NEWS FROM GALACTIC XRBs

„Cloudy” weather in SgXRBs••

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The mystery of the missing population of Be/BH XRBs

Spins of compact objects in XRBs

„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 < 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 = vw

2 + 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 (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).

X-ray emission from Be XRBs

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.

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)

Stellar population synthesis (SPS) calculations

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).

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 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.

The answer is shown in Fig. 2

NSBH

NS BH

MBe,min = 3 MSUN 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 !

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.

CAUTION: there might be also other factors !

SPINS of COMPACT OBJECTS in XRBs

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BHs

NSs

a* € (< 0.26, > 0.98)

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)

Reis et al. (2008) determined the spin of BH in GX 339-4 (from RXTE & XMM):

a* = 0.935 ± 0.02 at 90 % confidence (!)

rin = 2.02+0.02-0.06 rg at very high state

rin = 2.04+0.07-0.02 rg at low/hard state

Miller et al. (2008) did this from Suzaku & XMM:

a* = 0.93 ± 0.05

NEW ERA OF PRECISION

GX 339-4

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

kHz QPOs Name νQPO [Hz] MBH [MSUN]

a ≈ 0.7 ÷ 0.99a* ≈ 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 X-ray bursts)

Spins determined from X-ray pulses

They are, at present, in the range:

1.67 ms IGR J100291+5934

to 10 008 s 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

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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 νSPIN = 314 HzSAX J1808.4-3658 νSPIN = 401 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)

Aql X-1 νSPIN = 550 Hz

Pair QPOs are less obvious to interpret

SAX J1808.4-3658 ΔνPAIR ≈ 200 Hz (= 0.5 νSPIN)

XTE J1807-294 ΔνPAIR ≈ 200 Hz (= 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)

Aql X-1 ΔνPAIR ≈ 275 Hz (= 0.5 νSPIN)

Parametric epicyclic resonance theory (Abramowicz & Kluzniak, since 2001) seems to be able to explain this behaviour

13 additional spin determinations