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NEWS FROM GALACTIC XRBs

Cloudy weather in SgXRBs

The mystery of the missing population of Be/BH
XRBs

Spins of compact objects in XRBs
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CLOUDY WEATHER IN SgXRBs
Clumpy winds in SgXRBs

Cyg X-3 SFXTs

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Cyg X-3 (Szostek Zdziarski, 2008)
Analysing X-ray spectra from Beppo SAX authors
found that strong wind from WR component must be
very clumpy. The shapes of the spectra imply
that it consist of two phases ? hot tenuous
plasma carrying most of the wind mass ? cool
dense clumps (filling factor lt 0.01)
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SFXTs (Negueruela et al., 2008)
Authors analyse the variability of SgXRBs (
wind fed systems)
Classical SgXRBs are persistent systems New
cathegory of SFXTs show very rapid variability
They display outbursts with a rise time scale of
tens of minutes lasting a few hours
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Two sources (XTE J1739-3020 IGR J17544-2619)
increased their X-ray luminosity by a factor gt
100 in a few minutes !
Such increase cannot be explained by an orbital
motion through a smooth medium
MEDIUM IS NOT SMOOTH !
Clumpy winds will help
Clumpy wind model of Oskinova et al. (2007)
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vw (r) v8 (1- R/r)ß
raccr 2 G Mx/vrel2
vrel2 vw2 vorb2
Ncl 4p ?t/L03 L0 - porosity length
?t ?r/vw ?t 2 raccr/vw
when r? then raccr? and vw? Ncl??
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Number of clumps in a ring of width 2 raccr and
height 2 raccr
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The change of the slope of Ncl(r) at r 2 R is
very abrupt
It defines two regimes of accretion
The inner regime, where NS almost always sees a
clump ( clasical SgXRBs)
The outer regime, where NS very rarely sees a
clump ( SFXTs)
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The Mystery of the Missing Population of BeBlack
Hole X-Ray Binaries
At present, 117 Be/NS binaries are known in the
Galaxy and the Magellanic Clouds, but not a
single Be/BH binary was found so far
117 0 !!!
In all XRBs the ratio of NSs to BHs is 41
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Short tutorial on Be XRBs

These systems composed of a Be star and a compact
object form the most numerous class of XRBs in
our Galaxy

At present, 117 systems of that type are known in
the Galaxy and the Magellanic Clouds

In all systems (discovered so far), the compact
object is a NS.
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BeNS system consists of a NS orbiting a Be type
star on a rather wide (orbital periods in the
range of 10 to 300 days), frequently
eccentric, orbit. NS has a strong magnetic field
and, in vast majority of cases, is observed as an
X-ray pulsar (with the spin periods in the range
of 34 ms to 6000 s). The Be component is deep
inside its Roche lobe. This is a distinct
property of Be XRBs. In all other types of XRBs,
the optical component always fills or almost
fills its Roche lobe (even if the accreted matter
is supplied by the winds).
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X-ray emission from Be XRBs
X-Ray emission (with a few exceptions) has
distinctly transient nature with rather short
active phases (a flaring behaviour). There are
two types of flares, which are classified as Type
I outbursts (smaller and regularly repeating) and
Type II outbursts (larger and irregular).
Type I bursts are observed in systems with highly
eccentric orbits. They occur close to periastron
passages of NS. They are repeating at intervals
Porb.
Type II bursts may occur at any orbital phase.
They are correlated with the disruption of the
exretion disc around Be star (as observed in Ha
line). They repeat on time scale years.
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Type I burst
trec Porb
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Type II burst
trec years
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At present, 117 Be/NS binaries are known in the
Galaxy and the Magellanic Clouds (which is almost
a half of the total number of the known NS
binaries), but not a single Be/BH binary was
found so far (although 58 BH candidate systems
are known).
This disparity (117 Be/NS type systems out of 252
known NS XRBs vs. not a single Be/BH type system
among 58 known BH XRBs) called the attention of
the researchers already for some time.
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In particular, Zhang et al. (2004) noted that,
according to stellar population synthesis
calculations by Podsiadlowski et al. (2003), BH
binaries are formed predominantly with relatively
short orbital periods (Porb lt 10 days). If this
is the case, then, according to Zhang et al., the
excretion disc truncation mechanism (Artymowicz
Lubow, 1994) might be so efficient, that the
accretion rate is very low and the system remains
dormant (and therefore invisible) for almost all
the time. One should note, however, that
Podsiadlowski et al. considered, essentially, BH
systems with Roche lobe filling secondaries,
which definitely is not the case of Be XRBs.
Therefore, their results are not relevant for the
case of Be/BH XRBs.
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Stellar population synthesis (SPS) calculations
Sadowski, Ziólkowski, Belczynski Bulik carried
out calculations using the STAR TRACK code
(Belczynski, Kalogera Bulik, 2002 Belczynski
et al., 2008). State of art SPS code (Vicki
Kalogera)
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Definition of a Be star (for the purpose of SPS
calculations)
The primary property of Be stars distinguishing
them from other B stars is rapid rotation. All
other properties (in particular, the presence of
an excretion disc, which permits the efficient
accretion on the compact companion) are the
consequences of the fast rotation.
It is not clear how Be stars achieved their fast
rotation (although different hypothesis like
rapid rotation at birth or spin-up due to binary
mass transfer are advanced see e.g. McSwain
Gies, 2005). The fraction of Be stars among all B
stars is similar for single stars and for those
in binary systems (one quarter to one third).
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For simplicity, we assumed that one quarter of
all B stars are always Be stars and that these
stars are always efficient mass donors,
independently of the size of the binary orbit (as
is, in fact, observed in Be/NS XRBs).
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The preliminary results of our calculations,
showing the expected ratio of the number of Be/NS
binaries to the number of Be/BH binaries are
shown in Fig. 1.
Fig. 1 clearly shows that the expected numbers of
Be/NS and Be/BH binaries should be roughly
comparable. The estimated masses of observed Be
stars cover the range from 2.3 MSUN (Lejeune
Schaerer, 2001) to 25 MSUN (McSwain Gies,
2005).
Independently of the value of the minimum mass
assumed for a Be star, it is obvious that,
according to our calculations, Be/NS systems
should not outnumber Be/BH systems by more than a
factor of about 2.5.
What is the cause of so dramatic discrepancy ?
The answer is shown in Fig. 2
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NS
BH
BH
NS
MBe,min 3 MSUN
MBe,min 8 MSUN
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According to our calculations, the distribution
of the orbital periods is completely different
for Be/NS and Be/BH systems. Within the orbital
period range where Be XRBs are found ( 10 to
300 days), Be systems are formed predominantly
with a NS component. The ratio of the expected
number of Be/NS systems to the expected number of
Be/BH systems is, for this orbital period range,
larger than 50. The systems with a BH component
are formed predominantly with much longer orbital
periods. Such systems are very difficult to
detect, both due to very long orbital periods and
due to, probably, very low luminosities (the
accretion at such large orbital separations must
be very inefficient).
CAUTION there might be also other factors !
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CAUTION there might be also other factors !
Another possible factor may be related to the
previous evolution of a Be star. If, indeed, a
B star must be a member of a binary system and
undergo a mass transfer in order to become a Be
star, then one can imagine that the systems
composed of a Be star and a relatively less
massive companion (which collapses to a NS)
remain bound, while those composed of a Be star
and a relatively more massive companion (which
collapses to a BH) are disrupted in the process
of a supernova explosion.
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SPINS of COMPACT OBJECTS in XRBs

BHs
a (lt 0.26, gt 0.98)

NSs
Pspin (0.89 ms, 104 s)
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SPINS of BHs

• Spins of accreting BHs could be deduced from
• X-ray spectra (continua)
• 2. X-ray spectra (lines)
• 3. kHz QPOs

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Specific angular momentum for circular orbits
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X-RAY SPECTRA
Zhang et al. (1997) GRO J1655-40 a
0.93
GRS 1915105 a 1.0 Gierlinski et al.
(2001) GRO J1655-40 a 0.68
0.88 McClintock et al. (2006) LMC X-3
a lt 0.26
GRO J1655-40 a 0.65 0.80
4U 1543-47
a 0.70 0.85
GRS 1915105 a gt 0.98
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SPECTRAL LINES
MODELING THE SHAPE OF Fe Ka LINE
Miller et al. (2004) GX339-4 a
0.8 0.9 Miller et al. (2005) GRO J1655-40
a gt 0.9
XTE J1550-564 a gt 0.9 Miller et al.
(2002) XTE J1650-500 a 1.0

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Miller (2004)
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NEW ERA OF PRECISION
Reis et al. (2008) determined the spin of BH in
GX 339-4 (from RXTE XMM)
rin 2.020.02-0.06 rg at very high
state rin 2.040.07-0.02 rg at
low/hard state
a 0.935 0.02 at 90 confidence
(!)
Miller et al. (2008) did this from Suzaku XMM
a 0.93 0.05
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GX 339-4
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kHz QPOs
Name ?QPO
Hz MBH MSUN
GRO J1655-40 300 23
6.3 0.5
450 20
XTE J1550-564 184 26
10.5 1.0
272 20 H
1743-322 166 8
240 3 GRS
1915105 41 1
14 4.4
67 5
113
164 2
4U 1630-472 184 5
XTE J1859226 193 4
9 1 XTE
J1650-500 250
5 2
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a 0.7 0.99
a 0.7 0.99
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BHs SPINS (summary)
(1) GRS 1915105 has a rotation close to nearly
maximal spin ( a gt0.98) (2) several other
systems (GRO J1655-40, 4U 1543-47, XTEJ1550-564,
XTE J1650-500 and GX 339-4 have large spins (a
0.65) (3) not all accreting black holes have
large spins (robust result a lt 0.26 for LMC X-3)
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? almost all BHs for which there is a reliable
estimate of high spin (with the sole exception of
4U 1543-47) are microquasars ? perhaps, it is
so, that all microquasar BHs have high spins,
while other accreting BHs might or might not have
high spins.
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SPINS of NSs

• Spins of accreting NSs could be deduced from
• X-ray pulses

2. kHz QPOs (especially during the tails of X-ray
bursts)
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Spins determined from X-ray pulses
They are, at present, in the range 1.67
ms IGR J1002915934 to 10 008 s 2S
0114650
(135 X-ray pulsars known so far (among them 9
millisecond pulsars)
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Spins determined from kHz QPOs
There are two types of kHz QPOs

burst QPOs
?B ?SPIN

pair QPOs
??PAIR ?SPIN or ??PAIR 0.5 ?SPIN
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Burst QPOs reflect the true spin frequency
This is known from three sources which display
simultaneously X-ray pulses
XTE J1814-338 ?SPIN 314 Hz
SAX J1808.4-3658 ?SPIN 401 Hz
Aql X-1 ?SPIN 550 Hz
Spin periods from burst QPOs are known for 13
(14) other bursters. These spin periods are in
the range 1.62 to 10.5 ms
New discovery XTE J1739-285 PSPIN 0.89 ms
(Kaaret et al., 2007)
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Pair QPOs are less obvious to interpret
XTE J1807-294 ??PAIR 200 Hz
( ?SPIN)
SAX J1808.4-3658 ??PAIR 200 Hz
( 0.5 ?SPIN)
Aql X-1 ??PAIR 275 Hz
( 0.5 ?SPIN)
fast rotators (?SPIN 400 Hz) seem to have
??PAIR 0.5 ?SPIN, while slow rotators (?SPIN lt
400 Hz) have ??PAIR ?SPIN (sample of 7
sources)
13 additional spin determinations
Parametric epicyclic resonance theory (Abramowicz
Kluzniak, since 2001) seems to be able to
explain this behaviour