Chapter 3 Observational Characteristics of X-ray Binaries

Chapter 3 Observational characteristics of x-ray binaries

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Chapter 3 Observational Characteristics of X-ray Binaries

1. Introduction More than 300 galactic X-ray

binaries, with Lx ~ 1034 - 1038 ergs-1.

Concentrated towards the galactic center and the galactic plane, some in globular clusters.

The all-sky map generated by Uhuru


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There are also extragalactic X-ray sources discovered in LMC, SMC and other galaxies.

The Chandra's image of M83 shows numerous point-like neutron star and black hole X-ray sources

scattered throughout the disk of this spiral galaxy (http://chandra.harvard.edu).


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The standard model for galactic X-ray sources was first suggested by Salpeter (1964), Zeldovich (1964), Zeldovich & Guseynov (1965): accreting neutron stars or black holes in binary systems.


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This model was confirmed by the discovery of the pulsating X-ray binary source Cen X-3 in 1971. Regular pulsations with a period of

4.84 s (neutron star’s rotational period).

Both the X-ray eclipses and the

pulse period have a same modulation with the period of 2.087 days (orbital period).


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The mass function 3 3opt 7 3 2 3/ 2

X orb2X opt

sin( ) 1.035 10 (1 ) 15.6

( )M i

f M K P e MM M

⊙

implies a lower limit of 15.6 M⊙ for the mass of the companion star of Cen X-3. Thus the X-ray pulsar is moving in a very close orbit around a massive companion star.


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Classification of XRBs

Reig 2011


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Classification of X-ray binaries XRBs are conventionally divided into two classes: high-mass X-ray binaries (HMXBs) and low-mass X-ray binaries (LMXBs)


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Characteristics HMXBs LMXBs

X-ray spectra Relatively hard X-ray spectra with kT ≥ 15 keV in exponential fits or a power law energy index of ~0 –1

Soft X-ray spectra with kT ~ 5-10 keV in exponential fits

Time variability Often regular pulsations; no X-ray bursts; often X-ray eclipses.

Often X-ray bursts and quasi periodic oscillations; regular pulsations in a few cases


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Optical counterparts

Massive (>10 M ⊙ ) and early type (O and early B) stars; Lopt/Lx ~0.1-1000

Faint blue optical counterparts (late type or degenerate stars) ; Lopt/Lx ~0.001-0.01 ( low mass, <1 M⊙)

Optical spectrum

Stellar like Reprocessing

Orbital period 1 d - 1 yr 10 min – 10 d


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Accretion disk Usually no, sometimes small

Yes

Space distribution

Concentrated towards the Galactic plane; young stellar population, age < 107 yr

Concentrated towards the Galactic center; fairly wide spread around the Galactic plane; old stellar population, age (5-15)109yr


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2. Neutron star high-mass X-ray binaries

Supergiant/X-ray binaries Be/X-ray binaries


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Characteristics SG/X-ray binaries Be/X-ray binaries Optical counterparts

Evolved (giant, Of or blue super giant) stars, with R~10-30R⊙, Lopt > 105L⊙ and M > 20M⊙

Un-evolved stars of spectral type O9Ve to B2Ve characterized by emission lines (predominantly the Balmer lines), with R<5-10 R⊙, Lopt <3×104L⊙ and M ~8-20 M⊙

Orbits Nearly Circular orbits with Porb~1.4-10 days

Eccentric orbits with Porb~10 days to 1 year


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Time variability

X-ray eclipse and periodic ellipsoidal light variations in many cases

Rare X-ray eclipse and periodic ellipsoidal light variations

Mass transfer and X-ray emission

The optical stars fill, or nearly fill their Roche lobes; Persistent and transient X-ray Sources; Mass transfer is due to incipient Roche lobe overflow and/or capture of spherical wind

The Be stars underfill their Roche lobes; Often transient X-ray sourcesi; The outbursts are due to mass accretion from the equatorial wind


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X-ray luminosities

Group of HMXBs Average log Lx (ergs-1)

SG/X 36.7±0.9 (0.3) Be/X in outburst 36.6±0.8 (0.3) Be/X in quiescence 33.3±1.2 (0.3) Magellanic sources 38.4±0.5 (0.2)


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Pulse period distribution The pulse periods are distributed between 0.069 s and 104 s. So far about 40 systems both reliable orbital and spin periods are

known. In the Ps-Porb diagram (Corbet diagram) LMXBs, supergiant

HMXBs and Be/X-ray binaries distribute in different regions. For Be/X-ray binaries there is a strong correlation between Ps and Porb.


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The Corbet diagram

Townsend et al. 2011


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INTEGRAL X-ray Binaries The INTEGRAL observatory has performed a detailed survey of the Galactic plane. The most important result is the discovery of many new high-energy sources – concentrated in the Galactic plane, and in the Norma arm. Many of them are HMXBs hosting a neutron star orbiting around an O/B companion, in most cases a supergiant.

Bodaghee et al. 2007


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They are divided into two classes: (1) Obscured HMXBs exhibiting a huge intrinsic and local extinction;

Bodaghee et al. 2007


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(2) Supergiant Fast X-ray Transients (SFXTs) exhibiting fast and transient outbursts They sporadically undergo bright flares up to peak luminosities of 1036 –1037 ergs-1, with duration of a few hours for each single flare. In quiescence the luminosity can be as low as 1032 ergs-1.

Romano et al. 2009


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Pulse period changes Long term monitoring of the pulse periods of X-ray pulsars have revealed three types of behavior: Disk-fed sources: a secular decrease with time (spin up) with

erratic variations around the trend.

Bildsten et al. 1997


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Observations with CGRO and other telescopes show that the long-term spin-up trend is in fact the result of alternating (≤10-100 days ~ 10-20 yrs) intervals of steady spin-up and spin-down with a magnitude several times larger than the long-term spin-up trend! This has challenged the traditional model for neutron star – accretion disk interaction.


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Persistent wind-fed sources: random walk or secular spin-down. Probably they are also undergoing rapid switching behavior.


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Transient sources: rapid spin-up during (giant) outbursts and spin-down during quiescence..


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Pulse profiles There are dramatic differences in pulse shape and amplitude

between one pulsar and another. In some cases, the pulse profile shows a dependence on energy and luminosity.


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The pulse profiles of many pulsars can be reproduced simply by a beam pattern.

If the angles of the magnetic axis and the light of sight with respect to the rotation axis are and , respectively, then depending on whether , or, one or both magnetic poles are visible. This will lead to a single pulse or double pulses.


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If the beam has a maximum either along or perpendicular to the magnetic axis, the beam can be described as a pencil beam or a fan beam respectively.

In some cases an offset of the magnetic axis and/or two polar cap

emission regions with different sizes are required to give an asymmetric pulse profile.


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Pulsar spectra Continuum: The energy spectrum of an X-ray pulsar is characterized by

cfc

c

EEEEEENEEEN

EN ]/)(exp[

)(

0

0

where N is the photon number, the power-law energy index Typical values for the cutoff energy Ec and Ef both lie in the range 10-20 keV.


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Other spectrum components: Soft X-ray absorption by circumstellar matter hydrogen column density NH

distance. Fluorescent iron lines at ~6.4-7.0 keV, due to the fluorescent emission from less ionized iron in the cool, circumstellar matter.

Becker & Wolff 2007


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Cyclotron lines due to resonant scattering of the line of sight X-ray photons against electrons embedded in magnetic fields detected at energies ~ 7-120 keV (E0=11.6B12/(1+z) keV), corresponding to neutron star surface magnetic field strengths of (1-10)×1012

G.

The spectrum of Her X-1 with a cyclotron

absorption line feature at around ~38 keV.


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The cyclotron line energy vs. the cutoff energy


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3. Neutron Star Low-mass X-ray binaries (LMXBs)

LMXBs

HMXBs


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Space distribution The LMXBs contain the globular cluster X-ray sources, X-ray

bursters, soft X-ray transients, and the bright galactic bulge X-ray sources.

There are tens of luminous (>1035ergs-1) X-ray sources located in globular clusters, which are two orders of magnitude more than expected from the total mass in globular clusters relative to that in the Galaxy (Katz 1975).

Tidal capture or exchange encounter in cluster core may favor the production of X-ray binaries (Fabian, Pringle, and Rees 1975).


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(Pooley et al. 2003 ApJ 591, L131)


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Orbital period distribution The orbital periods range from ~12 min to 17 days. Comparison of the distribution of orbital periods of LMXBs and

cataclysmic variables (CVs) shows that, there are very few CVs with Porb< 1 hr, and within the period gap between 2 and 3 hrs; in the case of the LMXBs the period gap may extend down to < 1 hr (why?).


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X-ray orbital light curves Eclipses in LMXBs are fewer than expected if the mass transfer

occurs via a thin accretion disk. Thick accretion disks block the X-ray sources in edge-on systems. The accretion disk is thick because the incoming gas stream

creates turbulence at the outer edge of the disk.


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Dips in light curves are ascribed to material that is projected up above the disk plane by a splash point where the gas stream hits the accretion disk.


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For some systems that are viewed almost edge on, the compact X-ray source is hidden by the disk, but X-rays are still seen because they are scattered in a photo-ionized corona above the disk (accretion disk corona, ADC). This makes the source appear extended and results in partial eclipse.


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The observed properties of an LMXB depend on the viewing angle


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Viewing angle Properties Low inclination (<70) No X-ray dips or eclipses, optical

modulation (from the heated companion).

Intermediate inclination (70-80)

Periodic dips, sometimes brief eclipses.

High inclination (>80) Partial eclipses of ADC.


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Spectral emission NS LMXBs are divided into Z and atoll sources, according to their

patterns in the X-ray “colour-colour” diagrams (CDs).

Soft color: log of count ratio (3.5-6.4)/(2.0-3.5) keV

Hard color: log of count ratio (9.7-16.0)/(6.4-9.7) keV


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Z sources display the horizontal branch (HB), the normal branch (NB) and the flaring branch (FB); atoll sources follow a curved branch, with a “banana” at the right and one of more “islands” at the left.

The motion of an LMXB through the CD patterns is

one-dimensional: a source always moves smoothly following the pattern rather jumps through the diagram.


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It is believed that the position and sequence of the pattern is reflected by the mass accretion rate and its change (?).

The optical and UV emission suggest that for atoll sources, mass

accretion rate increases from the island to the left of the banana branch and then from left to right along the banana, and mass accretion rate increases in the sense HBNBFB for Z sources (?).


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(1) High luminosity/Z-source systems The continuum spectra can be represented by a two component

model: isothermal blackbody from a boundary layer, and the sum of blackbodies from a multi-color disk (Eastern model), or a Comptonized disk spectrum (Western model).


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(2) Low luminosity/atoll systems The spectra are harder than in Z sources, and are modeled as simple power-law spectra with energy index ~1 and an exponential high-energy cutoff with a temperature of 5-20 keV. (3) Line emission. Iron K line emission at 6.7 keV was observed in LMXBs, most likely arising in a photo-ionized ADC. (4) Atomic spectral lines

Cottam et al. 2002


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(4) Optical emission Typical Lx/Lopt ratio ~100-1000, except for the ADC sources,

which have ratio of Lx/Lopt ~20 (because the central X-ray source is hidden).

Most of the optical emission of LMXBs originates in an accretion

disk as in CVs. Contrary to CVs, whose disks radiate internally generated energy, the main source of optical emission of LMXBs is the absorption of X-rays by the accretion disk and the subsequent re-radiation of this energy as low energy photons (reprocessing of X-rays).


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X-ray bursts (1) Type I bursts Burst profiles Rising times <1 s to 10 s, decaying times ~ 10 s to minutes. X-ray spectra The time-dependent spectra are well described by a blackbody spectrum with an approximately constant radius and a decreasing temperature during burst decay.

Average spectra for 1724-307


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The blackbody temperature can be used to determine the apparent blackbody radius of the burst-emitting region through Rbb = d (Fbol/T4)1/2 In some bursts the Eddington critical luminosity is exceeded and atmospheric layers are lifted off the star’s surface, leading to a photospheric expansion, followed by a gradual recontraction while the luminosity remains close to the Eddington limit.


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Burst intervals ~ hours-days energy release ~1036-1038 ergs-1

The longer the waiting time, the higher the burst fluence.

GS 1826-238


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Interpretation Thermonuclear explosions of H and/or He on NS surface. Type I bursts are a distinct feature of neutron stars (Narayan astro-ph/0310692).


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At the lowest mass accretion rates, / ≲ 0.01, hydrogen burning is unstable and in turn triggers unstable helium burning. At intermediate mass accretion rates, 0.01 ≲ ≲ ~0.1, hydrogen burns stably into helium between bursts, forming a helium layer at the base of the accreted material. The temperature of the fuel layer rises until the point of helium ignition is reached. At high accretion rates, 0.1 ≲ ≲ ~0.9, helium ignites unstably in a hydrogen rich environment because steady burning of hydrogen does not proceed fast enough to convert all of the hydrogen into helium. At even higher mass accretion rates, helium burning is also stable and thermonuclear flashes are not expected and rarely observed (Ozel arXiv/1210.0916)


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Burst energetics The total amount of energy emitted is ~1039-40ergs. The observed ratio of the luminosity in persistent emission between bursts and in X-ray bursts range from ~10 to 103. Theoretically, = Lp/Lb = [G (dM/dt)c2]/[N (dM/dt)c2] ~ 20-100 The lower and upper values correspond to hydrogen and helium burning respectively.


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Mass-radius relation of neutron stars For a uniform spherical emitter of radius R, at a distance d, we have Lb = 4R

2T4 = 4d2Fb

Since the energy of each photon and also the rate at which photons arrive undergo gravitational redshift, R

= R(1+z) = R(1-2GM/(Rc2)]-1/2 Thus the measurement leads to information about the radius R and the mass M of the neutron star, and thereby the EOS of neutron star matter.


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(2) Type II X-ray bursts Type II X-ray bursts have been observed in the Rapid Burster (MXB 1730-335) and possibly in GRO J1744-28. For the RB both type I and II bursts have been observed. Burst duration from ~2 s to ~ 11 min, interval from ~1 s to ~1 hr. The fluence in a type II burst is approximately proportional to the interval to the next type II burst.


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Spectra The burst peak luminosities range from ~5×1037 to ~4×1038 ergs-1. The spectra can be approximated by that of a blackbody with constant temperature ~1.8 keV (no spectral softening during burst decay). Interpretation Accretion disk instability (?)


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Periodic and quasi-periodic oscillations (QPOs) QPOs are intensity fluctuations with a preferred frequency. Approximately symmetric peaks in the power spectrum whose ratio of full width at half maximum (FWHM) and centroid frequency does not exceed 0.5 (broader features are termed noise).


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There are several kinds of QPOs in NS LMXBs.


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(1) Horizontal-branch QPOs (HBOs): QPO frequency ranges from 5 to 60 Hz. There is a positive correlation between the HBO frequency and the X-ray intensity. (2) Normal- and flaring-branch QPOs (NBOs and FBOs). NBOs frequencies ~ 4.5 and 7 Hz. FBOs occur on a small part (~10% of the total extent) of the FB nearest the NB. Their frequencies increase from ~ 6 Hz near the NB-FB junction to ~20 Hz up the FB.


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(3) Millisecond oscillations (i) Millisecond pulsations


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(ii) Burst oscillations In the initial phase of type I bursts when the burning front is spreading, the energy generation is inherently very anisotropic (also due to magnetic fields and patchy burning), leading to periodic or quasi-periodic phenomena.


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It is widely accepted that the burst oscillations arise due to a hot spot or spots in an atmospheric layer of the neutron star rotating slightly slower than the star itself because it expanded by 5-50 m in the X-ray burst but conserved its angular momentum. So the oscillation frequency (or its sub-harmonics) represents the spin frequency of the neutron star.


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(iii) Kilo-Hertz QPOs More than 20 sources have shown kHz QPOs. There are two simultaneous QPO peaks (“twin peaks”) in the 300-1300 Hz range.

Sco X-1 4U1608-52


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The frequency of both peaks usually increases with X-ray flux, with the peak separation ν remaining roughly close to (but not equal to) νs or νs/2.


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On timescales of hours to a day, kHz OPO frequency usually increases with X-ray flux. But the same positive relation does not apply to another source or the same source with longer timescale.


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Saturation of kHz QPO frequency (?)


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There are good correlations between kHz QPOs and low-frequency phenomena in WDs, NSs and BHCs, implying a common mechanism for the origin of the QPOs in these systems.


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Theoretical models (a) The beat-frequency model (Miller et al. 1998, ApJ, 508, 791) The frequency ν2 of the upper peak is the Keplerian orbital frequency νorb of accreting matter at some preferred radius (e.g. the magnetospheric radius or the sonic radius) in the accretion disk. The lower peak frequency ν1 is the beat frequency between ν2 and the spin frequency νs, so ν1 = νorb – νs, but this relation is inconsistent with observations (ν changes with ν1 or ν2).


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The clump with its spiral flow, the emission from the flow’s footpoint (dashed lines), and the clump’s interaction with the

pulsar beam (lighter shading) in the Miller et al (1998a) model.


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(b) The relativistic precession disk model (Stella & Vietri 1998, ApJ, 492, L59; 503, 350) Inclined eccentric free-particle orbits around a spinning neutron star show both nodal precession (a wobble of the orbital plane) due to relativistic frame dragging and relativistic periastron precession. The frequency ν2 of the upper peak is the orbital frequency νorb of accreting matter at some radius. The lower peak frequency ν1 is the periastron precession of the orbit.


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(c) The two-oscillator model (Titarchuk & Osherovich, 1999, ApJ, 518, L95; Osherovich & Titarchuk 1999, ApJ, 523, L73) The lower peak frequency ν1 is the Keplerian frequency at the outer edge of a viscous transition layer between the disk and neutron star surface. Blob being thrown out of this layer into the magnetosphere oscillate both radially and perpendicular to the disk, producing two harmonics of another low-frequency QPO as well as the upper kHz peaks:


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(d) The MHD oscillation model (Zhang 2004; Li & Zhang 2005; Shi & Li 2009) The upper and lower kHz QPO frequencies are identified to be the rotational frequencies and the MHD (e.g. standing kink modes of) loop oscillations at the inner edge of the accretion disk, respectively.


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See Lin et al. (2011, ApJ, 726, 74) for a comparison of fitted frequency relations for 4U 1636-53 and Sco X-1 with all currently proposed QPO models.


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4. Black Hole binaries


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Casares (2006)


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(Remillard & McClintock 2006 ARA&A)


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Soft X-ray transients A large fraction of the BHCs are X-ray transients, called soft X-ray transients or X-ray novae. The X-ray flux increases by more than two orders of magnitude within several days. The flux declines on time scales of several tens of days to more than one hundred days.


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Many (possibly all) transients are recurrent. The intervals between outbursts vary from less than one month to tens of years or more,

the duty cycle ≤10-2. There are similarities between the outbursts of LMXBs and those of dwarf novae (DNe). Thermal disk instability model has been proposed for transient outbursts.


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BH diagnostics (1) Lack of pulses and Type I X-ray bursts. (2) Spectra Ultrasoft spectra NS LMXBs can have very soft spectra (e.g. Cir X-1) though such soft spectra in NS systems are much less common. High energy power law tail above 20 keV. High-soft and low-hard states (transition around ~1037 ergs-1).


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(3) For a given orbital period, quiescent BH binaries are ~100 times dimmer than quiescent NS binaries

(from astro-ph/0903.0322)


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State transitions

Spectra of XTE J1550-564

(from astro-ph/0508494) See, however, Miller et al. 2006 (ApJ, 652, L113; 653, 525) for possible

presence of a cool disk component in low/hard state.


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BH transient outburst sources tend to follow a harder trajectory from low to high luminosity, and make the main hard-soft (LSHS) transition at much higher luminosity than the soft-hard (HSLS) one, behavior that is often called hysteresis.


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Variability Large amplitude (rms variation ~30%) flickering, QPOs. Low-frequency QPOs (LFQPOs; roughly 0.1–30 Hz) have been detected on one or more occasions for 14 BHBs They are seen in the Steep Power-law (SPL) state, and in some hard states, particularly when the X-ray luminosity is high. This behavior clearly ties LFQPOs to the non-thermal component of the X-ray spectrum.


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High-frequency QPOs (HFQPOs; 40–450 Hz) have been detected in seven sources. They do not shift freely in frequency in response to sizable luminosity changes. All of the strong detections of HFQPOs above 100 Hz occur in the SPL state.


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Name BH Mass (Msun) HF QPO frequencies (Hz) GRO J1655-40 ~6 300, 450 XTE J1550-564 ~10 184, 276, 188, 249~276 GRS 1915+105 ~14 41, 67, 113, 165, 328 H 1743-322 ~ 160, 240, 166, 242 XTE J1859+226 ~9 150~200 4U 1630-472 ~ 184, 100~300 XTE J1650-500 ~ 110~270


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Measuring Black Hole Spin (1) Continuum Fitting Continuum radiation Rin RISCO (depending on M and a) a

McClintock et al. 2011,

arXiv:1101.0811

Zhang et al. 1997, ApJ, 482, L155


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(2) Fe K line profile

Reis et al. 2008, 2009

GX 339−4 0.935 ± 0.01

SWIFT J1753.5−0127 0.76+0.11−0152

GRO J1655−40 >0.9


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(3) High Frequency Quasi-Periodic Oscillations


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5. Luminous X-ray sources in galaxies X-ray source population in galaxies It was a surprise to find with Einstein observations that many “normal” spiral galaxies also had central X-ray sources.


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The subsequent X-ray observatories ROSAT and ASCA expanded our knowledge of the X-ray properties of galaxies, but did not produce the revolutionary leap originated by the first Einstein observations. Only with Chandra’s sub-arcsecond angular resolution, populations of individual X-ray sources can be separated from the diffuse emission of hot interstellar gases, both spatially and spectrally at the distance of the Virgo Cluster and beyond.


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NGC 1313 Observed with Einstein, ROSAT, Chandra


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Two typical observations of galaxies with Chandra: the spiral M83 and the elliptical NGC4697


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X-ray luminosity Functions (XLFs) XLFs have been used to characterize different XRB populations in the Milky Way, but these studies have always suffered the distance uncertainty. External galaxies, instead, provide clean source samples. Moreover, the detection of X-ray source populations in a wide range of different galaxies allows to explore global population differences connected with the age and/or metallicity.


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LMXBs in Early-type galaxies LMXBs can account for a large fraction of the X-ray emission of some early type galaxies.


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Two breaks have been reported in the XLFs of E and S0 galaxies: the first is a break at ~2-5 x 1038 ergs-1, near the Eddington limit of an accreting neutron star, which may be related to the transition in the XLF between neutron star and black hole binaries.


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Association of LMXBs with globular clusters (GCs)? Sarazin et al. (2003) point out that the fraction of LMXBs associated with GCs increases from spiral bulges (MW, M31 ~10-20%), to S0s ~20%, E ~50%, and cD~70%, suggesting that most of LMXBs may originate from GCs.


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Other authors conclude that the relationship between the fraction of LMXBs found in GCs and the GC specific frequency is consistent with the simple relationship expected if field LMXB originate in the field while GC LMXB originate in GCs (Juett 2005; Irwin 2005).


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LMXBs preferentially are found in luminous GCs, and in red, younger and/or metal rich, clusters (V-I >1.1), rather than in blue, older and/or metal poor, ones (Kundu, Maccarone & Zepf 2002, 2003). (Age or metallicity effect?)


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Young XRB populations Luminous HMXBs are expected to dominate the emission of star forming galaxies. These sources, resulting from the evolution of a massive binary system where the more massive star has undergone a supernova event, are short-lived (~ 106-7 yr), and constitute a marker of recent star formation: their number is likely to be related to the galaxy star formation rate (SFR).


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The HMXB XLF is overall flatter than that of LMXBs, with a cumulative power-law slope of −0.6 to −0.4; in other words, young HMXBs populations contain on average a larger fraction of very luminous sources than the old LMXB populations.


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Grimm, Gilfanov & Sunyaev (2003) show that there is a universal XLF of star-forming populations with a simple power-law with cumulative slope −0.6.


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Both the number and total luminosity of HMXBs in a galaxy are directly related to the SFR and can be used as an independent SFR indicator.

for

for


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The differences of the XLF in different regions of a galaxy, and in galaxies with different SFR


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These differences may be related to the aging of the X-ray source population, which will be gradually depleted of luminous young (and short-lived) sources associated with more massive, faster-evolving, donor stars, and also to metallicity effects.


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Ultraluminous X-ray sources (ULXs) For a recent review, see Feng & Soria 2012, NewAR, 55, 166 (arXiv:1109.1610) ULXs are the most luminous point-like extra-nuclear X-ray sources found in nearby galaxies. Observed (isotropic) X-ray luminosities in excess of 1039 ergs-1, the Eddington luminosity for a 10 M⊙ BH.

138TpE ergs )/(103.1/4

MMGMmL


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Location in galaxies ULXs are associated with both star-forming regions in spiral and irregular galaxies, and the old stellar population in elliptical galaxies. In spirals, ULXs are often near, but distinct from the dynamical centers of the galaxies. In ellipticals, ULXs are almost exclusively in the halos of the galaxies. The brightest ULXs are in the brightest FIR galaxies.


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The mean X-ray luminosity of ULXs in elliptical galaxies is less than in spiral galaxies.

Swartz et al. 2003


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Association of ULXs with active star-forming stellar populations


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Spectra and spectral variability The ASCA X-ray spectra generally consist of a power law (~1.2) and a disk-blackbody component with Tin=1.1-1.8 keV. They also show

spectral transitions between soft/high and hard/low states, similar to the state changes observed in Galactic black hole binaries.


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XMM-Newton spectra of two ULXs in NGC 1313 (X-1 and X-2) led to highly significant detections of soft accretion disk components, with temperatures of kT~150 eV, consistent with accretion disks of IMBHs (Miller et al. 2003, 2004).

1/ 2 1/ 4 1/ 4inKT m m


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Intermediate Mass Black Holes or XRBs? (1) Stellar-mass BHs with real super-Eddington X-ray emission (Begelman 2002) or anisotropic emission (King et al. 2001) or relativistically beamed emission (Kordng, Falcke & Markoff 2002) LX=bLsph=1040bL40,

b=/4, or b= [(1-cos)]-p (2) Intermediate mass (102-105M⊙) BHs with sub-Eddington

luminosities (IMBHs, Colbert & Mushotzky 1999; Miller & Hamilton 2002).


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HMXBs and IMXBs as ULXs

Podsiadlowski et al. 2003 Rappaport et al. 2004


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LMXBs as luminous X-ray sources (1) Transient (long orbital period) LMXBs Trudolyubov & Priedhosky (2004) report only one recurrent transient in their study of GC sources in M31, although 80% of these sources show some variability; however, they also find six persistent sources in the 1038 ergs-1 luminosity range. (2) Ultra-compact binaries (3) Persistent LMXBs with a CB disk


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Difficulties for the Beaming Model Emission nebulae of a few hundred pc diameter are found to be present at or around several ULXs


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The problem(s) with IMBHs (1) The XLF has an unbroken power-law form for 5 decades up to a luminosity of~21040ergs-1, this break occurs at ~10% of the Eddington luminosity for the~1000M⊙ black holes. (2) Association of ULXs with star formation.


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Formation Scenarios for IMBHs (1) Merging of stars in a young dense stellar cluster followed by direct collapse into an IMBH (Portegies Zwart et al. 1999). (2) Merging of binaries that have a black hole with initial mass of ~50 M⊙ in a globular cluster (Miller

& Hamilton 2002). (3) Evolution of primordial population III stars (Madau & Rees 2001).


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Most ULXs could be explained by stellar-mass black holes that are either super-Eddington, or subject to some sort of beaming. However, it is still difficult to reconcile the most extreme ULXs - those above 1040ergs-1 – with simple stellar-mass black hole systems.


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References 1. Bhattacharya, D. and van den Heuvel, E. P. J. 1991, Phys. Rep.,

203, 1 2. Lewin, W. H. G., van Paradijs, J., and van den Heuvel, E. P. J. 1995,

X-ray binaries (chapters 1-4, 6) 3. Nagase, F. 1989, PASJ, 41, 1 4. Bildsten, L. et al. 1997, ApJS, 113, 367 5. Van der Klis, M. 2000, ARA&A, 38, 717 (astro-ph/0001167) 6. Remillard, R. D. and McClintock, J. E. 2006, ARA&A,

(astro-ph/0606352) 7. Fabbiano, G. 2006, astro-ph/0511481

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Chapter 3 Observational Characteristics of X-ray Binaries