stellar source and the presence of an elongated structure
formed only from the ejecta in the frames taken with adaptive optics
suggest that components of the system lost their mass and might be
destroyed totally or partly. The interpretation of the event as a helium
flash is of no doubt. The dynamic theory of such an explosion predicts
detonation of the CO white dwarf core and destruction of the white dwarf.
Such an explosion will undoubtedly result in the loss of the binary's
common envelope. The remnant may be the core of the helium donor which remains
the source of ionizing radiation. In the hypothesis of double detonations,
ejecta of the two detonations, the helium shell, and the white dwarf core may be
An image of V445 Pup taken on 20 October 2009 with the SAO 1-m Zeiss
telescope. This is a sum of V and RC frames.
This image was obtained on the night with the best seeing,
FWHM = 1".5. No elongation is seen in the V445 Pup image (left). To the right,
we compare pixel brightness vs. pixel distance from the star center relations
for the V445 Pup remnant (bottom) and a nearby comparison star (top). The curves are
averaged star image profiles; the vertical lines indicate star image radii at
the half intensity.
We analyzed our CCD frames taken with the SAO 1-m reflector. The
best-resolution frame with the seeing FWHM = 1".5 was obtained
on 20 October 2009. A part of this frame is shown in Fig. 9
(left). The star was faint, and to increase the signal in pixels, we co-added
In the right panel, the profile of a standard star (top) is compared to
the profile of V445 Pup (bottom). These are normalized dependences of pixel
counts on the distance between the pixel center and the center of the
star image localized by the method of star-profile dispersion minimization.
This distance is measured in the units of pixel size, 0".44.
Both images are round, and the image of V445 Pup
is marginally resolved. No image elongation is seen, whereas the expected
elongation based on Fig. 6 by Woudt et al. (2009) should be 2:1.
One should note that the radiation in emission lines was resolved
by Woudt et al. (2009) with VLT spectra, and very probably, it
coincides spatially with the lobes seen in the infrared K band.
In both cases, either CO white dwarf detonation and destruction or
big mass loss from the helium donor,
the future SN Ia scenario for V445 Pup becomes impossible. The helium flash
accompanied with core detonation seems to be a single event in the life
of a white dwarf in the binary system with a helium star. It does not
look like a SN Ia. At least, it is too early to discuss
the SN Ia scenario for V445 Pup before the possible detection of restored
mass exchange in the binary system.
We found that V445 Pup progenitor was a binary system. The most
probable orbital period was 0d
The total variability amplitude was 0m
.7 in the B
and the total amplitude of the periodic variability component was
.4. Solutions of the light curves with
different periods are also possible.
The shape of the progenitor's light curve suggests that it was a common
envelope binary with a spot on the surface of the envelope and with
variable surface brightness.
The energy distribution of the progenitor can be successfully fitted with
that of a single A0V star with a small infrared excess.
The shape of the light curve after the outburst suggests that the dust
absorption minimum finished in 2004. We do not see any stellar remnant
either in optical range or in the infrared. No stellar object, like a
reddened A-type star, is visible. We argue that V445 Pup was an event
of double detonation, both of the helium envelope of the CO white
dwarf and the CO white dwarf's core, accompanied with the destruction of
the common envelope of the binary.
In such an event, one of the components of V445 Pup lost a
part of its mass, and the other one might be totally destroyed. Thus,
V445 Pup, as a remnant of the outburst, cannot be a progenitor of a future
The authors thank A. Retter and A. Jones for their observations of V445 Pup in
outburst provided for our analysis. We also used the Digitized Sky Survey,
the SuperCOSMOS Sky Surveys, the United States Naval Observatory (USNO) Image
and Catalog Archive, the USNO B1.0 astrometric catalog, and the All Sky
Automated Survey (ASAS) data. The information from deep sky
surveys is very useful for understanding the nature of astrophysical
objects and phenomena.
E.A.B., A.V.Zh. and V.P.G. thank Russian Foundation for Basic
Research (RFBR) for financial support through grant
No. 07-02-00630. S.Yu.Sh. is grateful to the RFBR for grants No. 08-02-01229
and 09-02-90458, to the Slovak Academy of Sciences for the VEGA Grant 2/0038/10,
and to the Russian Ministry of Education of Science for the grant RNP-2906.
Ashok, N.M., & Banerjee, D.P.K., 2003, Astron. & Astrophys.
Deeming, T.J., 1975, Astrophys. & Space Sci.
Fink, M., Roepke, F.K., Hillebrandt, W., et al., 2010,
Astron. & Astrophys.
Glasner, A., & Livne, E., 2002, AIP Conf. Proc.
Guillochon, J., Dan, M., Ramirez-Ruiz, E., & Rosswog, S., 2010,
Hachisu, I., Kato, M., & Saio, H., 1989, Astrophys. J.
Henden, A.A., Wagner, R.M., & Starrfield, S.G., 2001,
, No. 7730
Iben, I., Jr., & Tutukov, A.V., 1994, Astrophys. J.
Iijima, T., & Nakanishi, H., 2008, Astron. & Astrophys.
Kato, M., Saio, H., & Hachisu, I., 1989, Astrophys. J.
Kato, M., Hachisu, I., Kiyota, S., & Saio, H., 2008, Astrophys. J.
Kato, T., & Kanatsu, K., 2000, IAU Circ.
, No. 7552
Lafler, J. & Kinman, T.D., 1965, Astrophys. J., Suppl. Ser.
Liller, W., 2001, IAU Circ.
, No. 7561
Lynch, D.K., Russell, R.W., & Sitko, M.L., 2001, Astron. J.
Lynch, D.K., Rudy, R.J., Mazuk, S., et al., 2004, IAU Circ.
Meng, X. & Yang, W., 2010, Astrophys. J.
Neckel, Th., Klare, G., & Sarcander, M., 1980, Astron. & Astrophys.,
Platais, I., Kozhurina-Platais, V., Zacharias, M.I., &
Zacharias, N., 2001, IAU Circ.
, No. 7556
Smith, N., Ganeshalingam, M., Chornock, R., et al, 2009, Astrophys. J.
Straizys, V., 1982, Zvezdy s defitsitom metallov (Metal
, Vilnius: Mokslas (in Russian)
Terebizh, V.Yu., 1992, Analiz vremennykh ryadov v astrofizike
(Analysis of Time Series in Astrophysics)
, Moscow: Nauka (in
Wagner, R.M., Foltz, C.B., & Starrfield, S.G., 2001, IAU Circ.
, No. 7556
Williams, R.E., 1992, Astron. J.
Wooden, W.H., II, 1970, Astron. J.
Woudt, P.A., Steeghs, D., Karowska, M., et al., 2009, Astrophys. J.
Colours and magnitudes of comparison stars measured by A. Henden for VSNET
Moscow B magnitudes of V445 Pup
Sonneberg B magnitudes of V445 Pup
Archive photographic magnitudes of V445 Pup
from Digital Sky Survies
CCD observations of V445 Pup after the outburst
1SO: Special Astrophysical Observatory 1-m Zeiss
observation; VA: SAI Crimean Station 0.6-m Zeiss telescope.
Periods and frequencies revealed
2d.w.: a double-wave light curve.