Issue |
A&A
Volume 513, April 2010
|
|
---|---|---|
Article Number | A6 | |
Number of page(s) | 15 | |
Section | Stellar structure and evolution | |
DOI | https://doi.org/10.1051/0004-6361/200913480 | |
Published online | 13 April 2010 |
A survey for pulsating subdwarf B stars
with the Nordic Optical Telescope![[*]](/icons/foot_motif.png)
R. H. Østensen1 - R. Oreiro1 - J.-E. Solheim2 - U. Heber3 - R. Silvotti4 - J. M. González-Pérez5 - A. Ulla6 - F. Pérez Hernández5,7 - C. Rodríguez-López6,8 - J. H. Telting9
1 - Instituut voor Sterrenkunde, K. U. Leuven,
Celestijnenlaan 200D, 3001 Leuven, Belgium
2 - Institutt for Teoretisk Astrofysikk, Universitetet i Oslo, 0212
Blindern-Oslo, Norway
3 - Dr. Remeis-Sternwarte, Astronomisches Institut der
Univ. Erlangen-Nürnberg, 96049 Bamberg, Germany
4 - INAF-Osservatorio Astronomico di Torino, Strada dell'Osservatorio
20, 10025 Pino Torinese, Italy
5 - Instituto de Astrofísica de Canarias, 38200 La Laguna, Tenerife,
Spain
6 - Departemento Física Aplicada, Universidade de Vigo, 36310 Vigo,
Spain
7 - Departamento de Astrofísica, Universidad de La Laguna, 38205 La
Laguna, Tenerife, Spain
8 - Laboratoire d'Astrophysique de Toulouse-Tarbes, Univ. de
Toulouse, 14 Av. Edouard Belin, 31400 Toulouse, France
9 - Nordic Optical Telescope, 38700 Santa Cruz de La Palma, Spain
Received 15 October 2009 / Accepted 6 January 2010
Abstract
Context. A search programme for pulsating subdwarf B
stars was conducted with the Nordic Optical Telescope on La Palma over
59 nights between 1999 and 2009.
Aims. The purpose of the programme was to
significantly extend the number of rapidly pulsating sdB stars to
better understand the properties of this new group of variable compact
stars.
Methods. Candidates were selected initially from the
HS and HE surveys, but were supplemented with additional objects from
other surveys. Short sequences of time-series photometry were made on
the candidates to determine the presence of rapid pulsations.
Results. In total twenty new pulsators were found in
this survey, most of which have already been published and some
extensively studied. We present four new short period pulsators,
bringing the total of such pulsators up to 49. We also give limits on
pulsation amplitudes for 285 objects with no obvious periodic
variations, summarise the results of the survey, and provide improved
physical parameters on the composite pulsators for which only
preliminary estimates were published earlier.
Key words: subdwarfs - surveys - stars: oscillations
1 Introduction
Hot subluminous stars are considered to be extreme horizontal branch
(EHB)
stars or closely related objects, with effective temperatures
20-35 kK.
The EHB models imply that they
are core helium burning objects with an extremely thin (
)
inert hydrogen dominated envelope (Saffer et al. 1994; Heber 1986).
This structure prevents them from ascending the asymptotic giant branch
(AGB), and they must evolve instead towards higher temperatures and
surface
gravities after their core helium is exhausted. Thus, an sdB star
crosses the hotter sdO domain before reaching degeneracy and cooling as
a normal white
dwarf star (Dorman et al. 1993).
However, important questions remain concerning the exact evolutionary
paths and the appropriate timescales.
How they evolve to the EHB configuration is controversial. The
problem is how the mass loss mechanism in the progenitor manages
to remove all but a tiny fraction of the hydrogen envelope at
about the same time as the core has attained the mass
(0.47
)
required for the He flash. About half of the sdB stars reside in close
binary systems, with either
a white dwarf or an M-dwarf as companion
(Maxted et al. 2001;
Napiwotzki et al. 2004),
and a significant
fraction of the rest are in wider binaries with a main sequence F-K
star
as companion. Therefore mass transfer in close binary evolution must be
an important evolutionary pathway leading to the formation of sdB stars
(Han
et al. 2003,2002). A recent review of hot
subdwarf stars is provided by
Heber (2009).
The discovery of multimode pulsations in sdBs opened an
attractive opportunity
of probing their interiors with seismological methods.
The properties of the sdB pulsators (sdBVs or
V361 Hya
stars after the prototype; Kilkenny
et al. 1997) are characterised by relatively short
pulsation periods
(
1
to 10 min) and low pulsation amplitudes attributed to
low order pressure (p) modes.
Most V361 Hya stars have been found with pulsation amplitudes
of around
ten millimodulation amplitudes (mma
),
although a few objects show a main peak amplitude up to
60 mma
(V338 Ser, Kilkenny
et al. 1999; Balloon 090100001,
Oreiro et al. 2004),
while the lowest level pulsator discovered up to now is
LM Dra
(Silvotti et al. 2000),
with no peaks above 2 mma.
The number of detected periods varies from a single one to more than
fifty,
strongly correlated with the accuracy of the measurements.
In 2001, some sdBs were discovered to show long-period (1 h)
photometric modulations (Green
et al. 2003), which were interpreted as high
radial order gravity (g) modes. These stars are
referred to as long period sdBV
stars or V1093 Her stars after the prototype.
V1093 Her stars are cooler than their short-period
counterparts, although
their instability regions appear to overlap around
kK, where they can
display simultaneous short
and long period pulsations. These stars are referred to as hybrid sdBVs
or DW Lyn stars after the prototype (Schuh
et al. 2006).
Non-adiabatic computations by Charpinet
et al. (1997) predicted excited low
degree (),
low radial order (n) p-modes in
sdB models,
driven by an iron opacity bump, in the temperature range
29-37 kK. This mechanism is inefficient at solar metallicity,
but gravitational settling and radiative levitation can work together
to
locally enhance metals in a driving zone in the envelope.
This
mechanism has been successfully invoked to explain both
the p-mode pulsations in V361 Hya stars
and the g-mode
pulsations in V1093 Her stars (Fontaine
et al. 2003).
Jeffery & Saio (2006)
were able to expand the instability regions
by including opacities for nickel, to the point where the
p-mode and g-mode domains
overlap, thereby explaining
the presence of hybrid oscillations.
However, it is still not understood why most stars in this temperature
range do not appear to vary.
As the number of known short-period and hybrid sdBV stars have increased, several peculiar features are emerging:
- Mode density: the number of pulsation frequencies detected vary from a single one to more than 50. LS Dra (Østensen et al. 2001a) was confirmed to be monoperiodic well below the millimagnitude level from 127 h of observations by Reed et al. (2007b). For others, detailed follow-up observations have increased the number of detected frequencies to more than 50.
- Period grouping: pulsation periods for the whole V361 Hya group lie between 90 and 800 s, but distinctions appear to be present between the stars in the shorter and longer period groups. All the sdBVs with periods between 300 and 400 s except one are DW Lyn type hybrid pulsators. The odd one out is KL UMa, which also stands out as the only binary of the pulsators with periods in this range (O'Toole et al. 2004).
- Amplitude variations: amplitudes are found to change with time in many cases. However, it is unclear whether this is due to true amplitude variations or beating of closely spaced modes (Kilkenny 2010).
- Dominant modes: according to their amplitudes the stars can be grouped into high amplitude pulsators, for which a mode was observed to exceed 30 mmag at least once, and low-amplitude ones for which such dominant modes were never reported. All the high amplitude pulsators are DW Lyn type, except the unique V338 Ser, which lies well above the canonical EHB region.

Nevertheless, sdBVs with sparse pulsation spectra are important as they may allow the measurement of light travel time variations or period changes due to stellar evolution, if the modes are sufficiently stable, as exemplified by the planet-hosting V391 Peg (Silvotti et al. 2007), discovered by our survey (Østensen et al. 2001b). Hence, photometric surveys to discover new V361 Hya stars are still rewarding.
The theoretical location of the sdBV instability strip in the
-
plane
is quite well established.
However, no more than one out of ten sdB stars located within the
instability region are actually found to pulsate. This situation is
quite different from that of white dwarf pulsators of the ZZ Ceti type,
for which evidence is
growing that all white dwarfs in their instability strip are pulsating
(Gianninas et al. 2007).
For sdBVs it has been suggested that younger EHB stars may not have
enough iron accumulated in the driving region for the
-mechanism
to be efficient enough to drive pulsations, but time-dependent
diffusion calculations by Fontaine
et al. (2006) have
demonstrated that sufficient iron for pulsations to occur accumulates
after
a few hundred thousand years, and after 1 Myr no further
accumulation
is achieved. Compared to the EHB lifetime of 100-150 Myr this
is not significant. An interesting speculation was made by Jeffery & Saio (2007),
that the iron group element enhancements
may be disrupted by the atmospheric motions as pulsations
build up to some level. They note that since p-modes
mostly involve
vertical motion, while g-modes are dominated by
horizontal motion,
it is possible that p-modes are more effective at
redistributing the
iron group elements out of the driving zone. This could explain the
observation that most cool EHB stars are g-mode
pulsators
while at the hotter end of the branch most are not (Green et al. 2003).
It may also be that the shortage of pulsators is due to observational
biases, either that the pulsations have too low amplitudes, or that
they
have too high
to be easily detectable photometrically, an issue
that needs further investigation.
To this end it also is important to publish observational constraints
on
stars that were not found to be variable in order to guide future more
sensitive surveys.
Here we present the final results of a long-term programme to detect new short-period sdBVs using the Nordic Optical Telescope ( NOT) on La Palma.
2 The survey
Shortly after the discovery of the first pulsating sdBs by the South
African
group (Koen
et al. 1997; Stobie et al. 1997a; O'Donoghue
et al. 1997; Kilkenny et al. 1997),
it was decided to
initiate a programme to search for such pulsators using the
2.5 m NOT on La Palma.
The principal aim was to extend the number of sdB pulsators,
determine the fraction of variables within this class of stars,
and explore the boundaries of their instability region.
The South African group had at the time discovered and confirmed
pulsations
in 14 objects from a sample of around 600 spectroscopically
confirmed
hot subdwarf stars, and a Canadian team found four pulsators in 74
stars
after selecting candidates based on
(Billères et al. 2002).
Since those samples were overlapping considerably, two of the pulsators
were found in both samples.
To avoid searching an already depleted sample, we drew our
targets primarily from a new spectroscopic study targeting sdB star
candidates from the Hamburg-Schmidt (HS) survey (Hagen
et al. 1995), although
it was complemented with targets from other surveys as explained below.
The complete list of observational runs fully or partly used for this survey is presented in Table 1. In Table 5 we provide coordinates and observational limits on all sdB stars checked for variability during the survey, except those that were found to pulsate or vary due to eclipses and reflection effects. The eclipsing and reflection variables are published in separate papers, and the 20 pulsators are listed in Table 9, together with pulsators found by other surveys.
2.1 Chronology and results published to date
The rapid pulsations in sdBVs are a challenge to observe with standard CCD instrumentation due to the pulsation periods being comparable to the typical readout times of a full CCD frame. An essential part of this project was the design of a system capable of efficient windowed CCD photometry (Østensen & Solheim 2000). A dedicated camera was built for this purpose, the Tromsø CCD Photometer (TCP), and the control system implemented for this camera was adopted to both the HIRAC and ALFOSC cameras at the NOT.
The first search run was made in the summer of 1999 (cf. Table 1), using the Tromsø-Texas Photometer (TTP) and revealing the first sdB pulsator from this programme (Silvotti et al. 2000, Paper I). Already the same autumn the CCD system was operational, and the improved detection efficiency allowed us to reach acceptable noise levels in a much shorter time, thereby resulting in the discovery of three new pulsators in only five nights of observations (Østensen et al. 2001b, Paper II). Two observation runs in 2000 led to the discovery of four more pulsators (Østensen et al. 2001a, Paper III). At this time we started to run out of targets from the HS survey at some RA's, and supplemented our list with stars from the PG survey (Green et al. 1986). This was a concern, as many of the known sdB stars had already been targeted by the South African group, and we expected our detection efficiency to drop since we avoided targeting stars already reported to pulsate. Still, the sample was not completely depleted, and during the summer 2001 run we discovered pulsations in QQ Vir (PG 1325+101) and EP Psc which originated from the HS sample (HS 2303+0152) but also appear in the PG catalogue as PG 2303+019 (Silvotti et al. 2002b, Paper IV).
Table 1: Observation runs at the Nordic Optical Telescope contributing to this survey.
To get a sample that was both sufficiently large to properly determine the extent of the sdB instability region and the fraction of pulsators across this region, and at the same time undepleted, it was decided to supplement our target lists with stars from the SDSS survey (Stoughton et al. 2002). The first three SDSS stars were observed in October 2002, and one of them, SDSS J171722.10+580559.9 or J1717+5805 for short, was found to pulsate (Solheim et al. 2004, Paper V).
In spite of the initial success of our programme, the Nordic time-allocation committee started turning down our proposals for further search time, and the last regular search run on the original programme was allocated for February 2005, but was completely lost due to bad weather. The last successful search run was made in June 2004, resulting in three new pulsators. These are PG 1419+081, SDSS J144514.93+000249.0 (J1445+0002) and SDSS J164214.21+425234.0 (J1642+4552). Only brief sequences were made on these objects, and the results were presented in Solheim & Østensen (2006, Paper VI).
At this time a collaborative search was initiated with a Spanish team (RO, AU, FPH, CRL), which permitted additional access to the NOT through Spanish time. With the Spanish group, we were able to explore a larger number of stars drawn from the Subdwarf Database (Østensen 2004,2006). The first pulsator discovered with the NOT on Spanish time was another star drawn from the SDSS sample PG 1657+416 (SDSS J165841.83+413115.6), and was presented by Oreiro et al. (2007, Paper VII). Observation runs allocated on Nordic time are indicated by run numbers prefixed by N, and on Spanish time by S in Table 1.
Nordic time still continued to be granted for follow-up programmes on various pulsating stars, including one of the most interesting pulsators discovered during this programme, V391 Peg, which was demonstrated to be host to a planet by Silvotti et al. (2007), by using the orbital period modulation introduced on the two main pulsation modes. A few unexplored pulsator candidates were observed during these runs (N6, N8, N9 in Table 1), and also during the WET X COV 24 campaign on the pulsating sdB star UY Sex (Vuckovic et al. 2006, run X1), whenever the prime target was not reachable. Upgrades and tests on the fast CCD photometry system also gave several opportunities to observe some additional targets during technical time at the NOT (listed as T1 to T8 in Table 1). One new pulsator, PG 1033+201, was found during run T7.
The limitations on telescope time from 2004 onward forced us to abandon the faint SDSS targets. Instead, we focused on detecting more pulsators in the cool region where the particularly interesting objects V391 Peg and Balloon 090100001 had been found. The discovery of hybrid pulsations in DW Lyn by Schuh et al. (2006) provided further encouragement for increasing the sample in this region of the instability strip. Although no more such cool pulsators were detected, a total of five more pulsators in the hotter end of the instability strip were found. One, from the BG sample (Bok-Green, Green et al. 2008), was recently published in Oreiro et al. (2009, Paper VIII), and four are presented for the first time in this paper.
In addition to these 20 pulsators our photometric observations have revealed several of the relatively rare sdB stars with M-dwarf companions from their strong reflection effect. The discovery of an eclipsing sdB+dM system, HS 0705+6700, first detected during run N4 was presented by Drechsel et al. (2001), and the non-eclipsing sdB+dM system HS 2333+3927 discovered during run N2 was presented by Heber et al. (2004). Two more sdBs with dM companions were found during run S2: the eclipsing HS 2231+2441 (Østensen et al. 2007), and HS 2043+0615 (Østensen et al. 2010, A&A, in prep.).
Our survey was not originally designed to detect long period variability, since we aimed for short photometric sequences with high S/N. However, the large photometric amplitudes of sdB+dM binaries produce a strong trend in the differential photometry, of which we saw several examples. However, since follow-up of such long-period variable objects can easily be done with smaller telescopes using long integration times, no follow-up was attempted on these stars with the NOT. The long period g-mode sdBVs also pulsate on time-scales around one hour, but with much lower amplitudes. As they had yet to be discovered at the start of our survey, our observation strategy was not designed with these objects in mind.
2.2 The HS/HE/SPY samples
When our programme was initiated, the main source of targets was the HS survey, but we also included targets from the equatorial part of the Hamburg-ESO (HE) survey (Wisotzki et al. 1996). Follow-up spectroscopy made at Calar Alto with the TWIN spectrograph (Heber et al. 1999; Edelmann et al. 2003) allowed us to preselect candidates with effective temperatures and gravities in the domain predicted for the pulsational instability (Charpinet et al. 1997). This provided us with a sample of pulsator candidates that was both tuned to the instability region and relatively undepleted by previous surveys. The latter was particularly important since the South African group had already explored a considerable fraction of the sdB stars known in the literature (and compiled in the Catalogue of spectroscopically identified hot subdwarf stars, Kilkenny et al. 1988). However, it has turned out that the HE survey has a significant overlap with unpublished parts of the Edinburgh-Cape (EC) blue object survey (Stobie et al. 1997b), which was one of the main sources surveyed by the South African group. This can explain why out of 80 stars from the HS survey and 20 stars from the HE survey observed during the initial five runs [as reported in][]solheim04, pulsations were found in eight HS stars and in none of the HE stars. Due to the difficulty of establishing accurate effective temperatures and gravities for sdB+F-K stars, these were not included in our initial sample, and were only added later when the priority sample started to become depleted. Thus, unlike the South African group, who found that most of their first pulsators were in binary systems, due to the selection biases involved only one of our initial 10 were. Note that another composite, our first pulsator LM Dra, is in a well-separated visual binary with an F3 companion, and therefore not considered to be an sdB+F system since the stars cannot have interacted during their evolution.
Only eight stars from the HS sample observed during this survey have not yet been published elsewhere. Their coordinates are listed in the on-line table together with those of all the other targets (Table 5), and their physical parameters are listed in Table 2. We computed the stellar atmospheric parameters (effective temperature, surface gravity, and photospheric helium abundance) by fitting model atmosphere grids to the hydrogen and helium lines visible in the spectra. The procedure used for this fitting was the same as that of Edelmann et al. (2003), but using only the LTE models described in Heber et al. (2000), as the purpose of our modelling was only to establish which stars are roughly within the sdB instability region rather than to establish precise physical parameters.
The parameters of HS 1909+7004 are
uncertain since the spectrum is
of low resolution and the fit is quite poor, but the fits to the other
stars
appear to be quite good.
HS 2320+0840 has a very strong companion that
contributes
clearly visible line features in the spectrum.
For this reason, we performed a spectral decomposition,
subtracting a main sequence model from Munari
et al. (2005)
with = 6.5 kK,
= 4 and
solar metallicity, which
we found to contribute about 40% of the flux at 5000 Å.
Table 2: Analysis for eight stars from the HS survey, not included in Edelmann et al. (2003).
A total of 46 objects from the HE survey were observed. For 23 of these the physical parameters were determined by Edelmann (2003), and none of these were found to pulsate. Cross-referencing them with the literature reveals that only nine are new, the rest are subdwarfs with identifications from the Montreal-Cambridge-Tololo (MCT) Survey (Lamontagne et al. 2000), or earlier surveys compiled by Kilkenny et al. (1988).
A number of white dwarf candidates from the HE survey were
observed as part of the SPY survey (Napiwotzki et al. 2001)
and published by Lisker
et al. (2005), where physical parameters are
provided also for the composite objects.
23 HE stars from the SPY sample were
observed, and two were found to
pulsate. Also, four HS stars from this sample were observed, and one
more pulsator was found.
These three new variables are presented in Sect. 5.
Lisker et al. (2005)
also present a list of sdB stars formerly misclassified
as white dwarfs. We observed four of these, but did not detect signs of
variability.
2.3 The SDSS sample
More than 16 000 spectra of UV excess stars were downloaded
from the
SDSS survey (Stoughton et al. 2002)
up to data release 6 (DR6), and the hot subdwarfs
were identified by inspection.
The intermediate resolution (1 Å)
SDSS spectra cover all
wavelengths from 3800 to 9000 Å, ideal for classification
purposes.
We identified about 900 objects as hot subdwarf stars
and classified them
into sdB, sdO, sdOB, He-sdB, He-sdO, and He-sdOB, depending on whether
hydrogen or helium dominates the spectrum, and if
they contain lines from He I, He II
or both.
318 of the sdB stars
were classified as single sdB stars and 165 as sdB+F-K binaries (34%)
based on
spectroscopic signatures and their position in (u'-g', g'-r')
colour-colour
space. A similar fraction was found for the sdOB stars;
78 single and 35 binaries (31%). All these stars were
fitted to model spectra to determine their position in the
/
plane,
using the Balmer lines H to H , as well as
the most prominent helium lines at 4472 Å and 4026 Å
for the sdB stars,
and including He II 4686 for the
sdOBs.
The formal fitting errors are about 250 K for
,
0.05 dex for
,
and 0.1 dex for
=
at
a g' magnitude around 16.0,
but increase by a factor of two
for the fainter stars between g' = 17.5
and 18. For the stars with
F-G companions the uncertainties are larger, but not so large
that the parameters are not useful for establishing whether or not the
stars
are located within the instability region.
We note also that for the hottest stars in the sample, NLTE effects
start to
become significant (Napiwotzki 1997),
and that the LTE models can
underestimate the temperatures by up to 1000 K in some cases.
Only 40 stars from this large sample were observed due to
observing time
restrictions. Five were found to pulsate and are listed in
Table 4
with additional details in Table 9.
The remaining 35 objects are listed in Table 3, with the
physical
parameters we derived from the SDSS archive spectra.
The sample includes eleven spectroscopic binaries, and curiously all
five pulsators
found from the SDSS sample were among these. The physical parameters
for the
binaries have considerable systematic shifts depending on the
contribution from the main
sequence companion. For this reason we have made a spectral
decomposition of
the five pulsators in order to place them reliably in the
/
plane.
The details of the spectroscopic decomposition are given in
Table 4.
The six non-pulsating sdB+F-K systems were fitted as if they were
single stars.
Table 3: The stars from the SDSS survey observed during this programme, excluding the pulsators listed in Table 4.
Table 4: Spectral decomposition of the binary pulsators from the SDSS sample.
It is non-trivial to estimate the spectroscopic parameters of
the
companion, but a reasonable compromise can usually be made that
approximates
the spectral contributions from the main sequence companion to the
Ca II H and K lines as well as the g-band
and the Mg I lines.
For the spectral decomposition we used the SDSS spectra processed with
the
DR7 pipeline, which have a significantly improved flux calibration
compared
with earlier data releases. After dereddening the spectrum by the
E(B-V)
reddening coefficients provided by the Schlegel
et al. (1998) dust maps,
the flux was
fitted with one of our subdwarf model spectra and a main sequence model
spectrum from Munari et al.
(2005). The parameters of the subdwarf are then refitted
to the observed spectrum after subtraction of the main sequence model,
and the
procedure was iterated until it converged.
Note that the reddening coefficients are for lines of sight to infinity
and
that we have no guarantee that the particular subdwarf under
consideration is
not in front of the bulk of the dust. There is also a significant
uncertainty
in the E(B-V)
values due to the limited resolution of the Schlegel
et al. (1998)
maps and the lumpiness of the interstellar medium.
But due to the low resolution and limited signal in the SDSS spectra,
it
is hard to make a complete disentanglement of the spectra while leaving
the reddening as a free parameter, so we will stick with the
dust map values.
In the case of J1445+0002, which has the highest
reddening coefficient of the five, it is quite clear that the Ca II
lines
are too narrow to originate from the main sequence companion, and no
solution
can be found that removes the Ca II
contribution from the
subdwarf spectrum. Therefore, the Ca II
signature is mostly
interstellar in origin, consistent with the high reddening value.
The strongest companion by far is that of J1642+4552, in
which it
contributes almost half the light at
Å
(f5000
in
Table 4).
The system is also in
a region of very low extinction, reducing the ambiguity of the spectral
decomposition. The strong companion and good signal in this particular
SDSS spectrum allow us
to make some further constraints on the companion. It is clear that a
solar metallicity makes a far stronger contribution to the metal
bands than allowed by the observed spectrum. A better fit is achieved
with a spectrum metal depleted relatively to solar by [M/H] = -1.5.
The temperatures derived here are somewhat lower than the
temperatures
we originally used to select the stars, but not enough to make a
difference with respect to the sample selection,
as we suspected oreiro07.
For J1445+0002 the temperature drops from 37.6 to a more
reasonable 35.9 K. It is suspicious that the derived
values for all the five
composite stars are so similar
(within
0.1 dex),
and might be due to degeneracies when
fitting so many free parameters to spectra with rather low resolution.
It has also been noted that more reliable estimates of the surface
gravity are obtained when the spectra cover all the high order Balmer
lines, a requirement the SDSS spectra do not satisfy.
2.4 The Bok-Green sample
Eleven stars in our survey were selected from a spectroscopic campaign to obtain a large unbiased sample of hot stars based on 2 MASS colours, undertaken by Green et al. (2008) with the University of Arizona 2.3 m Bok telescope. Physical parameters were derived from a subsample of 89 stars from this campaign by Winter (2006, Appendix C). None of the stars we observed have reliable spectroscopic classifications in the literature, but three have been listed as faint blue stars by earlier surveys (Ton 930, FBS 1133+754, and FBS 1224+780). However, none of these were known to be hot subdwarf stars, so we consider all eleven as belonging to an undepleted sample. One of the sdBs in the BG sample was found to pulsate, 2M0415+0154 = 2 MASS J04155016+0154209, as reported in oreiro09. The remaining seven stars are listed with 2M designation in Table 5.
2.5 Literature
To supplement our sample with targets to fill all RAs we added stars with temperatures and gravities published in the literature. Most of these additional objects are well-known subdwarfs from the PG survey and are included in many studies. The largest sample used was that of Saffer et al. (1994), where 24 targets were surveyed and one found to pulsate (LM Dra).
Seventeen PG stars from the radial velocity survey of Maxted et al. (2001), and four stars from the HST study of Heber et al. (2002) were also included in our sample, but none were found to be variable.
Rough temperature estimates can also be made on the basis of
photometric
measurements alone, and several works provide such estimates.
The survey of Beers et al.
(1992) includes temperatures computed from UBV
colours, and we observed eight of these stars.
A handful of composite sdB+MS stars from the list of Allard et al. (1994),
where their estimate of the temperature for the primary lie within the
instability region, were also observed.
Another four stars were targeted based on the temperature estimates
of Bixler et al. (1991).
Similarly, we observed stars from the sample of Moehler
et al. (1990), after
estimating the temperatures based on their Strömgren photometry
and the spectral features visible in their spectroscopic atlas.
In total, about fifty stars in the sample were observed based on such
photometric temperature estimates, and only one,
PG 1419+081,
was found to vary. The SDSS spectroscopy of this object was released
after our discovery, allowing us to refine the original colour
temperature
estimate and correctly identify the star as an sdB+G5 composite.
We found that the colour temperature (
= 33.8 kK;
Beers et al. 1992)
is encouragingly close to our model fit value (
= 33.3 kK).
If we look at the complete sample from Beers
et al. (1992) and consider only
the stars classified as sdB with temperatures above 27.5 kK we
are left
with 23 stars, of which we have observed nine and found one variable.
But
EC 20338-1925 is also in this sample and was in our
original list,
although excluded after pulsations were discovered by the South African
team
(published by Kilkenny
et al. 2006, but preliminary results from the 1998
discovery
were distributed much earlier).
Another star in the sample is
BPS CS 22169-1,
analysed by Geier et al. (2010, A&A, in prep.) and
found to have = 39.1 kK,
much hotter than the 33.9 kK estimate by Beers
et al. (1992),
and hotter than the predicted sdBV instability region.
The star BPS CS 22890-94 =
PG 1525+024 was not
observed since the colour temperature estimate by Beers
et al. (1992)
placed it below our cut at
= 22.9 kK.
However, recent SDSS
spectroscopy gives us
= 27.9 kK,
which would have placed the
object inside our cuts, further emphasising the uncertainties
associated with using colours as a temperature estimator.
Table 5: Truncated version of the catalogue of sdB stars observed during our search programme.
3 Observations
During our observations we systematically aimed at identifying as many clear pulsators as possible, rather than trying to push the limit on the pulsation amplitudes as low as possible. Thus, most of our detected variables have rather high pulsation amplitudes, and the sample may still contain objects with undetected amplitudes below 2 mma. Due to the shortness of our typical photometric sequences, there may also be multiperiodic pulsators among our null-detections, unidentified simply because their periods happened to be canceling each other out in a beat phase at the short time they were observed. The stars in Table 5 should therefore not be considered non-variable, rather, we prefer the term not observed to vary (NOV).
For the aims of this campaign we consider a star as NOV if
there is no peak
in its Fourier Transform (FT) with an amplitude higher than 3.0 times
the
mean amplitude level ().
However, we require an amplitude of 3.7
(Kuschnig
et al. 1997, 99% confidence level)
in order to call the object a pulsator. Objects that fall between these
two
tiers are considered pulsator candidates until further observations can
confirm
the detection.
Most pulsator candidates were reobserved in subsequent runs, but
11 objects
remain in Table 5
with this classification,
and will be discussed separately in Sect. 4.1.
Only the frequency region 1.5-10 mHz was used to
compute
homogeneously.
The low frequency range was excluded to avoid the higher mean amplitude
level caused by small trends in the light-curve due to extinction
effects
or sky transparency variations.
3.1 PMT photometry
The first observation run (N1) was made with the Tromsø-Texas 3-channel
photoelectric photometer, equipped with Hamamatsu R647
photomultipliers.
The data were reduced with the standard quilt
software package
developed for the WET (Whole Earth
Telescope; Nather et al. 1990).
The reductions include smoothing of the sky background, and
correction for extinction and linear trends.
A total of 13 stars were observed, and one pulsator detected.
However, due to problems with the guiding many of the sequences were
rather poor, giving a
of about 2 mma
in the FT. Several stars showed a behaviour that could
be interpreted as pulsations, but was actually due to periodic
motion in the apertures. For this reason three stars were reobserved
during the N2 run, and only the sequence with the lowest
is listed
in Table 5.
3.2 CCD photometry
Further NOT observations were done with
the High Resolution Adaptive Camera ( HIRAC)
for the second run
(N2 in Table 1),
and the Andalucía Faint Object
Spectrograph and Camera ( ALFOSC) for all
the remaining runs.
Both HIRAC and ALFOSC
were equipped with Loral, Lesser thinned,
2048 2048
CCD chips, up until January 2004, when ALFOSC
was
upgraded with an E2V CCD 42-40 device.
The sky area available for locating reference stars was limited to
and
arcmin2
respectively
for each of the two cameras. All CCD observations used the
multi-windowed
fast CCD photometry (Østensen
& Solheim 2000). Efficient fast photometry is easily
achieved with this mode, since only small regions
(windows) of the chip are read.
Observations were made with either a Bessell B-band
filter
( NOT #74, Bessell
1990), for the brighter targets, and
a much wider filter
( NOT #92) for the fainter ones.
This filter (hereafter referred to as the W, or
wide filter)
has the same centre as that of the V-filter
(5500 Å),
but is a full 2750 Å wide, effectively covering all bands from
B to R with more than 90%
transmission.
The cycle times were mostly set to 20 s, except for some of
the brightest
objects, which can have cycle times as short as 5 s, and the
faintest ones,
which may have times up to 40 s. Between 2.5 and
6.5 s of the cycle times
are readout and shutter overheads, the exact number depending on the
number of reference stars selected and the size of the readout windows.
The CCD data were reduced on-line with the Real Time Photometry (RTP) program developed by one of us as part of his Ph.D-project (Østensen 2000, see also Paper II). This software was particularly useful for the purposes of this survey, as it performs real-time reduction of the data during sequencing, displays the light-curve of the target, and computes the FT. Thus, the noise level and the presence of clear variability could be checked at the telescope, although careful data reduction was also done afterwards. The processing includes bias level removal, flat fielding, sky subtraction, extinction correction and aperture photometry using apertures that track each star's geometrical centre.
The optimal aperture for each sequence was selected after processing all data sets with apertures of a wide range of diameters and choosing the one that gave the lowest noise in the FT. The apertures tested ranged between 10 and 40 pixels in radius, which corresponds to 1.9 and 7.6 arcsec on the sky, while the best results usually stayed within the range of 18 to 22 pixels (3.4 to 4.2 arcsec), but on nights with particularly poor seeing the optimal radius could be as high as 30 pixels (5.7 arcsec).
4 Results
A total of 309
stars deemed to lie in the predicted instability region for sdB stars
were checked for variability during 25 observing runs from 1999 to
2009.
We did not apply a specific cutoff for
and
since our
estimate of these parameters were preliminary and sometimes crude.
Instead
we prioritised our targets by their distance from the centre of the
instability region at 34 kK. In the later runs we shifted our
focus more
towards the region where the hybrid pulsators are located.
Of all our targets, 20 were observed to pulsate
with frequencies corresponding to the V361 Hya stars,
which corresponds to a
ratio
of fast pulsators/NOV sdBs.
A further four stars were found to have such strong reflection effects
that the slope was obvious even in our short light-curves. This
fraction
of 1.3% is only a lower limit for the sdB+dM population, as there may
be quite a few more objects with longer periods, which our survey for
short period pulsators was not sensitive to.
In fact, we know of one more reflection variable in our survey,
BPS CS 22169-1, which we found to be
constant in our
half-hour light-curve.
Since these results were not obtained at the NOT
as part of this
survey, they will be reported in a future paper.
The 285 stars for which no obvious variability was detected
are presented in Table 5.
A different detection threshold was achieved for each target,
depending
on its magnitude, the filter used for the observations (B
or W), the
length of the photometric time-series, and the weather conditions in
each particular case. In the top panel of Fig. 1 the
mean amplitude level in the FT ()
obtained for the sample is shown as function of each star's magnitude.
A general trend of larger
for
fainter objects can be distinguished, although in most cases the
time-series sequence was stopped when the real-time FT showed no
significant peaks in the relevant region and the highest peaks reached
amplitudes
1 mma.
Theoretical
levels as a function of
magnitude for light-curves with 1, 5, 20, 100 and 300 data
points are also
plotted in this panel. They represent the contribution to
of both
the scintillation and photon shot noise, and thus should be understood
as
the limiting accuracy in the FT obtainable from a time series with the
same
number of data points.
The scintillation noise was computed following Young (1993), and the photon shot
noise was estimated using the signal-to-noise calculator
of the NOT. In both cases, a typical
airmass X=1.2 and a typical
exposure time of 20 s were used. Most points follow the
predicted
trends for light-curves with between 20 and 300 points, and the
outliers
are mostly due to poor weather conditions.
The middle panel of Fig. 1 shows, again
as function of ,
the maximum amplitude obtained in the FT for every observed target
within the
1.5-10 mHz frequency range. The continuous line indicates the
3.7
(99%) threshold, meaning that if
is above this line, the star is considered as pulsator. The dashed line
indicates the 3.0
(90%) threshold, and stars whose
falls below
this line are considered as NOV, according to our definition.
Fourteen of the twenty sdBVs discovered in this
survey are marked with stars (published earlier) and four with diamonds
(the new ones presented here); the remaining two pulsators,
EP Psc and QQ Vir, have such
large pulsation amplitudes
that they lie beyond the range shown in the figure.
Note that we have used only the discovery data for these plots, so the
lowest amplitude pulsator, LM Dra, and the one on
the
99% confidence line at
,
PG 1657+416,
have been confirmed in follow-up runs with higher confidence levels.
The objects between the 3.0 and 3.7
confidence lines are
discussed below, and the pulsators marked with diamonds will be
discussed
in Sect. 5.
The bottom panel of Fig. 1 represents an
histogram of the mean
amplitude level achieved along the survey. For most of the targets, the
noise
level in the FT was below 1 mma
(239 targets = ), and in 50% of the
cases a
mma
was obtained.
![]() |
Figure 1:
Top panel: photographic B
magnitudes of the targets
as function of |
Open with DEXTER |
4.1 Pulsator candidates
Ten objects in Table 5
have amplitude peaks between 3.0
and 3.7,
higher than our NOV requirement.
They must therefore still be considered pulsator candidates.
Three of these have a
of 0.33 mma or less, and have been ignored.
The others deserve particular mention.
One object, PG 1248+164 has such a high
that it
is off the scale in Fig. 1.
The object comes from the sample
of Maxted et al. (2001)
and deserves further observations.
The object on the 99% confidence line at
= 3.4 is
PG 1725+252. Normally it would have been reobserved,
but as it was
realised that this object is also in the survey of Billères et al. (2002)
with
a comfortable 0.03% limit, no follow-up was performed.
The object just below the 99% limit at
is
KPD 2215+5037, observed during run N5.
This star was also part of the survey of Billères
et al. (2002)
with a limit of 0.2%. We note a rising trend in our light-curve on
the 1% level over the 20 min covered by our observations,
which could
indicate that the object is a long period variable.
The three objects between the 3
and 3.7
confidence lines
at
1
are PG 1313+132, PG 1722+286 and
J1351+0234. PG 1313+132 shows variations
at low frequencies,
consistent with g-mode pulsations, which is likely
considering its low
temperature (
= 25.6 kK,
=5.41; Saffer et al. 1994),
but uncertain as the observations were done at relatively high airmass.
If we compute the noise level only for frequencies higher than
5 mHz we
get
= 0.53
and the highest peak is well below three times this value.
The light-curve of PG 1722+286 was obtained during
the problematic
N1 run. An attempt was made to reobserve it during T2, but poor weather
resulted in an even
worse light-curve. After a limit on pulsation of 0.08% was presented
by Billères et al. (2002),
no further attempts were made to observe this target.
The last of the three,
J1351+0234, is just noisy due to its low brightness (g' = 17.1).
The three objects that lie on the 3
confidence line around
= 2
are HS 1813+7247, EGGR 491, and
HS 2029+0301 (only the last of these is actually
above the
3
line, but we will discuss all three).
HS 1813+7247 has a slope in the light-curve, but it
is not consistent
between the different reference stars, and so is most likely due to
differential
extinction as the W-band filter was used.
EGGR 491 is another faint
target from the SDSS sample, so the light-curve is just noisy, with no
obvious
features. HS 2029+0301 is also from the N1
photoelectric run, and
has not been reobserved. Further observations would be required to give
a more
useful limit on any pulsational behaviour of this object.
5 New pulsators
The four pulsators found in our sample and not published elsewhere are HE 2151-1001, HS 2125+1105, PG 1033+201, and HE 1450-0957. The details of the discovery observations as well as follow-up observations on the first two done with the NOT are provided here. Table 6 includes the detailed log of observations for all these new sdBVs.
For all the pulsators we derived the frequency content of the
light-curves
by performing a non-linear least-squares fit to a sine function, using
the
highest amplitude frequency in the amplitude spectrum (top panels
of Figs. 2-5). This fit was
then subtracted from the light-curve
to compute the residual amplitude spectrum (bottom panels).
Notice that the noise level is recomputed after each prewhitening step,
so that the 3.7
level lines in the figures drop after removing
the established signal.
If there was still any peak left above the significance threshold,
the original light-curve was fitted to the sum of
two sine
functions with the established frequencies.
Table 7
lists the best fitting frequencies and amplitudes,
with the associated errors from the least-squares procedure.
Table 6: Details from the discovery runs for the four new sdBVs.
![]() |
Figure 2: Amplitude spectrum of HE 2151-1001. The lower panel shows the residual spectrum after prewhitening the two frequencies listed in Table 7 from the light-curve. The dashed line indicates 3.7 times the noise level. |
Open with DEXTER |
5.1 HE 2151-1001
This target was observed for 40 min
during the run S2.
Its amplitude spectrum is shown in Fig. 2, where some
signal above the threshold can
be distinguished at
8 mHz.
This frequency range is typical for a pulsating sdB with
these
-
parameters
(see Table 9).
Our non-linear least-squares frequency analysis revealed two
low amplitude
(3 mma) peaks, as listed in
Table 7.
Note that the frequency distance between the two
peaks (140
Hz) is
within the frequency resolution (
325
Hz).
Another short sequence on HE 2151-1001 was obtained
during S3
(see Table 6),
but the noise level achieved was not sufficient to detect the
oscillations above the significance threshold.
This target was observed spectroscopically as a white dwarf
candidate by
the SPY survey. The spectrum was analysed
by
Lisker et al. (2005),
who found no spectral indication of a companion,
but noted that it has a peculiar H
profile.
The 2 MASS photometry (J = 16.53,
H = 16.37, K = 15.91)
indicates a main sequence companion, but this is uncertain as the
values
are close to the faint limit of the 2 MASS
survey.
The NOMAD catalogue
(Naval Observatory Merged Astrometric Dataset, Zacharias et al. 2005)
provides reasonably good photographic magnitudes (B = 15.25,
V = 15.57, R = 16.09),
and the UV magnitudes (FUV = 14.55,
NUV =14.87) from the GALEX
satellite (Martin et al. 2005)
all indicate colours
typical for a hot subdwarf.
![]() |
Figure 3:
Amplitude spectrum of HS 2125+1105. The lower panel
shows residual spectrum after subtracting one frequency from the
original light-curve.
The dashed line indicates 3.7 times the noise level.
The peak at 7.3 mHz is just above the 3.7 |
Open with DEXTER |
![]() |
Figure 4: Amplitude spectrum of PG 1033+201. The lower panel shows the residual spectrum after prewhitening the two frequencies listed in Table 7 from the light-curve. We have also plotted the amplitude spectrum of the first half (red dashed curve) and second half (blue short dash, long dash curve), to demonstrate that the main peak is present with the same amplitude in both halves of the dataset. The dashed line indicates 3.7 times the noise level. |
Open with DEXTER |
5.2 HS 2125+1105
The faint (B = 16.3) HS 2125+1105 was first observed during run S2. A 6.4 mma peak was detected at 6.89 mHz, although its significance level was exactly 3.7 times the noise level (see Table 6). A year later, a slightly longer light-curve with better photometric conditions was obtained during run ST, which reduced the noise level to 0.95 mma. However, a peak detected at the same frequency also dropped its amplitude to 3.6 mma, again at exactly the 3.7 times the noise level. During run S3, the object was observed on each of the three nights as indicated in Table 6, and the good conditions finally allowed us to get reliable confirmation of the main peak. The amplitude spectrum obtained from the combined light-curve is shown in the upper panel of Fig. 3. In Table 7 the results of a frequency analysis are listed. A 4 mma peak at 6.83 mHz is detected, but this time at a reassuringly high S/N. In this sequence a second peak at 7.34 mHz was also detected, as listed in Table 7. The bottom panel of Fig. 3 shows the residual amplitude spectrum after prewhitening the main frequency from the light-curve.
In order to check the stability of this mode, an additional short sequence was taken in S4. The same f1 was found, this time with a 4.7 mma amplitude.
This HS star was not part of the original sample of sdB stars from Edelmann (2003), but was observed as part of the SPY survey and was analysed by Lisker et al. (2005). The star has reliable optical photometry (but no spectrum) in the SDSS survey (u' = 16.13, g' = 16.36, r' = 16.78, i' = 17.09, z' = 17.33). There is no indication of a main sequence companion from the spectroscopy or SDSS photometry. The object is well below the faint limit of the 2 MASS survey, so no IR colours are available. The GALEX photometry is FUV =15.6, NUV = 15.8.
![]() |
Figure 5:
Amplitude spectrum of HE 1450-0957 ( top)
and residual amplitude spectrum
after subtracting the main peak ( bottom), with 3.7 |
Open with DEXTER |
Table 7: Pulsation properties for the four new pulsators.
5.3 PG 1033+201
PG 1033+201 was one among 15 targets observed with
the B-band filter during the technical night T7.
During the run we had some problems with the observing system, which
resulted
in some frames in the photometric sequence being dropped. For this
reason
the real-time processing program was unable to process the entire
sequence.
Due to these problems we stayed on target
almost 40 min, which is
considerably longer than our standard sequence on such bright targets.
Eventually, we
terminated the observations in order to reinitialise the CCD system,
but
we never returned to reobserve this star. During the careful
reprocessing
of the photometry for this paper, we excluded two lost and one bad data
point
from the 118 frame sequence, and performed a regular frequency analysis
of
the cleaned light-curve.
The amplitude spectrum is shown in the top panel of
Fig. 4
as a continuous line. A clear peak at 5.8 mHz is
well above the 3.7
(horizontal dashed) line. We divided the
light-curve into two equal halves, and we can see from the
corresponding
spectra (also shown in Fig. 4 as red dotted
and blue
dash-dotted curves), that the main period is present in both halves
with
the same amplitude. A second frequency at 6.8 mHz is
significant only in the first
part, indicating that it is either spurious or a combination of
frequencies
suffering cancellation from beating within our short sequences.
This exercise demonstrates how short-period pulsations can be detected
with
reasonable confidence even in short (
20 m) light-curves, as long as
one acquires adequate signal-to-noise in the individual data-points.
The bottom panel displays the residual amplitude spectrum after
subtracting
the frequency at 5.8 mHz. The remaining peak at
6.8 mHz
becomes
significant when the noise is computed in the prewhitened spectrum
(see Table 7).
We conclude that PG 1033+201 appears to be
a typical short-period sdB pulsator with a main pulsation amplitude of
5 mma, with likely further frequencies at the 2 mma
level.
This object was included in our sample based on the decomposition by Allard et al. (1994), who estimated the system to be an sdB+F9 composite from their Cousins BVRI photometry (V = 15.67, B-V=-0.17, V-R=-0.01, R-I=+0.06), and provided a colour temperature for the primary of 31.5 kK. The object occurs in the SDSS survey as J103638.93+195202.2, and has photometry (u' = 15.14, g' = 15.39, r' = 16.80, i' = 16.01, z' = 16.21) but no spectroscopy, which prevents us from acquiring a better temperature and a gravity estimate. The 2 MASS IR photometry clearly supports the notion of a main-sequence companion (J = 15.46, H = 15.16, K = 15.13, i.e. J-H > +0.3, see Reed & Stiening 2004).
5.4 HE 1450-0957
HE 1450-0957 was one of the three targets surveyed during the final run of the programme (S6). A rather long light-curve was obtained (1h5min), since its oscillations were already clear as the data were processed as they arrived. Its amplitude spectrum (upper panel of Fig. 5) reveals the multimode behaviour of this target. At almost 7 mma, the main peak at 7.2 mHz has the highest amplitude of the four new pulsators presented here. A significant second peak is also detected in the amplitude spectrum of HE 1450-0957 (see bottom panel of Fig. 5), but more low amplitude modes are likely to remain in the range 6-9 mHz. The frequencies derived from this discovery light-curve are also listed in Table 7.
This object came into our sample from the HE stars surveyed by SPY (Lisker et al. 2005). It is included in the Edinburgh-Cape survey as EC 14507-0957 where reliable photometry can be found (V = 15.27, U-B=-1.04, B-V=-0.21). As noted by Reed & Stiening (2004), its 2 MASS magnitudes (J = 15.58, H = 15.41, K = 15.36) are indicative of a main sequence companion, but this companion must be very weak in the optical as Lisker et al. (2005) do not detect any trace of it in their VLT/ UVES spectrum.
6 Summary and discussion
The NOT search programme for pulsating sdB
stars has been a great
success,
contributing twenty new pulsators to the known population of short
period
sdBVs. Together with the two new pulsators presented by Kilkenny et al. (2009),
another two recently found by Barlow et al. (2010,2009),
and HE 0230-4323,
our latest tally brings the total up to 49.
Our effort to implement an efficient system for CCD photometry at
the NOT has certainly been fruitful,
producing on average one new
pulsator for every three nights of observations. This is a
substantially
higher efficiency than any other group has reported.
6.1 Follow-up
Many of our pulsators have already received detailed follow-up. Most effort has been put in on V391 Peg, which was the subject of several campaigns, and its pulsation periods are still being monitored regularly (Silvotti et al. 2007). The original aim of this monitoring was to detect evolutionary period changes, but the periodic variation that was found is more likely to be caused by an orbiting planet inducing period variations through the light travel time effect. Several other stars in our sample are now regularly monitored for variations in the pulsation period under the EXOTIME project (Lutz et al. 2009). As a substantial time base of these observations develops, we expect to detect more planet-hosting hot subdwarf stars.
QQ Vir was the first sdBV for which spectroscopic line-profile variations were used to constrain its main pulsation mode (Telting & Østensen 2004), and was the subject of an extensive photometric campaign in 2003 (Silvotti et al. 2006) revealing 15 pulsation frequencies. The detailed frequency spectrum from this campaign allowed Charpinet et al. (2006) to obtain an asteroseismic solution using the forward method. QQ Vir also became the first star from our sample to be targeted with high-resolution spectroscopy with the VLT in 2008 (Telting et al. 2010).
LM Dra was followed extensively by Reed et al. (2007a) with multisite observations between 2003 and 2004, and in spite of its extremely low pulsation amplitudes they detected six pulsation periods with amplitudes between 1.0 and 2.2 mma. V429 And and V1636 Ori were followed by Reed et al. (2007b) in Nov.-Dec. 2005, revealing a rich frequency spectrum with fourteen significant peaks between 0.45 and 4.35 mma in V429 And, and a simpler spectrum with only three significant peaks between 0.8 and 11.1 mma in V1636 Ori. Reed et al. (2007b) also report that LS Dra appears to be monoperiodic, but with a highly variable amplitude, changing from a maximum of 5.26 to a minimum of 0.87 mma over their campaign spanning 47 days. They further confirm three of the four pulsation periods in V387 Peg.
J1717+5805 was targeted with high-speed simultaneous 3-channel photometry using ULTRACAM on the WHT, as reported by Aerts et al. (2006), resolving the broad peak described in solheim04 into two separate frequencies.
Table 8: Comparison with the results of other surveys for pulsating sdBs (upper part) and our own survey divided into subsamples (lower part).
6.2 Summary
The four new pulsators presented in this paper certainly deserve more attention as well. Confirming the pulsations in PG 1033+201 would be a priority, but an intensive campaign would probably be required to disentangle its multiperiodic nature. The pulsations observed in HE 2151-1001 are barely above our significance threshold and need to be confirmed. HS 2125+1105 is clearly a stable pulsator, as the same oscillation period was detected in six sequences spanning four years, but we would suspect more frequencies to be discovered with more dedicated follow-up. HE 1450-0957 is also a convincing case where residual power in the amplitude spectrum between 6 and 10 mHz indicates the presence of unresolved pulsations that could be revealed with dedicated follow-up. The existence of pulsations in this star was recently confirmed by observations from South Africa (C. Koen, priv. comm.).
In Table 8 we compare the fraction of pulsators detected in each of the different samples we have drawn our candidates from. A fraction of pulsators between 5 and 12% is clear when targets are selected based on spectroscopic temperature estimates. We estimate that it should be possible to obtain a pulsation fraction of about 10%, when starting with an undepleted sample and reaching a noise level below 1 mma. This could probably be increased by focusing only on the hot end of the EHB, between 31 and 36 kK, at the cost of sacrificing discoveries of the rarer pulsators at the cool end of the instability region.
That the colour selected sample has a lower success fraction than the samples based on spectroscopy is not surprising, but it is surprising that there are so few pulsators in the sample from Saffer et al. (1994). The total number of sdB stars in that sample hotter than 28 kK are 44, of which we observed 16 (plus another eight that were outside the borders of the instability region). Excluding the known g-mode pulsators on the cool end of that sample, we have two short period sdBVs in 41 stars; QQ Vir found by us and V1078 Her found by Bonanno et al. (2003). A large number of stars from this sample were surveyed by Billères et al. (2002), and the remaining stars have been checked by other observers (but no limits have been published), so we do not expect more clear pulsators in this sample, although some low amplitude variables may have been missed. Similarly, the number of sdB stars in the instability region in the sample of Maxted et al. (2001) is 34. We found no pulsators in this sample, but it has been heavily exploited by other groups; Billères et al. (2002) surveyed 14 of them and found one: UY Sex. When combining information from this and other unpublished surveys, we can only find a single star that has yet to be surveyed. Thus, the best fraction from the Maxted et al. (2001) sample we can infer is 3%, which is lower than what we have found in the HS, SPY and SDSS samples, but with such small samples this is barely significant.
6.3 The big picture
Table 9: All short period sdBV stars published to date.
We have compiled all the known short period sdB pulsators in Table 9, using data from the literature and added the results from this paper. We included preliminary physical parameters on two new pulsators published by Kilkenny et al. (2009), based on recent low-resolution spectroscopy from NOT/ALFOSC (3250-6150 Å at
The first sdB pulsator candidate in a globular cluster was
reported by
Randall et al. (2009a),
but while the single clear period at 114 s makes
this star a very likely sdBV, spectroscopic confirmation remains to
be done.
The same is the case for the pulsator candidate reported by Silvotti et al. (2009),
which was found in the Kepler satellite field-of-view, with a likely
pulsation period of 125 s found at the 3.5
level.
![]() |
Figure 6:
a)
|
Open with DEXTER |
We note that the number of pulsators in Table 9 with F-K companions are 18 of 49 = 37%, which is compatible with the fraction of spectroscopic composites we found in the SDSS as a whole. Pulsators also occur in systems with close white dwarf or M-dwarf companions, which means that pulsations appear independently of whether the subdwarf was formed through stable Roche-lobe overflow or via common envelope ejection. The stability of the pulsations for stars such as V391 Peg has revealed that some sdBs are definitely single, which implies that a formation channel must exist that produces single sdB stars that are otherwise indistinguishable from those produced via the channels that produces wide sdB+F-K or short period binaries.
In Fig. 6a
we show the location in the (
,
)
plane
of the pulsators from our sample (blue) together with pulsators from
the
literature (red), as listed in Table 9,
with symbol size proportional to
.
As our sample stars are fairly evenly spaced along the canonical EHB
strip,
it is immediately
clear that pulsations are more common on the hot side than on the cool
side, but the amplitudes are often much higher in the cool, low gravity
region.
The figure does not really reproduce the gap around
30 kK
between the main group
of pulsators on the hot end of the EHB and the cool group of g-mode
and hybrid pulsators seen in Fig. 3 of Østensen (2009), where
the spectroscopy was obtained from the independent BG survey
(Green et al. 2008).
This may be due to the fact that diverse methods
and model grids are used in the temperature determinations collected
from
the literature, and this can produce scatter as high as 2 kK
in
and
0.2 dex in
.
If we plot only the stars for which we ourselves have obtained the fits
on a
consistent model grid, the gap is indeed present.
The main cluster of pulsators correlates very well with the initial instability strip predicted by Charpinet et al. (1997), with only a few rare pulsators appearing to be anomalous. While the sample presented here contains no stars hotter than 39 kK, we have observed quite a few such stars without finding any with significant variability. But these are classified as sdO or He-sdO stars, and will be presented in a future paper. Candidates in that sample were selected from very diverse surveys and ended up having a wide range of physical parameters, as stars matching the properties of the only known sdO pulsator, J16007+0748 (Woudt et al. 2006) are extremely rare, preventing an effective preselection such as we made for the sdB sample.
Figure 6b
shows the surface gravities plotted against the
range of pulsation periods listed in Table 9.
As the
values are derived with many different methods, the systematic
errors can be substantially larger than
0.1 dex. Still, the expected
relationship between pulsation periods and surface gravity seems to
hold
up quite well (see Koen
et al. 1999b).
The relationship appears to break down for objects with surface
gravities higher than 5.8 dex, but this may not be a
real effect. Evolutionary
tracks for EHB stars with any envelope thick
enough to sustain
pulsations can hardly reach higher surface gravities than about
5.9 dex
(note the lower end of the canonical ZAEHB in Fig. 6).
On the other hand,
it is well known that the very high rotation rates in binary systems,
high metallicity, as well as the pulsations themselves,
give rise to broadening of atmospheric lines.
If these broadening effects are not accounted for in the models used
for
estimating the atmospheric parameters, the fitting procedure may
converge
towards a too high gravity to compensate.
7 Conclusions and outlook
The interest in sdB pulsators can only increase in the future, when satellite missions such as Kepler (Christensen-Dalsgaard et al. 2007) start to find numerous new sdBVs and observe them with unprecedented frequency resolution and at lower pulsation amplitudes than have ever been achieved from the ground. As we enter the age of space based asteroseismology for the faint subdwarf stars, we can expect our understanding of the incidence of pulsations and their amplitude variability over time and across the instability region to improve, as MOST (Walker et al. 2003) and CoRoT (Michel et al. 2008) have done for brighter pulsating stars. Meanwhile, a new generation of ground based instruments are coming on-line, capable of high temporal resolution and simultaneous multi-colour photometry, by splitting the light with dichroics and reading out several CCDs simultaneously. ULTRACAM has pioneered this technology, permitting reliable amplitude ratios to be measured and used directly for mode identification.
As theoretical models of the interior structure of these stars become sophisticated enough to accurately predict the pulsation frequencies (Hu et al. 2008,2009), new high-precision photometric observations will help us distinguish between the possible formation scenarios. When this happens we will finally be in position to unravel the complete evolutionary history of the hot subdwarf stars.
AcknowledgementsThe authors thank the staff at the NOT for excellent support over a decade of observations.
The time-series data presented here have been taken using ALFOSC, which is owned by the Instituto de Astrofisica de Andalucia (IAA) and operated at the Nordic Optical Telescope under agreement between IAA and the NBIfAFG of the Astronomical Observatory of Copenhagen.
The research leading to these results has received funding from the European Research Council under the European Community's Seventh Framework Programme (FP7/2007-2013)/ERC grant agreement N
227224 ( PROSPERITY), as well as from the Research Council of K.U. Leuven grant agreement GOA/2008/04.
C.R.L. acknowledges an Ángeles Alvariño contract of the regional government Xunta de Galicia.
References
- Aerts, C., Jeffery, C. S., Fontaine, G., et al. 2006, MNRAS, 367, 1317 [NASA ADS] [CrossRef] [Google Scholar]
- Aguilar-Sánchez, Y. 1998, Ph.D. Thesis, University of Bonn [Google Scholar]
- Allard, F., Wesemael, F., Fontaine, G., Bergeron, P., & Lamontagne, R. 1994, AJ, 107, 1565 [NASA ADS] [CrossRef] [Google Scholar]
- Baran, A., & Fox Machado, L. 2010, Ap&SS, accepted [arXiv:0912.4332] [Google Scholar]
- Barlow, B. N., Dunlap, B. H., Lynas-Gray, A. E., & Clemens, J. C. 2009, AJ, 138, 686 [NASA ADS] [CrossRef] [Google Scholar]
- Barlow, B. N., Dunlap, B. H., Clemens, J. C., et al. 2010, MNRAS, 403, 324 [NASA ADS] [CrossRef] [Google Scholar]
- Beers, T. C., Doinidis, S. P., Griffin, K. E., Preston, G. W., & Shectman, S. A. 1992, AJ, 103, 267 [NASA ADS] [CrossRef] [Google Scholar]
- Bessell, M. S. 1990, PASP, 102, 1181 [NASA ADS] [CrossRef] [Google Scholar]
- Billeres, M., Fontaine, G., Brassard, P., et al. 1997, ApJ, 487, L81 [NASA ADS] [CrossRef] [Google Scholar]
- Billeres, M., Fontaine, G., Brassard, P., et al. 1998, ApJ, 494, L75 [NASA ADS] [CrossRef] [Google Scholar]
- Billéres, M., Fontaine, G., Brassard, P., et al. 2000, ApJ, 530, 441 [NASA ADS] [CrossRef] [Google Scholar]
- Billères, M., Fontaine, G., Brassard, P., & Liebert, J. 2002, ApJ, 578, 515 [NASA ADS] [CrossRef] [Google Scholar]
- Bixler, J. V., Bowyer, S., & Laget, M. 1991, A&A, 250, 370 [NASA ADS] [PubMed] [Google Scholar]
- Bonanno, A., Catalano, S., Frasca, A., Mignemi, G., & Paternò, L. 2003, A&A, 398, 283 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Brassard, P., Fontaine, G., Billères, M., et al. 2001, ApJ, 563, 1013 [CrossRef] [Google Scholar]
- Charpinet, S., Fontaine, G., Brassard, P., et al. 1997, ApJ, 483, L123 [NASA ADS] [CrossRef] [PubMed] [Google Scholar]
- Charpinet, S., Fontaine, G., Brassard, P., et al. 2005, in PASPC, 14th European Workshop on White Dwarfs, ed. D. Koester, & S. Moehler, 334, 619 [Google Scholar]
- Charpinet, S., Silvotti, R., Bonanno, A., et al. 2006, A&A, 459, 565 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Christensen-Dalsgaard, J., Arentoft, T., Brown, T. M., et al. 2007, Commun. Asteroseismol., 150, 350 [NASA ADS] [CrossRef] [Google Scholar]
- Dorman, B., Rood, R. T., & O'Connell, R. W. 1993, ApJ, 419, 596 [CrossRef] [Google Scholar]
- Drechsel, H., Heber, U., Napiwotzki, R., et al. 2001, A&A, 379, 893 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Dreizler, S., Schuh, S. L., Deetjen, J. L., Edelmann, H., & Heber, U. 2002, A&A, 386, 249 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Edelmann, H. 2003, Ph.D. Thesis, University of Erlangen-Nurnberg [Google Scholar]
- Edelmann, H., Heber, U., Hagen, H.-J., et al. 2003, A&A, 400, 939 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Eisenstein, D. J., Liebert, J., Harris, H. C., et al. 2006, ApJS, 167, 40 [NASA ADS] [CrossRef] [Google Scholar]
- Fontaine, G., Brassard, P., Charpinet, S., et al. 2003, ApJ, 597, 518 [Google Scholar]
- Fontaine, G., Brassard, P., Charpinet, S., & Chayer, P. 2006, Mem. Soc. Astron. Ital., 77, 49 [Google Scholar]
- Geier, S., Nesslinger, S., Heber, U., et al. 2007, A&A, 464, 299 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Gianninas, A., Bergeron, P., & Fontaine, G. 2007, in 15th European Workshop on White Dwarfs, ed. R. Napiwotzki, & M. R. Burleigh, PASPC, 372, 577 [Google Scholar]
- Green, E. M., Fontaine, G., Reed, M. D., et al. 2003, ApJ, 583, L31 [NASA ADS] [CrossRef] [Google Scholar]
- Green, E. M., Fontaine, G., Hyde, E. A., For, B.-Q., & Chayer, P. 2008, in Hot Subdwarf Stars and Related Objects, ed. U. Heber, C. S. Jeffery, & R. Napiwotzki, PASPC, 392, 75 [Google Scholar]
- Green, R. F., Schmidt, M., & Liebert, J. 1986, ApJS, 61, 305 [NASA ADS] [CrossRef] [Google Scholar]
- Hagen, H.-J., Groote, D., Engels, D., & Reimers, D. 1995, A&AS, 111, 195 [NASA ADS] [Google Scholar]
- Han, Z., Podsiadlowski, P., Maxted, P. F. L., Marsh, T. R., & Ivanova, N. 2002, MNRAS, 336, 449 [NASA ADS] [CrossRef] [Google Scholar]
- Han, Z., Podsiadlowski, P., Maxted, P. F. L., & Marsh, T. R. 2003, MNRAS, 341, 669 [NASA ADS] [CrossRef] [Google Scholar]
- Heber, U. 1986, A&A, 155, 33 [NASA ADS] [Google Scholar]
- Heber, U. 2009, ARA&A, 47, 211 [NASA ADS] [CrossRef] [Google Scholar]
- Heber, U., Edelmann, H., Lemke, M., Napiwotzki, R., & Engels, D. 1999, in 11th European Workshop on White Dwarfs, ed. J.-E. Solheim, & E. G. Meistas, PASPC, 169, 551 [Google Scholar]
- Heber, U., Reid, I. N., & Werner, K. 2000, A&A, 363, 198 [NASA ADS] [Google Scholar]
- Heber, U., Moehler, S., Napiwotzki, R., Thejll, P., & Green, E. M. 2002, A&A, 383, 938 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Heber, U., Drechsel, H., Østensen, R., et al. 2004, A&A, 420, 251 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Hu, H., Dupret, M., Aerts, C., et al. 2008, A&A, 490, 243 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Hu, H., Nelemans, G., Aerts, C., & Dupret, M.-A. 2009, A&A, 508, 869 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Jeffery, C. S., & Saio, H. 2006, MNRAS, 372, L48 [NASA ADS] [CrossRef] [Google Scholar]
- Jeffery, C. S., & Saio, H. 2007, MNRAS, 378, 379 [NASA ADS] [CrossRef] [Google Scholar]
- Kawaler, S. D., & Hostler, S. R. 2005, ApJ, 621, 432 [NASA ADS] [CrossRef] [Google Scholar]
- Kilkenny, D. 2010, Ap&SS, submitted [Google Scholar]
- Kilkenny, D., Heber, U., & Drilling, J. S. 1988, South African Astronomical Observatory Circular, 12, 1 [Google Scholar]
- Kilkenny, D., Koen, C., O'Donoghue, D., & Stobie, R. S. 1997, MNRAS, 285, 640 [Google Scholar]
- Kilkenny, D., O'Donoghue, D., Koen, C., Lynas-Gray, A. E., & van Wyk, F. 1998, MNRAS, 296, 329 [NASA ADS] [CrossRef] [Google Scholar]
- Kilkenny, D., Koen, C., O'Donoghue, D., et al. 1999, MNRAS, 303, 525 [NASA ADS] [CrossRef] [Google Scholar]
- Kilkenny, D., Billères, M., Stobie, R. S., et al. 2002, MNRAS, 331, 399 [NASA ADS] [CrossRef] [Google Scholar]
- Kilkenny, D., Reed, M. D., O'Donoghue, D., et al. 2003, MNRAS, 345, 834 [NASA ADS] [CrossRef] [Google Scholar]
- Kilkenny, D., Stobie, R. S., O'Donoghue, D., et al. 2006, MNRAS, 367, 1603 [CrossRef] [Google Scholar]
- Kilkenny, D., O'Donoghue, D., Crause, L., et al. 2009, MNRAS, 396, 548 [NASA ADS] [CrossRef] [Google Scholar]
- Koen, C. 1998, MNRAS, 300, 567 [NASA ADS] [CrossRef] [Google Scholar]
- Koen, C. 2007, MNRAS, 377, 1275 [NASA ADS] [CrossRef] [Google Scholar]
- Koen, C., Kilkenny, D., O'Donoghue, D., van Wyk, F., & Stobie, R. S. 1997, MNRAS, 285, 645 [NASA ADS] [CrossRef] [Google Scholar]
- Koen, C., O'Donoghue, D., Kilkenny, D., et al. 1998a, MNRAS, 296, 317 [NASA ADS] [CrossRef] [Google Scholar]
- Koen, C., O'Donoghue, D., Pollacco, D. L., & Nitta, A. 1998b, MNRAS, 300, 1105 [NASA ADS] [CrossRef] [Google Scholar]
- Koen, C., O'Donoghue, D., Kilkenny, D., Stobie, R. S., & Saffer, R. A. 1999a, MNRAS, 306, 213 [NASA ADS] [CrossRef] [Google Scholar]
- Koen, C., O'Donoghue, D., Pollacco, D. L., & Charpinet, S. 1999b, MNRAS, 305, 28 [NASA ADS] [CrossRef] [Google Scholar]
- Koen, C., O'Donoghue, D., Kilkenny, D., & Pollacco, D. L. 2004, New A, 9, 565 [Google Scholar]
- Kuschnig, R., Weiss, W. W., Gruber, R., Bely, P. Y., & Jenkner, H. 1997, A&A, 328, 544 [NASA ADS] [Google Scholar]
- Lamontagne, R., Demers, S., Wesemael, F., Fontaine, G., & Irwin, M. J. 2000, AJ, 119, 241 [NASA ADS] [CrossRef] [Google Scholar]
- Lisker, T., Heber, U., Napiwotzki, R., et al. 2005, A&A, 430, 223 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Lutz, R., Schuh, S., Silvotti, R., Kruspe, R., & Dreizler, S. 2009, Commun. Asteroseismol., 159, 94 [NASA ADS] [CrossRef] [Google Scholar]
- Martin, D. C., Fanson, J., Schiminovich, D., et al. 2005, ApJ, 619, L1 [Google Scholar]
- Maxted, P. f. L., Heber, U., Marsh, T. R., & North, R. C. 2001, MNRAS, 326, 1391 [NASA ADS] [CrossRef] [Google Scholar]
- McCook, G. P., & Sion, E. M. 1987, ApJS, 65, 603 [NASA ADS] [CrossRef] [Google Scholar]
- Michel, E., Baglin, A., Weiss, W. W., et al. 2008, Commun. Asteroseismol., 156, 73 [NASA ADS] [CrossRef] [Google Scholar]
- Moehler, S., de Boer, K. S., & Heber, U. 1990, A&A, 239, 265 [NASA ADS] [Google Scholar]
- Munari, U., Sordo, R., Castelli, F., & Zwitter, T. 2005, A&A, 442, 1127 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Napiwotzki, R. 1997, A&A, 322, 256 [NASA ADS] [Google Scholar]
- Napiwotzki, R., Edelmann, H., Heber, U., et al. 2001, A&A, 378, L17 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Napiwotzki, R., Karl, C. A., Lisker, T., et al. 2004, Ap&SS, 291, 321 [NASA ADS] [CrossRef] [Google Scholar]
- Nather, R. E., Winget, D. E., Clemens, J. C., Hansen, C. J., & Hine, B. P. 1990, ApJ, 361, 309 [NASA ADS] [CrossRef] [Google Scholar]
- Noguchi, T., Maehara, H., & Kondo, M. 1980, Annals of the Tokyo Astronomical Observatory, 18, 55 [NASA ADS] [Google Scholar]
- O'Donoghue, D., Lynas-Gray, A. E., Kilkenny, D., Stobie, R. S., & Koen, C. 1997, MNRAS, 285, 657 [NASA ADS] [Google Scholar]
- O'Donoghue, D., Koen, C., Lynas-Gray, A. E., Kilkenny, D., & van Wyk, F. 1998, MNRAS, 296, 306 [NASA ADS] [CrossRef] [Google Scholar]
- Oreiro, R., Ulla, A., Pérez Hernández, F., et al. 2004, A&A, 418, 243 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Oreiro, R., Pérez Hernández, F., Ulla, A., et al. 2005, A&A, 438, 257 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Oreiro, R., Pérez Hernández, F., Østensen, R., et al. 2007, A&A, 461, 585 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Oreiro, R., Østensen, R. H., Green, E. M., & Geier, S. 2009, A&A, 496, 827 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Østensen, R. 2009, Commun. Asteroseismol., 159, 75 [Google Scholar]
- Østensen, R., & Solheim, J.-E. 2000, Baltic Astron., 9, 411 [Google Scholar]
- Østensen, R., Heber, U., Silvotti, R., et al. 2001a, A&A, 378, 466 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Østensen, R., Solheim, J.-E., Heber, U., et al. 2001b, A&A, 368, 175 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Østensen, R., Oreiro, R., Drechsel, H., et al. 2007, in 15th European Workshop on White Dwarfs, ed. R. Napiwotzki, & M. R. Burleigh, PASPC, 372, 483 [Google Scholar]
- Østensen, R. H. 2000, Ph.D. Thesis, University of Tromsø [Google Scholar]
- Østensen, R. H. 2004, Ap&SS, 291, 263 [NASA ADS] [CrossRef] [Google Scholar]
- Østensen, R. H. 2006, Baltic Astron., 15, 85 [Google Scholar]
- O'Toole, S. J., Heber, U., & Benjamin, R. A. 2004, A&A, 422, 1053 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Piccioni, A., Bartolini, C., Bernabei, S., et al. 2000, A&A, 354, L13 [NASA ADS] [Google Scholar]
- Ramsay, G., Napiwotzki, R., Hakala, P., & Lehto, H. 2006, MNRAS, 371, 957 [NASA ADS] [CrossRef] [Google Scholar]
- Randall, S. K., Green, E. M., van Grootel, V., et al. 2007, A&A, 476, 1317 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Randall, S. K., Calamida, A., & Bono, G. 2009a, A&A, 494, 1053 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Randall, S. K., Van Grootel, V., Fontaine, G., Charpinet, S., & Brassard, P. 2009b, A&A, 911 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Reed, M. D., & Stiening, R. 2004, PASP, 116, 506 [NASA ADS] [CrossRef] [Google Scholar]
- Reed, M. D., Eggen, J. R., Zhou, A.-Y., et al. 2006, MNRAS, 369, 1529 [NASA ADS] [CrossRef] [Google Scholar]
- Reed, M. D., O'Toole, S. J., Terndrup, D. M., et al. 2007a, ApJ, 664, 518 [NASA ADS] [CrossRef] [Google Scholar]
- Reed, M. D., Terndrup, D. M., Zhou, A.-Y., et al. 2007b, MNRAS, 378, 1049 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Reed, M., Tendrup, D., Østensen, R., et al. 2010, Ap&SS, submitted [Google Scholar]
- Saffer, R. A., Bergeron, P., Koester, D., & Liebert, J. 1994, ApJ, 432, 351 [NASA ADS] [CrossRef] [Google Scholar]
- Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500, 525 [NASA ADS] [CrossRef] [Google Scholar]
- Schuh, S., Huber, J., Dreizler, S., et al. 2006, A&A, 445, L31 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Silvotti, R., Solheim, J.-E., Gonzalez Perez, J. M., et al. 2000, A&A, 359, 1068 [NASA ADS] [Google Scholar]
- Silvotti, R., Janulis, R., Schuh, S. L., et al. 2002a, A&A, 389, 180 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Silvotti, R., Østensen, R., Heber, U., et al. 2002b, A&A, 383, 239 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Silvotti, R., Bonanno, A., Bernabei, S., et al. 2006, A&A, 459, 557 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Silvotti, R., Schuh, S., Janulis, R., et al. 2007, Nature, 449, 189 [NASA ADS] [CrossRef] [PubMed] [Google Scholar]
- Silvotti, R., Handler, G., Schuh, S., Castanheira, B., & Kjeldsen, H. 2009, Commun. Asteroseismol., 159, 97 [NASA ADS] [CrossRef] [Google Scholar]
- Solheim, J.-E., & Østensen, R. 2006, Baltic Astron., 15, 231 [NASA ADS] [Google Scholar]
- Solheim, J.-E., Østensen, R., Silvotti, R., & Heber, U. 2004, Ap&SS, 291, 419 [NASA ADS] [CrossRef] [Google Scholar]
- Stepanian, J. A., Green, R. F., Foltz, C. B., et al. 2001, AJ, 122, 3361 [NASA ADS] [CrossRef] [Google Scholar]
- Stobie, R. S., Kawaler, S. D., Kilkenny, D., O'Donoghue, D., & Koen, C. 1997a, MNRAS, 285, 651 [NASA ADS] [CrossRef] [Google Scholar]
- Stobie, R. S., Kilkenny, D., O'Donoghue, D., et al. 1997b, MNRAS, 287, 848 [NASA ADS] [CrossRef] [Google Scholar]
- Stoughton, C., Lupton, R. H., Bernardi, M., et al. 2002, AJ, 123, 485 [NASA ADS] [CrossRef] [Google Scholar]
- Telting, J. H., & Østensen, R. H. 2004, A&A, 419, 685 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Telting, J. H., Østensen, R. H., Oreiro, R., et al. 2010, Ap&SS, submitted [Google Scholar]
- Vuckovic, M., Kawaler, S. D., O'Toole, S., et al. 2006, ApJ, 646, 1230 [NASA ADS] [CrossRef] [Google Scholar]
- Vuckovic, M., Aerts, C., Østensen, R., et al. 2007, A&A, 471, 605 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
- Walker, G., Matthews, J., Kuschnig, R., et al. 2003, PASP, 115, 1023 [NASA ADS] [CrossRef] [Google Scholar]
- Winter, C. 2006, Ph.D. Thesis, Armagh Observatory & The Queen's University of Belfast [Google Scholar]
- Wisotzki, L., Koehler, T., Groote, D., & Reimers, D. 1996, A&AS, 115, 227 [NASA ADS] [Google Scholar]
- Woudt, P. A., Kilkenny, D., Zietsman, E., et al. 2006, MNRAS, 371, 1497 [NASA ADS] [CrossRef] [Google Scholar]
- Young, A. T. 1993, The Observatory, 113, 41 [NASA ADS] [Google Scholar]
- Zacharias, N., Monet, D. G., Levine, S. E., et al. 2005, VizieR Online Data Catalog, 1297, 0 [Google Scholar]
- Zhou, A.-Y., Reed, M. D., Harms, S., et al. 2006, MNRAS, 367, 179 [NASA ADS] [CrossRef] [Google Scholar]
Footnotes
- ... Telescope
- Full Table 5 is only available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/513/A6
- ...
V361 Hya
- Translations between the old survey names and the variable star names used here are provided in Table 9.
- ... (mma
- The mma units are 10-3 of the Fourier amplitudes of a light-curve in normalised intensity units. One mma translates to a peak-to-peak amplitude of two millimodulation intensity (mmi) units in the light curve.
- ...
variability
- Actually, one more of the fourteen stars in this list of misclassified WDs is included in our sample: Ton S 155 appears in Edelmann (2003) as HE 0021-2326.
- ... ones
- Filter details with transmission curves for all NOT filters can be found on their web-pages; http://www.not.iac.es/.
- ...
HE 0230-4323
- Originally reported as ``an unusual hot subdwarf pulsator'' by Koen (2007) but now confirmed to be a more regular sdBV by Kilkenny (priv. comm.).
All Tables
Table 1: Observation runs at the Nordic Optical Telescope contributing to this survey.
Table 2: Analysis for eight stars from the HS survey, not included in Edelmann et al. (2003).
Table 3: The stars from the SDSS survey observed during this programme, excluding the pulsators listed in Table 4.
Table 4: Spectral decomposition of the binary pulsators from the SDSS sample.
Table 5: Truncated version of the catalogue of sdB stars observed during our search programme.
Table 6: Details from the discovery runs for the four new sdBVs.
Table 7: Pulsation properties for the four new pulsators.
Table 8: Comparison with the results of other surveys for pulsating sdBs (upper part) and our own survey divided into subsamples (lower part).
Table 9: All short period sdBV stars published to date.
All Figures
![]() |
Figure 1:
Top panel: photographic B
magnitudes of the targets
as function of |
Open with DEXTER | |
In the text |
![]() |
Figure 2: Amplitude spectrum of HE 2151-1001. The lower panel shows the residual spectrum after prewhitening the two frequencies listed in Table 7 from the light-curve. The dashed line indicates 3.7 times the noise level. |
Open with DEXTER | |
In the text |
![]() |
Figure 3:
Amplitude spectrum of HS 2125+1105. The lower panel
shows residual spectrum after subtracting one frequency from the
original light-curve.
The dashed line indicates 3.7 times the noise level.
The peak at 7.3 mHz is just above the 3.7 |
Open with DEXTER | |
In the text |
![]() |
Figure 4: Amplitude spectrum of PG 1033+201. The lower panel shows the residual spectrum after prewhitening the two frequencies listed in Table 7 from the light-curve. We have also plotted the amplitude spectrum of the first half (red dashed curve) and second half (blue short dash, long dash curve), to demonstrate that the main peak is present with the same amplitude in both halves of the dataset. The dashed line indicates 3.7 times the noise level. |
Open with DEXTER | |
In the text |
![]() |
Figure 5:
Amplitude spectrum of HE 1450-0957 ( top)
and residual amplitude spectrum
after subtracting the main peak ( bottom), with 3.7 |
Open with DEXTER | |
In the text |
![]() |
Figure 6:
a)
|
Open with DEXTER | |
In the text |
Copyright ESO 2010
Current usage metrics show cumulative count of Article Views (full-text article views including HTML views, PDF and ePub downloads, according to the available data) and Abstracts Views on Vision4Press platform.
Data correspond to usage on the plateform after 2015. The current usage metrics is available 48-96 hours after online publication and is updated daily on week days.
Initial download of the metrics may take a while.