Transcript News from Galactic X

NEWS FROM GALACTIC XRBs
• „Cloudy” weather in SgXRBs
mystery of the missing population of
• The
Be/BH XRBs
Spins of compact objects in XRBs
•
„CLOUDY” WEATHER IN SgXRBs
„Clumpy” winds in SgXRBs
• Cyg X-3
• SFXTs
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 < 0.01)
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
Two sources (XTE J1739-3020 & IGR J17544-2619)
increased their X-ray luminosity by a factor > 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)
vw (r) = v∞ (1- R*/r)β
raccr ≈ 2 G Mx/vrel2
vrel2 = vw2 + vorb2
Ncl = 4π Δt/L03
L0 - porosity length
Δt = Δr/vw
Δt ≈ 2 raccr/vw
when
r↑ then raccr↓ and vw↑
Ncl↓↓
Number of clumps in a ring of width 2 raccr and height 2 raccr
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)
The Mystery of the Missing Population of
Be+Black 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 4:1
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.
Be+NS 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).
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 Hα line). They repeat on time scale ~
years.
Type I burst
trec ~ Porb
Type II burst
trec ~ years
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.
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 < 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.
Stellar population synthesis (SPS)
calculations
Sądowski, Ziółkowski, Belczyński & Bulik carried out
calculations using the STAR TRACK code (Belczynski,
Kalogera & Bulik, 2002; Belczynski et al., 2008).
„State of art” SPS code (Vicki Kalogera)
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).
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).
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
BH
NS
BH
NS
MBe,min = 3
MBe,min = 8 MSUN
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 !
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.
SPINS of COMPACT OBJECTS in XRBs
•
BHs
a* € (< 0.26, > 0.98)
• NSs
Pspin € (0.89 ms, 104 s)
SPINS of BHs
Spins of accreting BHs could be deduced
from:
1. X-ray spectra (continua)
2. X-ray spectra (lines)
3. kHz QPOs
Specific angular momentum for circular orbits
X-RAY SPECTRA
Zhang et al. (1997):
GRO J1655-40
a* = 0.93
GRS 1915+105
a* ≈ 1.0
Gierliński et al. (2001): GRO J1655-40
a* = 0.68 ÷ 0.88
McClintock et al. (2006): LMC X-3
a* < 0.26
GRO J1655-40
a* = 0.65 ÷ 0.80
4U 1543-47
a* = 0.70 ÷ 0.85
GRS 1915+105
a* > 0.98
SPECTRAL LINES
MODELING THE SHAPE OF Fe Kα LINE
Miller et al. (2004): GX339-4
a* ≥ 0.8 ÷ 0.9
Miller et al. (2005): GRO J1655-40
a* > 0.9
XTE J1550-564
a* > 0.9
Miller et al. (2002): XTE J1650-500
a* ≈ 1.0
Miller (2004)
NEW ERA OF PRECISION
Reis et al. (2008) determined the spin of BH in GX
339-4 (from RXTE & XMM):
rin = 2.02+0.02-0.06 rg
at very high state
rin = 2.04+0.07-0.02 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
GX 339-4
kHz QPOs
Name
νQPO [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 1915+105
41 ± 1
14 ± 4.4
67 ± 5
113
164 ± 2
4U
1630-472
184 ± 5
XTE
J1859+226
193 ± 4
9±1
XTE
J1650-500
250
5±2
a*
a ≈≈ 0.7 ÷ 0.99
BHs SPINS
(summary)
(1) GRS 1915+105 has a rotation close to nearly maximal
spin ( a* >0.98)
(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* < 0.26 for LMC X-3)
● 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.
SPINS of NSs
Spins of accreting NSs could be deduced
from:
1. X-ray pulses
2. kHz QPOs (especially during the tails of Xray bursts)
Spins determined from X-ray pulses
They are, at present, in the range:
1.67 ms
to 10 008 s
IGR J100291+5934
2S 0114+650
(135 X-ray pulsars known so far (among them 9
millisecond pulsars)
Spins determined from kHz QPOs
There are two types of kHz QPOs:
•
burst QPOs
•
pair QPOs
νB = νSPIN
ΔνPAIR = νSPIN
or
ΔνPAIR = 0.5 νSPIN
Burst QPOs reflect the true spin frequency
This is known from three sources which display
simultaneously X-ray pulses:
XTE J1814-338
SAX J1808.4-3658
Aql X-1
νSPIN = 314 Hz
νSPIN = 401 Hz
νSPIN = 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)
Pair QPOs are less obvious to interpret
XTE J1807-294
SAX J1808.4-3658
Aql X-1
ΔνPAIR ≈ 200 Hz
ΔνPAIR ≈ 200 Hz
ΔνPAIR ≈ 275 Hz
(= νSPIN)
(= 0.5 νSPIN)
(= 0.5 νSPIN)
fast rotators (νSPIN ≥ 400 Hz) seem to have ΔνPAIR ≈ 0.5 νSPIN,
while slow rotators (νSPIN < 400 Hz) have ΔνPAIR ≈ νSPIN
(sample of 7 sources)
13 additional spin determinations
Parametric epicyclic resonance theory (Abramowicz &
Kluzniak, since 2001) seems to be able to explain this
behaviour