University of Groningen Disc reflection in low-mass X-ray binaries Wang, Yanan

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Disc reflection in low-mass X-ray binaries

Wang, Yanan

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Wang, Y. (2018). Disc reflection in low-mass X-ray binaries. Rijksuniversiteit Groningen.

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Disc reflection in low-mass X-ray binaries

Proefschrift

ter verkrijging van de graad van doctor aan de Rijksuniversiteit Groningen

op gezag van de

rector magnificus Prof. dr. E. Sterken en volgens besluit van het College voor Promoties.

De openbare verdediging zal plaatsvinden op maandag 3 december 2018 om 9:00 uur

door

Yanan Wang

geboren op 08 mei 1988 te Liaoning, China

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Prof. M. Méndez Copromotor Dr. D. Altamirano Beoordelingscommissie Prof. S. Zhang Prof. C. Done Prof. R. Morganti

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iii

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Cover: Artist’s impression of Cygnus X-1 (Credit NASA, ESA, Martin Kornmesser) Extra design by: C. Zheng.

Printed by GVO Drukkers & Vormgevers on recycled paper. ISBN: 978-94-034-1255-9 (printed version)

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CONTENTS v

1

Thesis introduction

1.1 X-ray binaries

As astronomical objects can attract the surrounding matter through their gravitational force, this process, known as accretion, thus plays a fundamental role at all scales en-countered in the universe. X-ray binaries, as a class of binary stars emitting luminous X-rays, are excellent laboratories to study this process.

X-ray binary systems consist of a compact object, either a black hole (BH) or a neutron star (NS), and a companion star, both of which orbit a common centre of mass of the system. These systems can be subdivided into high-mass X-ray binaries (HMXBs) and low-mass X-ray binaries (LMXBs), based on the mass of the compan-ion star. In the former systems, matter is mainly accreted via a stellar wind, whereas in the latter ones matter is accreted onto a compact object via Roche-lobe overflow in an accretion disc (see Fig. 1.1). In the following, I focus on introducing the properties of LMXBs.

In LMXBs, the companion star is usually a late-type main sequence star or a white dwarf, with a mass typically lower than 1 M. Such systems can be formed in two different ways: i) two stars are directly gravitationally bound from birth, both of them surviving a supernova explosion (see Tauris & van den Heuvel 2006); ii) after a massive star collapses into a compact star in an environment with a high density of stars, a second star may be captured by it (see Bhattacharya & van den Heuvel 1991; Verbunt & Lewin 2006). The latter process, however, mainly happens in globular clusters. Most of the binaries in the Milky Way have been born in a bound state. In an LMXB system, the matter from the outer layers of the normal star is stripped off and is then accreted onto the compact object. With a large amount of angular mo-mentum, the transferred matter cannot be radially accreted onto the compact object,

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Figure 1.1 – Comparison of HXMB and LMXB. Each panel presents an artist’s view (top) and a schematic view (bottom) of the systems to illustrate the different accretion modes, respectively via stellar wind and via Roche lobe overflow through an accretion disc. Figure from Egron (2013).

but forms a rotating disc. The X-rays stem from the conversion of the gravitational energy of the in-falling material into radiation through viscous processes or shocks occurring in the accretion disc. As the accreted matter moves inwards, the disc tem-perature gradually increases, which can be up to several million degrees in the region closest to the compact object. However, X-ray energy spectra of these systems can extend up to 100 keV. This indicates that the disc is not the only emitting region. The region that produces this high-energy radiation is called corona, which is com-posed of hot electrons (T > 109_{K) and is responsible for producing the high-energy}

photons by inverse Compton scattering.

1.2 Components of binary systems

As was just mentioned in the previous section, the accretion disc plays a key role in the X-ray emission of LMXBs. The broadband energy spectrum also indicates that there is a significant contribution from the corona in X-rays. The study of the accre-tion disc and the corona, as well as their interacaccre-tion, helps us to better understand the physics of these systems.

1.2.1 The accretion disc

The accretion disc is composed of stellar plasma whose properties are determined by the ions and electrons that compose it. Depending on the temperature, the matter can be neutral, partially or fully ionized. The disc structure and geometry seem to be affected by the mass accretion rate and viscosity.

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1.3: X-ray spectra 3

The standard accretion disc model was introduced by Pringle & Rees (1972) and Shakura & Sunyaev (1973), who described an optically thick but geometrically thin accretion disc, which is dominated by gas pressure and is close to a blackbody. This standard disc model is applied to the radiatively efficient case, whereas for some sources with a low accretion rate, this model is not capable of reproducing the corre-sponding spectral energy distribution (Esin et al. 1996).

The Advection Dominated Accretion Flow (ADAF) model was initially introduced by Ichimaru (1977); the model assumes that once the disc is close to the central object, its inner part becomes optically thin and is replaced by a hot flow (Narayan & Yi 1994, 1995a,b). In this scenario, only a small amount of the energy is radiated away (Narayan 1996). Meanwhile, most of the energy is advected radially with the flow onto the compact object. Therefore, the radiation efficiency of ADAF is much lower than that of the standard disc (see Narayan, Mahadevan & Quataert 1998; Kato, Fukue & Mineshige 1998). ADAF is mainly applied to black-hole systems, as part of the accretion energy is advected into the black hole. In neutron-star systems the accretion is efficient, since the energy can eventually be released by the surface of the neutron star, even in the propeller state due to the near spherical ADAF (e.g. Zhang, Yu & Zhang 1998; Menou et al. 1999).

1.2.2 The corona

The corona consists of hot electrons and is the region where the high-energy X-ray photons, up to hundreds of keV, are produced. Both the origin and the geometry of the corona are still unclear. Shakura & Sunyaev (1973) suggested that the corona may be formed by evaporation of hot material from the disc. Alternatively, Liang & Price (1977) proposed that an accretion disc may pump energy into an outer tenu-ous layer of its atmosphere to form a high-temperature corona analogtenu-ous to the solar corona. As for its geometry, there are several main models: the ‘slab’ or ‘sand-wich’, the ‘sphere’ (ADAF-like) and the ‘patchy’ or ‘pill box’ corona. Besides those, the ‘lamppost’, an on-axis isotropic point, geometry (e.g. Martocchia & Matt 1996; Miniutti & Fabian 2004) has become a popular model for accreting black holes.

1.3 X-ray spectra

As explained previously, the accretion disc and the corona play major roles in the X-ray emission of LMXBs. By modelling X-X-ray spectra, one can study the interactions between matter and photons, the geometry of the disc and corona, and further explore the physics of X-ray binaries.

X-ray spectra of accreting binary systems generally exhibit two main components: a soft component associated with an accretion disc, and a hard component resulting from inverse Compton scattering of soft photons by high-energy electrons in a corona

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TOTAL COMPTONIZED REFLECTED DISK REFLECTED COMPTONIZED DISK

Figure 1.2 –The three main components of the X-ray emission from an accretion black hole (top) and a plausible geometry of the accretion flow in the hard spectral state (bottom), where the black point, the blue slabs and the magenta points represent the black hole, the accretion disc and the corona, respectively. Figures from Gilfanov (2010).

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1.3: X-ray spectra 5

(see Sunyaev & Truemper 1979; Sunyaev & Titarchuk 1980). These two components can be fitted by a multi-temperature disc blackbody (e.g. Mitsuda et al. 1984) and a power-law with a high-energy cutoff, respectively. In a neutron-star LMXB, a third soft thermal component is required by the data, which is slightly hotter than the ac-cretion disc and originates from the neutron-star surface or the boundary layer (e.g. Gilfanov & Revnivtsev 2005). Additionally, a reflection component is often detected in such spectra and is interpreted as the result of the high-energy coronal photons irradiating an optically thick accretion disc (Fabian et al. 1989). These typical com-ponents in a black-hole LMXB system are shown in Fig. 1.2.

The reflection spectrum is the superposition of Compton scattering, fluorescence, Auger effect and recombination emission produced in the accretion disc (e.g. Fabian et al. 1989; Ross & Fabian 1993). Fig. 1.3 shows the results of a Monte-Carlo sim-ulation modelling these processes when a power-law X-ray continuum with photon index = 2 is incident on a gaseous slab. The most notable signature of the reflection spectrum is the iron emission line at 6.4–7 keV plus a Compton hump at 10–30 keV. Fluorescent and recombination emission lines are generated by irradiation of the cold matter nearby the compact object. X-ray photons are photoelectrically absorbed by particles and hence their K-shell electron is removed. The K-hole is then filled by an electron from outer shells and subsequently, a K-shell photon (fluorescence) or a second electron (Auger effect) is emitted (Bambynek et al. 1972). The probability of emitting a fluorescent photon over an Auger electron is determined by the flu-orescent yield, which is proportional to the nuclear charge Z to the fourth power (i.e. ∝ Z4_{). The combination of the high fluorescent yield and the relatively high}

abundance makes the iron emission line, at 6.4 keV, the most prominent one in the reflection spectrum. Highly ionized iron can be detected at 6.67–6.70 keV and 6.95– 6.97 keV, associated with Fe XXV (He-like) and Fe XXVI (H-like), respectively. The He-like lines are formed partly by fluorescence and partly by recombination, whereas H-like lines are only formed by recombination (Hatchett, Buff & McCray 1976). The above process happens mainly when the incident photons have energies lower than 15 keV. Once the energy is higher than that, photons are substantially Compton scattered until they are emitted away or are photo-absorbed. Both processes result in a Compton hump at 10–30 keV in the spectrum.

Another consequence of photoelectric absorption is the absorption edge. When the energy of an incident photon is higher than the energy of an electronic transition or to the ionization potential of an atom, the probability of the atom absorbing that photon dramatically decreases with increasing photon energy. This resulting jump in cross section is cabled an absorption edge, e.g. the iron K edge at 7.1 keV in Fig. 1.3. The iron line in the reflection spectrum is intrinsically narrow, but may appear broad-ened and skewed towards lower energies. Since the iron line arises from the

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in-1

Figure 1.3–X-ray reflection from an illuminated slab. The dashed line shows the incident continuum and the solid line shows the reflected spectrum (integrated over all angles). The main reflected features are the iron fluorescence line at 6.4 keV, the iron absorption edge at 7.1 keV and the Compton hump peaking at 30 keV. Monte Carlo simulation from Reynolds (1996).

nermost regions of the accretion disc, its profile is modified by the combination of Doppler shifts, Doppler broadening and gravitational redshifts (Fabian et al. 2000), as illustrated in Fig. 1.4. At each radius of the disc, a symmetric double-peak line profile is produced by the emission from the matter approaching to (blue-shifted) and receding from (red-shifted) the observer. As matter proceeds closer to the central object, special relativistic beaming enhances the blue peak of the line. Moreover, gravitational redshift shifts the line to lower energy. Summing all the effects, the shape of the line appears to be broad and asymmetric.

Additionally, the line profile also depends on the inclination of the system with re-spect to the line of sight, the inner and outer radius of the disc, the emissivity index (Fabian et al. 1989) and the disc ionization state (e.g. Fabian et al. 2000).

Besides spectral analysis, timing analysis reveals different aspects of the accretion processes in X-ray binaries. The study of the rapid variability constitutes a powerful tool to probe the properties of the inner most regions of the accretion disc. The main

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1.4: X-ray states 7 0.5 1 1.5 Line profile Gravitational redshift General relativity

Transverse Doppler shift Beaming

Special relativity Newtonian

Figure 1.4 –Left panel: a schematic representation of the accretion disc from the top, with matter approaching (blue) and receding (red). Right panel: the first three panels showing the Doppler shift and relativistic effects on a narrow emission line profile at two different radii in the disc. The two line profiles indicate the emission from the two disc radii as shown in the left panel: the broadest profile corresponds to the ring at the smallest radius. The bottom panel shows a broad and skewed line profile after integrating all disc radii, resulting from the sum of the above effects. Figures taken from Fabian et al. (2000).

method to study the fast variability is the Fourier power density spectrum, which describes how the power of a time series is distributed with frequency.

As an energy spectrum is time averaged, one always needs to consider the influence of the source variability on its spectrum. Studying timing properties of a system also helps us to understand the source state. Combined results of spectral and timing analysis bring a more complete view of what happens in the source.

1.4 X-ray states

The existence of the X-ray states in LMXBs was first demonstrated by the work of Tananbaum et al. (1972) in which they observed that the soft X-ray flux (2–6 keV) of Cygnus X–1 decreased whereas its hard X-ray flux (10–20 keV) increased. Since then, similar X-ray transitions have been discovered in many sources which are called

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X-1

ray luminosity can be Lx∼ 1036−39erg s−1, and after several days to months or even

years, they become quiescent.

For BH X-ray binaries, there are two main states, the soft state and the hard state, depending on the source timing and spectral properties (see Remillard & McClintock 2006, for a detailed review). During the soft state, an optically thick and geometri-cally thin accretion disc is observed; the disc is normally thought to extend down very close to the compact object, i.e. the inner stable circular orbit (ISCO) in BH systems or the NS surface. This component is commonly described as thermal emission with a temperature of ∼ 1 keV and can be fitted by a blackbody or a multi-temperature blackbody. A non-thermal component is detected at energies above ∼ 10 keV, and is fitted with a power law by a photon indexΓ ∼ 2.1−3, whose contribution is limited to < 25% of the flux in the 2–20 keV energy range (also see Remillard & McClintock 2006). During the hard state, the spectrum is dominated by a power-law component with Γ ∼ 1.5 − 1.9 and a high-energy cutoff, which contributes ≥ 80% of the 2– 20 keV flux. A weak thermal component is detected at low energies, likely from a cool disc that appears to be truncated at relatively large distances from the compact object (Done, Gierli´nski & Kubota 2007).

For NS X-ray binaries the scenario is similar, except that one additional component from the NS surface or the boundary layer is required (e.g. Syunyaev & Shakura 1986; Inogamov & Sunyaev 1999). This boundary layer is formed when the rapidly rotating gas in the disc reaches the slowly/non rotating accreting NS. During this process, part of the gravitational potential energy has been radiated by the gas in the disc and the rest of the energy is released in the boundary layer. Since this energy comes from a small region nearby the NS, the boundary layer is consequently hotter than the disc, which results in harder radiation.

1.5 X-ray facilities

The study of astronomical objects at the energy ranges of X-rays and γ-rays began

in the early 1960s. Before then, the Sun was believed to be an intense source in these wavebands, but no other objects had been observed in X- orγ-rays. The Earth’s

atmosphere is opaque to electromagnetic radiations at certain wavelengths, e.g.

γ-rays, X-γ-rays, most UV and Infrared light (see Fig. 1.5). This means that, in order to observe extra-terrestrial sources of such emission, these sources must be observed from above the atmosphere.

With the launch of rocket flights, the first non-solar X-ray source was successfully de-tected in 1962 by Riccardo Giacconi, Herb Gursky, Frank Paolini, and Bruno Rossi. Later, this bright X-ray source was designated as Scorpius X-1. Riccardo Giacconi subsequently won the Noble Prize for Physics in 2002 for his contribution to X-ray astronomy. The first earth-orbiting X-ray telescope UHURU, which was also known

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1.5: X-ray facilities 9

Figure 1.5– Electromagnetic radiation through the Earth’s Atmosphere, fromγ-rays on the left (the most powerful, shortest wavelength) over to the Radio on the right (the longest wavelength). The wavelength and frequency of light determines how far it can penetrate into the Earth’s Atmosphere. Credits: Harris Geospatial Solutions.

as the Small Astronomical Satellite-1, based on Giacconi’s design, was launched on 12 December 1970, operated over two years and ended in March 1973. Ever since, new missions have been launched, and we entered into a new era of discovery in high energy astrophysics. In this section I briefly describe the main instruments used in this thesis.

1.5.1 NUSTAR

The Nuclear Spectroscopic Telescope Array (NUSTAR) is a National Aeronautics

and Space Administration (NASA) small explorer mission that carries the first focus-ing hard X-ray optics into space (Harrison et al. 2013); it was successfully launched on 13th June 2012. With a mass of 350 kg for a focal length of 10.14 meters, NU

S-TAR was launched into a near-equatorial, low-Earth orbit, to minimize exposure to the South Atlantic Anomaly (SAA) where there is a high concentration of trapped particles. NUSTAR has a pair of co-aligned focusing telescopes (Fig. 1.6), which

are designed to provide a broad-band coverage in the hard X-rays between 3–79 keV. Compared to the previous hard X-ray instruments (> 10 keV), NUSTAR results in

an order of magnitude improvement in angular resolution and has two orders of mag-nitude improvement in sensitivity. Fig. 1.7 shows a comparison of the NUSTAR

effective collecting area with some other focusing X-ray detectors currently on orbit. 1.5.2 XMM-NEWTON

XMM-NEWTON, the X-ray Multi-Mirror Mission, is an X-ray space observatory

launched by the European Space Agency (ESA) on 10th December 1999. The main objective of this observatory is to investigate galactic and extragalactic X-ray sources, perform narrow- and broad-range spectroscopy, and obtain the first simultaneous imaging of objects in both X-ray and optical (visible and ultraviolet) wavelengths. With a mass of 3,800 kg for a length of 10.8 meters, XMM-NEWTONwas launched into space aboard an Ariane 504 rocket and placed into a highly elliptical orbit with

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Opcs modules Deployed mast Focal plane bench Instrument star tracker

Mast canister Metrology detector (1 of 2) Metrology

lasers

Focal plane detector module (1 of 2)

Figure 1.6–Diagram of the NUSTAR observatory in the stowed (bottom) and deployed (top)

config-urations (adapted from Harrison et al. 2013).

Figure 1.7–Effective collecting area of NUSTAR compared to selected operating focusing telescopes

(CHANDRA, XMM-NEWTON, and SUZAKU). NUSTAR provides good overlap with these soft X-ray

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1.5: X-ray facilities 11

a perigee of 6,000 km and an apogee of 112,000 km. This mission carries three high throughput X-ray telescopes with an unprecedented effective area and an optical monitor (see Fig 1.8). The instruments on board XMM-NEWTONare:

• The European Photon Imaging Cameras

The three European Photon Imaging Cameras (EPIC) are the primary instru-ments aboard XMM-NEWTON. The system is composed of two MOS-CCD

cameras (Turner et al. 2001) and a single pn-CCD camera (Strüder et al. 2001), with a total field of view of 30arcmin and an energy sensitivity range between 0.15–12 keV, with an energy resolution of E/∆E = 20 − 50 and an angular resolution of 6arcsec.

• The Reflection Grating Spectrometer

The two Reflection Grating Spectrometers (RGS) are a secondary system on the spacecraft and are composed of two Focal Plane Cameras and two Reflec-tion grating Assemblies (den Herder et al. 2001). This system is used to collect high resolution X-ray spectral data and can determine the elements present in the target, as well as the temperature, abundance and other characteristics of those elements. The RGS system operates in the energy range of 0.33–2.5 keV, which allows the detection of carbon, nitrogen, oxygen, neon, magnesium, sil-icon and iron.

• The Optical Monitor

The Optical Monitor (OM) is a 30 cm Ritchey-Chretien optical/UV telescope designed to provide simultaneous observations alongside the spacecraft’s X-ray instruments (Mason et al. 2001). It consists of three optical and three UV filters over the wavelength range of 180–600 nm.

1.5.3 SWIFT

The Neil Gehrels SWIFT Observatory∗, previously called the SWIFT Gamma-Ray

Burst Mission, is a NASA space telescope launched on 20th November 2004. With a mass of 1,500 kg, SWIFT was launched into a low circular Earth orbit, aboard a

Delta II rocket. This telescope is a multi-wavelength space observatory dedicated to the study ofγ-ray bursts (GRBs). Its three instruments work together to observe

GRBs and their afterglows in theγ-ray, X-ray, UV, and optical wavebands. The three

instruments are:

• Burst Alert Telescope

The Burst Alert Telescope (BAT) is an alert system to detect GRB events and

∗ _{The name ‘Swift’ is not a mission-related acronym, but rather a reference to the instrument’s rapid}

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Figure 1.8–Sketch of the XMM-Newton payload. On the lower left: the aperture doors and the mirror

modules with the two Reflection Grating Arrays. The optical monitor is obscured by the lower mirror module. On the upper right: the EPIC MOS cameras (green), the EPIC pn camera (purple), and the RGS detectors (red). The direction of the incoming X-rays is indicated by the blue arrow on the left. Figure adapted from the VILSPA XMM-Newton Science Operations Centre (Dornier Satellitensysteme GmbH).

Figure 1.9 – SWIFT’s three scientific instruments work together to obtain as much information as

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1.6: Thesis motivation and overview 13

compute their coordinates in the sky. It is capable of observing an event in the range of 15–150 keV and to locate the position of each event with an accuracy of 1 to 4arcmin within 15 s.

• X-Ray Telescope

The X-ray Telescope (XRT) is a focusing X-ray telescope with an effective area of 110 cm2 _{at 1.5 keV and an 18arcsec spatial resolution covering the range}

of 0.2–10 keV. It can take images and perform spectral analysis of the GRB afterglow. This provides a more precise location of the GRB, with a typical error circle of approximately 2arcsec radius. The XRT is also used to perform long-term monitoring of GRB afterglow light-curves for days to weeks after the event, depending on the brightness of the afterglow.

• Ultraviolet/Optical Telescope

After SWIFT has slewed towards a GRB, the Ultraviolet/Optical Telescope

(UVOT) is used to detect an optical afterglow, in the range of 170–650 nm. The UVOT is also used to provide long-term follow-ups of GRB afterglow light-curves.

Since several X-ray sources are multi-wavelength emitters, it is extremely useful to perform simultaneous observations of a source in different energy bands, in order to better understand the physics and the connections between a compact star and the matter in its vicinity.

1.6 Thesis motivation and overview

The origin of the spectral components in X-ray binaries, such as the soft/hard X-ray excess and the iron Kα line complex are still unclear. Properties of the accretion

flow, e.g. the accretion geometry, composition of the accreted material and the mass accretion rate, can be primarily studied through X-ray observations. How to detect and measure them accurately remains challenging.

Disc reflection, the interaction between the corona and/or other components (e.g. the NS surface/boundary layer in NS systems or the jet or the distant reflector in AGN) and the accretion disc, is a powerful tool to diagnose the physical and dynamical con-ditions of accreting sources. Modelling the reflection features can lead to important constraints on the disc ionization state (e.g., Ross, Fabian & Young 1999; García & Kallman 2010; García, Kallman & Mushotzky 2011; García et al. 2013), the disc in-clination angle (e.g., Fabian et al. 1989; Mushotzky et al. 1995; Crummy et al. 2006), the inner radius of the disc and the spin parameter of the compact object (e.g., Fabian et al. 1989; Laor 1991; Brenneman & Reynolds 2006; Reynolds & Fabian 2008). However, the accuracy of the measurement of some parameters derived from reflec-tion models are debated. For instance, Pandel, Kaaret & Corbel (2008) and Sanna

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et al. (2013) studied the iron emission line in the NS 4U 1636–53 with XMM-NEWTON and found a very high inclination angle of the accretion disc, consistent

with 90◦_{, even though neither eclipses or dips have ever been observed in this source.}

Recently, a very high value of the iron abundance in the accretion disc of several sys-tems has been brought up. García et al. (2015b) and Parker et al. (2016) found iron abundances of, respectively, 4.0–4.3 and 4.7±0.1 in solar units in the BHC Cyg X–1; similarly, values of 5 ± 0.1 and 6.6+0.5_−0.6times solar for GX 339–4 have been reported by Parker et al. (2015) and Walton et al. (2016) as well.

Triggered by these findings and issues, my work is mainly focused on the study of the reflection component and the inner flow in the vicinity of a compact object in LMXBs. I conducted all my research with the instruments described in the previous section. The chapters can be summarised as follows:

• Chapter 2 – The XMM-NEWTONspectra of 4U 1630–47 revisited

In this chapter, we analysed two XMM-NEWTONobservations of the

black-hole candidate 4U 1630−47 during the 2012 outburst, for which there had been a claim of Doppler-shifted emission lines that were interpreted as arising from baryonic matter in the jet. We applied an alternative model that, without the need of Doppler-shifted emission lines, fits the data well. The fit to all the 2012 XMM-NEWTONobservations of this source require a moderately broad

emission line at around 7 keV plus several absorption lines and edges.

• Chapter 3 – The reflection spectrum of 4U 1636–53

To explore the anomalously high-inclination observed in the past from fits to the XMM-NEWTONspectra of the neutron star low-mass X-ray binary

4U 1636−53, we investigated three NUSTAR observations of this source in the

soft, transitional and hard state. By applying different models to the data, we got a reasonable inclination angle of ∼ 56◦_{derived from the lamppost}

compo-nentRELXILLLP. We also discuss the results for these models and the possible

primary source of the hard X-rays in this system.

• Chapter 4 – The reflection component in the average and heartbeat

spec-tra of IGR J17091–3624

Similar to the well known black hole candidate GRS 1915+105, IGR J17091– 3624 is a black hole candidate with heartbeat variability. In this chapter, we jointly fit the NUSTAR and SWIFT spectra of IGR J17091–3624 during its

2016 outburst to study the evolution of the spectral parameters in this outburst and, in particular, during the QPO cycle of the heartbeat.

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1.6: Thesis motivation and overview 15

• Chapter 5 – Study of the X-ray properties of the neutron-star binary

4U 1728–34 from the soft to hard state

We fitted five XMM-NEWTONspectra of the neutron-star LMXB 4U 1728–34

with several popular reflection models and found that almost all of them yield a supersolar iron abundance, up to 10 times solar. We also explore the possible reasons why the supersolar iron abundance is required by the data and found that this high value is probably caused by the absence of the hard photons in the XMM-NEWTONdata. We also found the change of the inner disc does not

support the standard accretion disc model.

• Chapter 6 – Conclusions and future work

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2

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2

The XMM-Newton spectra of the 2012 outburst of the

black-hole candidate 4U 1630–47 revisited

Yanan Wang1_{and Mariano Méndez}1

MNRAS, 2016, 456, 1579

1_{Kapteyn Astronomical Institute, University of Groningen, PO BOX 800, NL-9700 AV} Gronin-gen, the Netherlands

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2 Abstract

Recent XMM-NEWTONobservations of the black-hole candidate 4U 1630−47

dur-ing the 2012 outburst revealed three relativistically Doppler-shifted emission lines that were interpreted as arising from baryonic matter in the jet of this source. Here we reanalyse those data and find an alternative model that, with less free parameters than the model with Doppler-shifted emission lines, fits the data well. In this model we allow the abundances of S and Fe in the interstellar material along the line of sight to the source to be non solar. Among other things, this significantly impacts the emission predicted by the model at around 7.1 keV, where the edge of neutral Fe appears, and renders the lines unnecessary. The fits to all the 2012 XMM-NEWTON

observations of this source require a moderately broad emission line at around 7 keV plus several absorption lines and edges due to highly ionised Fe and Ni, which reveal the presence of a highly-ionised absorber close to the source. Finally, the model also fits well the observations in which the lines were detected when we apply the most recent calibration files, whereas the model with the three Doppler-shifted emission lines does not.

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2.1: Introduction 19

2.1 Introduction

The soft X-ray transient 4U 1630−47 (Jones et al. 1976; Parmar, Angelini & White 1995) shows regular outbursts every 600-690 days (Abe et al. 2005; Tomsick et al. 2014). The source has been classified as a black hole (Parmar, Stella & White 1986) because of the similarity of its spectral and timing properties to those of systems with measured black-hole masses (e.g., Barret, McClintock & Grindlay 1996; Abe et al. 2005). 4U 1630−47 shows strong absorption by neutral material along the line of sight, with a hydrogen column density NH=5 − 12 × 1022cm−2 (e.g., Tomsick,

Lapshov & Kaaret 1998), and both IR (Augusteijn, Kuulkers & van Kerkwijk 2001) and radio emission (Hjellming et al. 1999) were detected during the 1998 outburst of this source. The optical counterpart of 4U 1630−47 has not been identified, mostly due to the high reddening and the location of the source in a crowded star field (Par-mar, Stella & White 1986).

Absorption lines due to highly ionised material have been observed in the spectrum of 4U 1630−47 (Kubota et al. 2007; Ró˙za´nska et al. 2014; Díaz Trigo et al. 2014; Miller et al. 2015). Using SUZAKUobservations carried out in 2006, Kubota et al.

(2007) studied these absorption line features in relation to the accretion-disc parame-ters, and concluded that the lines were due to a wind. Using the same SUZAKUdata,

Ró˙za´nska et al. (2014) proposed that the absorption lines could be produced effec-tively in the accretion disc atmosphere. Using XMM-NEWTONobservations, Díaz

Trigo et al. (2014) found a thermally/radiatively driven disc wind in 4U 1636−47; the wind becomes more photoionised as the luminosity of the source increases. Recently, Miller et al. (2015) analysed CHANDRAobservations of 4U 1630−47 and three other

galactic black hole candidates. For 4U 1630−47, they found that the wind consists of at least two absorption zones with velocities of −200 km s−1_{and −2000 km s}−1_,

respectively. They also found that, in some respects, these zones correspond to the broad-line region in active galactic nuclei.

Díaz Trigo et al. (2013) analysed two XMM-NEWTONand two quasi-simultaneous

observations with the Australia Telescope Compact Array (ATCA) carried out during the 2012 outburst of 4U 1630−47. Díaz Trigo et al. (2013) found three relatively narrow emission lines in the X-ray spectrum of one of these observations that they identified as arising from baryonic matter in a jet. The three lines had energies of 4.04 keV, 7.28 keV and 8.14 keV, respectively, which Díaz Trigo et al. (2013) inter-preted as the red- and blueshifted component of Fe XXVILyα and the blueshifted

component of Ni XXVIII Lyα, respectively. From the radio data, Díaz Trigo et al.

(2013) confirmed that there was an optically thin jet in 4U 1630−47 at the time of that observation. Hori et al. (2014) investigated SUZAKUand Infrared Survey

Facil-ity observations of 4U 1630−47 during the same outburst, at a time when the source was in the very high state. These observations were carried out three to five days after the XMM-NEWTON observations of Díaz Trigo et al. (2013). The SUZAKU

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Table 2.1–XMM-NEWTONObservations of 4U 1630−47 used in this chapter

ObsID Observation Time (UTC) Observation mode RAWX source RAWX back (day/month/year hh:mm) 0670671501-1 04/03/2012 11:24 - 04/03/2012 12:27 Timing [27,46] [4,10] 0670671501-2 04/03/2012 13:43 - 05/03/2012 09:23 Timing [28,45] [4,10] 0670671301 20/03/2012 19:54 - 21/03/2012 02:30 Timing [28,45] [4,10] 0670672901 25/03/2012 04:14 - 25/03/2012 21:56 Timing [28,45] [4,10] 0670673001 09/09/2012 21:14 - 10/09/2012 07:49 Timing [28,45] [4,10] 0670673101 11/09/2012 20:56 - 12/09/2012 05:38 Burst [20,51] [4,10] 0670673201 28/09/2012 07:16 - 28/09/2012 21:48 Burst [20,51] [4,10] Notes. ObsID 0670671501 contains two separate event files in timing mode that We extracted and fitted separately. We called them 0670671501-1 and 0670671501-2, respectively. RAWX source and RAWX back indicate the extraction region in the CCD for the source and the background, respectively.

X-ray spectra, however, did not show the Doppler-shifted emission lines of the jet re-ported by Díaz Trigo et al. (2013). Using CHANDRAand ATCA observations taken

eight months prior to the XMM-NEWTONobservations of Díaz Trigo et al. (2013),

Neilsen et al. (2014) reported a similar result to that of Hori et al. (2014).

When fitting the CHANDRAdata of 4U 1630−47, Neilsen et al. (2014) allowed the

abundances of Si, S and Ni in the component that they fitted to the interstellar absorp-tion to be different from solar but, unfortunately, they do not report the best fitting values of these parameters. On the other hand, using the Reflection Grating Spec-trometer on board XMM-NEWTON, Pinto et al. (2013) measured the abundances of

O, Ne, Mg, and Fe in the interstellar medium (ISM) in the direction of nine low-mass X-ray binaries, not including 4U 1630−47. Interestingly, they found that the Fe abundance in the neutral ISM in the direction of these sources ranges between less than 0.02 and 0.50 times the solar abundance. Because the putative lines reported by Díaz Trigo et al. (2013) are close to the Kα edges of (neutral) Ca I (4.04 keV),

FeI(7.12 keV), and NiI(8.34 keV), and the column density toward 4U 1630−47 is

quite high (see above), the results of Pinto et al. (2013) suggest the possibility that the emission lines reported by Díaz Trigo et al. (2013) could in fact be an artefact of the model if the incorrect elemental abundance in the ISM is used in the fits. (Díaz Trigo et al. 2013, assumed solar abundance in their fits.)

In this chapter, we use the same XMM-NEWTONdata of 4U 1630−47 as Díaz Trigo

et al. (2013), but we explore an alternative model in which we allow the abundance of the ISM to be different from solar. We can fit the data well with a model that does not require any Doppler-shifted emission lines; instead, the fits yield non-solar abun-dances of S and Fe in the ISM along the line of sight to the source. The model not only fits the two observations in Díaz Trigo et al. (2013), but also the other XMM-NEWTONobservations during the 2012 outburst (Díaz Trigo et al. 2014), in

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2.2: Observations and data reduction 21

which four absorption lines, that we identify as being produced by FeXXV, FeXXVI

and Ni XXVIII(or Fe XXVLyβ), and two absorption K-edges, due to FeXXVand

FeXXVI, are detected. Furthermore, the putative Doppler-shifted emission lines are

not required either using the same model as Díaz Trigo et al. (2013) when we apply the new calibration files to those observations.

2.2 Observations and data reduction

The X-ray data that we used here consist of six observations of 4U 1630−47 with XMM-NEWTON (Jansen et al. 2001) taken between March 4 and September 28

2012. We report the details of the observations in Table 2.1. To reduce and anal-yse the raw data we used version 14.0.0 of the XMM-NEWTONScientific Analysis

Software (SAS) package following standard procedures.

We used the command epproc to calibrate the timing- and burst-mode photon event files. Following the recommendations of the XMM-NEWTONteam, for the

burst-mode data we also ran the command epfast. This command applies a correction to the energy scale due to charge transfer inefficiency in the CCD in burst mode. While there is some discussion in the literature regarding the applicability of this correction (see Walton et al. 2012, and the XMM-Newton Calibration Technical Note of November 2014∗_{), Díaz Trigo et al. (2013) applied this correction during their}

analysis and therefore, in order to compare to their results, we apply it here. For completeness, we also reduced the burst-mode observations without applying the

epfast correction.

We selected calibrated events with PATTERN≤4 in the central CCD of the EPIC-pn camera to get the spectrum of the source and we extracted a background spectrum from the outer columns of the central CCD (see Table 2.1 for the parameters of the extraction regions).

The difference between timing and burst mode in the process of extracting the data is that the spectra of the latter are influenced by the value of RAWY, i.e., the CCD row number (Kirsch et al. 2006). Following the recommendations of the XMM-NEWTONteam, we excluded events with RAWY > 140. We rebinned the average

EPIC-pn spectra before fitting in order to have a minimum of 25 counts in each bin. We created the redistribution matrix file (RMF) and the ancillary response file (ARF) using the SAS tasks rmfgen and arfgen, respectively. Following Díaz Trigo et al. (2013), we fitted the EPIC-pn spectra between 2 and 10 keV.

We used the spectral analysis package XSPEC v12.8.2 to fit the data, adding a 1% systematic error to the model to account for calibration uncertainties. The models that we used in this chapter include a component to account for photoelectric

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Figure 2.1 – X-ray spectra of ObsIDs 0670673101 and 0670673201 of 4U 1630−47 fitted with the model of Díaz Trigo et al. (2013). The second panel is the residuals of the best-fitting of ObsID 0670673101; the third panel is the residuals of ObsID 0670673201 when the strength of the three

GAUSSIANcomponents is set to zero. For this, and the other figures, the residuals are the data minus the model divided by the error.

.

tion of the interstellar material along the line of sight. For this component we used eitherTBABSorVPHABS; the latter allows variable abundances in the interstellar

ma-terial. For the emission component we used DISKBB, a multi-colour disc blackbody

(Mitsuda et al. 1984), POWERLAW, a simple power law, and GAUSS, to account for

possible Gaussian emission lines. We added several Gaussian absorption lines and edges (EDGE), when necessary. Throughout the chapter, we give the 1σ errors for all

fitted parameters and, when required, the 95% confidence upper limits.

2.3 Results

To compare the results with those of Díaz Trigo et al. (2013), we first fitted the two burst-mode observations, separately from the timing-mode observations, using the same calibration files that Díaz Trigo et al. (2013) used. We then fitted the model simultaneously to the two burst- and the four timing-mode observations. Finally we fitted the same model only to the burst-mode observations using the latest calibration. 2.3.1 Fits to the two burst-mode observations using the old calibration files Following Díaz Trigo et al. (2013), we first used the modelTBABS*(DISKBB+

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2.3: Results 23

Table 2.2–Best-fitting parameters for the two burst-mode observations of 4U 1630−47 based on the old calibration using two models

Model of Díaz Trigo et al. (2013) The model ObsID 0670673101 0670673201 0670673101 0670673201 TBABS/VHPABS NH(1022cm−2) 8.80 ± 0.05 8.85 ± 0.08 15.0 ± 0.2[1] 15.0 ± 0.2[1] S/S 1.0f 1.0f 1.32 ± 0.06[2] 1.32 ± 0.06[2] Fe/Fe 1.0f 1.0f 0.54 ± 0.07[3] 0.54 ± 0.07[3] DISKBB kTin(keV) 1.73 ± 0.01 1.75 ± 0.02 1.67 ± 0.01 1.68 ± 0.02 kdbb 90.7 ± 2.0 91.9 ± 3.6 107.7 ± 3.0 110.6 ± 3.8 POWERLAW Γ 2 f ₂f ₂f ₂f kpow <0.23 1.03 ± 0.14 <0.44 1.10 ± 0.09 GAUSS₁ E (keV) 4.02 ± 0.06[4] _{4.02 ± 0.06}[4] _7.0+0p −0.05[5] 7.0+0p−0.05[5] σ (eV) 165+47 −53[6] 165+47−53[6] 183+108−79 [7] 183+108−79 [7] kgau <2.3 5.8 ± 1.8 <2.3 1.61 ± 0.6 W (eV) <15.9 13.8 ± 3.6 <12.2 15.4 ± 6.9 GAUSS₂ E (keV) 7.24 ± 0.04[8] _{7.24 ± 0.04}[8] σ (eV) 165+47 −53[6] 165+47−53[6] kgau <0.9 3.0 ± 0.7 W (eV) <21.7 31.2 ± 6.7 GAUSS3 E (keV) 8.12 ± 0.10[9] _{8.12 ± 0.10}[9] σ (eV) 165+47 −53[6] 165+47−53[6] kgau <0.5 1.5 ± 0.6 W (eV) <42.8 23.7+7.5 −8.0 χ2

ν 0.98 for 256 d.o.f. 0.86 for 259 d.o.f.

NHis the column density of the neutral absorber.

S/Sand Fe/Feare, respectively, the sulphur and iron abundances, in solar units, of the

absorber along the line of sight.

kdbb, equal to the cosine of the inclination of the accretion disc with respect to the line of

sight times the square of the ratio of the inner radius of the disc in km and the distance to the

source in units of 10 kpc, kpow, in units of photons keV−1cm−2s−1at 1 keV, and kgau, in

units of 10−3_{photons cm}−2_s−1_{, are, respectively, the normalisation of the}_DISKBB_,

POWERLAWandGAUSSIANcomponents.

W is the equivalent width of the line.

Parameters with the same number in between square brackets were linked to be the same during the fit.

f _{This parameter was kept fixed at the given value during the fits.}

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Figure 2.2–X-ray spectra of ObsIDs 0670673101 and 0670673201 of 4U 1630−47 fitted simultane-ously with the alternative model that we proposed here. The second panel shows the residuals of the best-fitting model to the two observations; the following panels are the residuals of best-fitting model when the strength of theGAUSSIANcomponent is set to zero (third panel from the top), the abundance of S inVPHABSis set to solar (fourth panel from the top), and the abundance of Fe inVPHABSis set to solar (bottom panel).

reproduce their procedures as close as possible, we allowed NHto vary between the

two observations. Similar to Díaz Trigo et al. (2013), we found three emission lines in ObsID 0670673201, but not in ObsID 0670673101. In the latter case we calculated the upper limit to those lines assuming that they had the same energy and width as the lines in the other observation. We got values of the parameters that, except for

NHand the normalisation of the DISKBBand the GAUSS components, were similar

to those in Díaz Trigo et al. (2013). The difference in the normalisation is likely due to the fact that we considered the background spectrum in the analysis, whereas Díaz Trigo et al. (2013) did not. We give the best-fitting parameters for this model in Table 2.2, and we plot the X-ray spectra and best-fitting model of the two burst-mode observations in Figure 2.1. As in Díaz Trigo et al. (2013), to highlight the three emission lines we set their normalisations to zero in the residuals plot.

We then fitted an alternative model to the same data, in which we replaced theTBABS

component by theVPHABScomponent, and we kept only one of theGAUSSemission

components. The other components were the same as those in the model of Díaz Trigo et al. (2013). In this case we fitted the same model to both spectra simultane-ously and, since the interstellar absorption along the line of sight to the source should not change, we linked the parameters ofVPHABSbetween the two observations. We changed the default solar abundances model (ABUNDin XSPEC) to the abundances

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2.3: Results 25

Figure 2.3–X-ray spectra of all the six XMM-NEWTON_{observations of 4U 1636−47 fitted}

simulta-neously. The second panel from the top shows the residuals of the best-fitting model to all observations; the following panels are the residuals of each observation when the strength of the emission and ab-sorption components are set to zero. Each colour corresponds to one observation in Table 2.1, in the sequence black, red, green, blue, light blue, magenta and yellow. We do not show a residual panel for ObsID 0670673101 (magenta) because there are no emission or absorption components in this obser-vation.

of Wilms, Allen & McCray (2000) and the default photoelectric absorption cross-sections table (XSECT in XSPEC) to that given by Verner et al. (1996). We fixed

the photon index,Γ, of the POWERLAWcomponent at 2, since it could not be well

constrained in the fits.

One by one, we let the abundances of C, N, O, Ne, Mg, Si, S, Ca, Fe, and Ni in

VPHABSfree to fit the data, while the other element abundances were kept fixed at

the solar values. Except for the case of S and Fe, the best-fitting abundances were consistent with solar, and hence we eventually left the S and Fe abundances free and fixed all the other abundances to solar to fit the data.

The best-fitting model contains a moderately broad Gaussian line at 7 keV, consistent with the Lyα line of FeXXVI. A marginal detection of a similar line, likely due to

reflection off the accretion disc, had been previously reported in this source (Tomsick & Kaaret 2000; Tomsick et al. 2014). Since Fe reflection lines should appear between 6.4 keV (Fe I) and 6.97 keV (Fe XXVI), we constrained the line to be in the range

6.4–7 keV during the fits. The best-fitting value of the energy of the line pegged at the upper limit of this range, which could be in partly due to an imperfect calibration

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2

of the energy scale in burst-mode. The reduced χ2 _{of the model is} _χ2

ν =0.86 for

259 d.o.f. We plot the data and the best-fitting model in Figure 2.2, and show the best-fitting parameters of these two observations using the model in Table 2.2. The residual panels in this figure show the effect of the different parameters to the fit. The fit does not require any relativistically Doppler red- or blueshifted emission line. Instead, the abundance of S in VPHABSis higher than solar, 1.32±0.06, and that of

Fe is lower than solar, 0.54±0.07.

As a final check, we also created calibrated event files without applying the epfast correction and fitted the spectra with the model with variable sulphur and iron abun-dances. Since the best-fitting parameters in this case are consistent, within errors, with those from the fits to the spectra for which we did apply the epfast correction, we do not show a plot of this analysis. The only difference between the two sets of spectra is that in the model for the data for which we do not apply epfast we need to add an extra, marginally significant, emission line at 6.65 ±0.09 keV in the model of ObsID 0670673101.

2.3.2 Fits to the burst- and timing-mode observations using the old calibration We subsequently fitted the new model to all seven spectra simultaneously. A quick inspection of the residual plots indicated, in some cases, the presence of absorption features at energies of ∼ 6.5 keV or higher. Therefore we added up to four Gaussian absorption lines, using negativeGAUSSin XSPEC, and two edges,EDGEin XSPEC,

to account for possible absorption from highly ionised material close to the source. Not all these components were required in all observations. To keep the model as simple as possible, when the best-fitting parameters of these absorption compo-nents turned to be similar within errors, we linked these parameters across the ob-servations. The model we fitted,VPHABS*(DISKBB+GAUSS+POWERLAW-GAUSS1

-GAUSS2-GAUSS3-GAUSS4)*EDGE1*EDGE2, gives an acceptable fit, withχν2=0.99

for 895 d.o.f.. We show the best-fitting parameters in Table 2.3 and plot the spec-tra and best-fitting model of all the observations in Figure 2.3. In order to show the emission and absorption lines and edges in each observation, we set the strength of these components to zero in the residual panels (see Figure 2.3). In Figure 2.4 we show a zoom in of the residual panels of Figure 2.3 in the energy range 6-10 keV. Compared to the parameters in Kubota et al. (2007), Ró˙za´nska et al. (2014) and Díaz Trigo et al. (2014), we find a higher value of NH than theirs and the temperature of

the disc is higher than that of Kubota et al. (2007). The energy of the absorption lines and edges are consistent with those of Fe XXVHeα (6.70 keV), FeXXVILyα

(6.97 keV), Ni XXVIIILyα (8.09 keV) or Fe XXVLyβ (7.88 keV), Fe XXVI Lyβ

(8.25 keV), FeXXVK-edge (8.83 keV) and FeXXVIK-edge (9.28 keV), similar to

the identification in Díaz Trigo et al. (2014). However, the results are not identi-cal; e.g., Díaz Trigo et al. (2014) reported an absorption edge at 8.67 keV in ObsID

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2.3: Results 27 −4 −2 0 2 χ −15 −10 −5 0 χ −20 −15 −10−5 0 χ −15 −10−5 0 χ −15 −10−5 0 χ −10 −5 0 χ 6 7 8 9 10 −4 −2 0 2 χ Energy (keV)

Figure 2.4–All the residuals of the burst- and timing-mode observations except ObsID 0670673101.

0670671301 that is not required in the fits. Ró˙za´nska et al. (2014) detected seven iron absorption lines with SUZAKU, four of which are FeXXVHeα, FeXXVILyα,

FeXXVLyβ and FeXXVILyβ, the same ones we report here. The remaining two

ab-sorption lines identified by Ró˙za´nska et al. (2014) are FeXXVLyγ and FeXXVILyγ,

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Table 2.3 – Parameters of the emission and absorption lines and edges in the X M M -N EW T ON observ ations of 4U 1630 − 47. ObsID 0670671501-1 0670671501-2 0670671301 0670672901 0670673001 0670673101 0670673201 aN H (10 22 cm − 2) 14.1 ± 0. 1 [1 ] 14.1 ± 0. 1 [1 ] 14.1 ± 0. 1 [1 ] 14.1 ± 0. 1 [1 ] 14.1 ± 0. 1 [1 ] 14.1 ± 0. 1 [1 ] 14.1 ± 0. 1 [1 ] bS/S 1. 47 ± 0. 02 [2 ] 1. 47 ± 0. 02 [2 ] 1. 47 ± 0. 02 [2 ] 1. 47 ± 0. 02 [2 ] 1. 47 ± 0. 02 [2 ] 1. 47 ± 0. 02 [2 ] 1. 47 ± 0. 02 [2 ] bFe/Fe 0. 95 ± 0. 04 [3 ] 0. 95 ± 0. 04 [3 ] 0. 95 ± 0. 04 [3 ] 0. 95 ± 0. 04 [3 ] 0. 95 ± 0. 04 [3 ] 0. 95 ± 0. 04 [3 ] 0. 95 ± 0. 04 [3 ] GA USS e E (k eV) 6.89 ± 0. 02 [4 ] 6.89 ± 0. 02 [4 ] 6.89 ± 0. 02 [4 ] 6.89 ± 0. 02 [4 ] 6.89 ± 0. 02 [4 ] 6.89 ± 0. 02 [4 ] 6.89 ± 0. 02 [4 ] σ (eV) 169.1 + 13 .5 − 23 .1 [5 ] 169.1 + 13 .5 − 23 .1 [5 ] 169.1 + 13 .5 − 23 .1 [5 ] 169.1 + 13 .5 − 23 .1 [5 ] 169.1 + 13 .5 − 23 .1 [5 ] 169.1 + 13 .5 − 23 .1 [5 ] 169.1 + 13 .5 − 23 .1 [5 ] kgau 4.6 + 1. 0 − 0. 5 [6 ] 4.6 + 1. 0 − 0. 5 [6 ] 4.6 + 1. 0 − 0. 5 [6 ] 4.6 + 1. 0 − 0. 5 [6 ] 1.3 ± 0. 5 [7 ] < 0. 9 1.3 ± 0. 5 [7 ] W (eV) 113.4 + 57 .5 − 53 .4 102.0 + 51 .1 − 52 .8 84.1 + 40 .6 − 39 .8 81.6 + 36 .4 − 39 .8 14.9 ± 5. 4 < 12 .4 11.5 ± 4. 4 GA USS a1 E (k eV) 6.78 ± 0. 01 [8 ] 6.78 ± 0. 01 [8 ] 6.78 ± 0. 01 [8 ] 6.78 ± 0. 01 [8 ] 6.78 ± 0. 01 [8 ] 6.78 ± 0. 01 [8 ] 6.78 ± 0. 01 [8 ] σ (eV) 10.0 + 12 .2 − 10 .0 p [9 ] 10.0 + 12 .2 − 10 .0 p [9 ] 10.0 + 12 .2 − 10 .0 p [9 ] 10.0 + 12 .2 − 10 .0 p [9 ] 10.0 + 12 .2 − 10 .0 p [9 ] 10.0 + 12 .2 − 10 .0 p [9 ] 10.0 + 12 .2 − 10 .0 p [9 ] kgau 3.4 ± 0. 5 [10 ] 3.4 ± 0. 5 [10 ] 3.4 ± 0. 5 [10 ] 3.4 ± 0. 5 [10 ] 1.2 ± 0. 3 < 0. 3 < 1. 1 W (eV) 65.5 + 25 .9 − 24 .2 60.3 + 23 .2 − 19 .0 51.7 + 18 .5 − 17 .6 50.4 + 18 .9 − 17 .5 21.2 ± 7. 2 < 140 .9 < 20 .9 GA USS a2 E (k eV) 7.03 ± 0. 01 [11 ] 7.03 ± 0. 01 [11 ] 7.03 ± 0. 01 [11 ] 7.03 ± 0. 01 [11 ] 7.03 ± 0. 01 [11 ] 7.03 ± 0. 01 [11 ] 7.03 ± 0. 01 [11 ] σ (eV) 10.0 + 4. 3 − 10 .0 p [12 ] 10.0 + 4. 3 − 10 .0 p [12 ] 10.0 + 4. 3 − 10 .0 p [12 ] 10.0 + 4. 3 − 10 .0 p [12 ] 10.0 + 4. 3 − 10 .0 p [12 ] 10.0 + 4. 3 − 10 .0 p [12 ] 10.0 + 4. 3 − 10 .0 p [12 ] kgau 4.4 ± 0. 6 [13 ] 4.4 ± 0. 6 [13 ] 4.4 ± 0. 6 [13 ] 4.4 ± 0. 6 [13 ] 3.1 ± 0. 3 < 0. 3 < 0. 5 W (eV) 95.2 + 17 .4 − 9. 5 87.5 + 13 .9 − 10 .7 74.9 + 14 .6 − 8. 9 73.0 + 12 .5 − 9. 9 45.2 + 4. 4 − 3. 3 < 12 .0 < 10 .3 GA USS a3 E (k eV) 7.93 ± 0. 01 [14 ] 7.93 ± 0. 01 [14 ] 7.93 ± 0. 01 [14 ] 7.93 ± 0. 01 [14 ] 7.93 ± 0. 01 [14 ] 7.93 ± 0. 01 [14 ] 7.93 ± 0. 01 [14 ] σ (eV) 10.0 + 27 .6 − 10 .0 p [15 ] 10.0 + 27 .6 − 10 .0 p [15 ] 10.0 + 27 .6 − 10 .0 p [15 ] 10.0 + 27 .6 − 10 .0 p [15 ] 10.0 + 27 .6 − 10 .0 p [15 ] 10.0 + 27 .6 − 10 .0 p [15 ] 10.0 + 27 .6 − 10 .0 p [15 ] kgau 0.5 ± 0. 06 [16 ] 0.8 ± 0. 1 1.0 ± 0. 07 0.5 ± 0. 06 [16 ] 0.3 ± 0. 1 < 0. 2 < 0. 4 W (eV) 29.2 ± 4. 5 30.2 + 2. 7 − 3. 2 13.6 ± 2. 1 13.2 ± 2. 1 6.0 + 2. 5 − 2. 1 < 249 .0 < 146 .3 GA USS a4 E (k eV) 8.32 ± 0. 01 [17 ] 8.32 ± 0. 01 [17 ] 8.32 ± 0. 01 [17 ] 8.32 ± 0. 01 [17 ] 8.32 ± 0. 01 [17 ] 8.32 ± 0. 01 [17 ] 8.32 ± 0. 01 [17 ] σ (eV) 67.6 + 15 .2 − 12 .7 [18 ] 67.6 + 15 .2 − 12 .7 [18 ] 67.6 + 15 .2 − 12 .7 [18 ] 67.6 + 15 .2 − 12 .7 [18 ] 67.6 + 15 .2 − 12 .7 [18 ] 67.6 + 15 .2 − 12 .7 [18 ] 67.6 + 15 .2 − 12 .7 [18 ] kgau 1.2 ± 0. 1 [19 ] 1.2 ± 0. 1 [19 ] 1.2 ± 0. 1 [19 ] 1.2 ± 0. 1 [19 ] 0.8 ± 0. 1 < 0. 2 < 0. 5 W (eV) 52.3 + 3. 5 − 2. 4 47.1 + 3. 3 − 2. 3 38.0 ± 2. 4 36.9 ± 2. 4 18.1 ± 2. 4 < 287 .6 < 182 .2 ED GE 1 E (k eV) 8.63 ± 0. 02 [20 ] 8.63 ± 0. 02 [20 ] 8.63 ± 0. 02 [20 ] 8.63 ± 0. 02 [20 ] 8.63 ± 0. 02 [20 ] 8.63 ± 0. 02 [20 ] 8.63 ± 0. 02 [20 ] τ 0.12 ± 0. 01 [21 ] 0.12 ± 0. 01 [21 ] 0.04 ± 0. 01 [21 ] 0.04 ± 0. 01 [21 ] 0.04 ± 0. 01 [21 ] < 0. 03 < 0. 01 ED GE 2 E (k eV) 9.05 ± 0. 04 [22 ] 9.05 ± 0. 04 [22 ] 9.05 ± 0. 04 [22 ] 9.05 ± 0. 04 [22 ] 9.05 ± 0. 04 [22 ] 9.05 ± 0. 04 [22 ] 9.05 ± 0. 04 [22 ] τ < 0. 04 0.04 ± 0. 01 [23 ] 0.04 ± 0. 01 [23 ] 0.04 ± 0. 01 [23 ] < 0. 004 < 0. 02 < 0. 02 χ 2 ν 0.99 for 895 d.o.f. Notes. The GA U SSe and GA USS a1 to GA U SSa4 components represent the emission and absorption lines of Fe XXV He α ,Fe XX VI Ly α ,Ni XX VII Ly α or Fe X XV He β and Fe XXV I Ly β ,respecti vely .The E DGE 1 and E DGE 2 components indicate the absorption K-edges of Fe X XV and Fe XX VI ,respecti vely .See Table 2.2 for the definition and units of the parameters. The symbols used in this table ha ve the same meaning as in Table 2.2. As in Table 2.2, superscripts indicate parameters that were link ed between observ at ions during the fits.

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2.3: Results 29

We plot the equivalent width of the emission and absorption lines as a function of the total unabsorbed flux in the 2−10 keV range in Figure 2.5. From this figure it appears that the equivalent width of the emission and absorption lines is anti-correlated with the total unabsorbed flux. To test this we fitted both a constant and a linear function to each of these relations to assess whether the decreasing trend is significant. For the case of the absorption lines the F-test probabilities range from 3 × 10−6 _{to 7 ×}

10−2_{, indicating that in most cases the decrease of the equivalent width with flux is}

significant. For the emission line, however, the F-test suggests that a linear function is not significantly better than a constant. Since the emission and absorption lines in this part of the spectrum respond mostly to the high-energy flux, we also examined the plots of the line equivalent widths vs. the 7−10 keV flux; the trends are the same as those in Figure 2.5 for which we used the 2 − 10 keV flux, and hence we do not show those plots here.

In Figure 2.6 we plot some of the fitting parameters as a function of time. In the top panel of Figure 2.6 we show the 2-10 keV energy band from MJD 55970 to MJD 56220. The other panels show, respectively, the time histories of the temperature, the normalisation and the flux of theDISKBBcomponent, and the flux of thePOWERLAW

component. We do not plot the emissionGAUSScomponent in this figure because we

linked the parameters of this component across several observations (see Table 2.3). Figure 2.6 shows that the temperature of theDISKBBcomponent generally increases

with time, whereas the normalisation shows the opposite trend; on average, the flux of theDISKBBandPOWERLAWcomponents appear to increase with time. The

temper-ature of the disc and the disc and power-law fluxes are generally correlated, whereas the disc normalisation is anti-correlated, with the 2-10 keV MAXI flux. This is con-sistent with the standard scenario of black-hole states, but given the long time gaps between the observations, and the complex changes of the light curve with time (top panel of Figure 2.6), we do not discuss these correlations further. ThePOWERLAW

component in ObsID 0670673101 is not significant, and therefore we plotted the up-per limit as a triangle. From this figure it is apparent that the emission in the 2-10 keV range is always dominated by theDISKBBcomponent.

2.3.3 Fits to two burst-mode observations using the new calibration

While we were analysing these data, the XMM-NEWTONteam released a new set

of calibration files (dated March 31 2015) for Epic-pn burst-mode observations. We therefore extracted the burst-mode spectra again using the new calibration, and fitted the model to these two observations of 4U 1630−47. Comparing to the previous fit-ting results of the two burst-mode observations using the old calibration, we added an extra negative GAUSScomponent to the model to account for a possible absorption

line at ∼ 7 keV. We list the best-fitting parameters in Table 2.4. The F-test probabili-ties for the emission and absorption lines are, respectively, 10−2_{and 10}−3_{, indicating}

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2

Figure 2.5–The equivalent width of the broad emission and narrow absorption lines of 4U 1630−47 as a function of the total unabsorbed flux in the 2-10 keV range. Weand Wa1to Wa4represent, respectively,

the equivalent width of the emission and absorption lines of FeXXVHeα, FeXXVILyα, NiXXVIILyα

or FeXXVHeβ and FeXXVILyβ.

model of Díaz Trigo et al. (2013), but the three emission lines are not significantly detected. Finally, we also reduced these two burst-mode observations using the new calibration without applying the epfast correction. The best-fitting parameters are consistent, within errors, with those from the other fits, and in this case we do not find any significant emission line in ObsID 0670673201 either.

2.4 Discussion

Recently, Díaz Trigo et al. (2013) reported the detection of three Doppler-shifted emission lines arising from the jet of 4U 1630−47 in an XMM-Newton observation obtained during the 2012 outburst of the source. Here we show that this same obser-vation can be well fitted with a model that does not require the three emission lines. The main difference between the model and that of Díaz Trigo et al. (2013) is that we allow the abundances of S and Fe in the interstellar material along the line of sight to the source to vary. The model also fits well the other observation in Díaz Trigo et al. (2013), in which they do not detect the emission lines. Fitting these two observations simultaneously, we find that the abundances of S and Fe in the interstellar medium toward the source are, respectively, 1.32±0.06 and 0.54±0.07, in solar units. Be-cause of the large value of the column density in the interstellar medium toward the source, a non-solar abundance of these elements impacts upon the model at energies around the neutral Fe edge, at ∼ 7.1 keV (see the two bottom panels in Figure 2.2),

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2

2.4: Discussion 31

Figure 2.6 – Light curve and time history of some fitting parameters of 4U 1630−47. The upper panel shows the MAXI light curve (in units of photons cm−2_s−1_{) in the 2-10 keV band. The green}

dashed-dotted and red dashed vertical lines indicate the times of the two burst-mode and the four timing-mode XMM-NEWTONobservations, respectively. The second and third panels show, respectively, the temperature (in units of keV) and normalisation of theDISKBBcomponent. The fourth and fifth panels show the 2-10 keV unabsorbed flux, in units of 10−9_{erg cm}−2 _s−1_{, of the}_DISKBB_and_POWERLAW

components, respectively. The triangle in the fifth panel denotes the 95% confidence upper limit of the flux of thePOWERLAWcomponent in that observation. Some of the error bars are too small to show up on this plot.

such that the emission lines are no longer required in the model. The model also fits the rest of the XMM-Newton observation of the 2012 outburst of 4U 1630−47 (Díaz Trigo et al. 2014); in this case, similar to Díaz Trigo et al. (2014), we need to add several absorption lines and edges in the 6.7–9.1 keV energy range, likely due to photo-ionised material close to the source.

Since the two models are fundamentally different, we cannot compare them from a statistical point of view (e.g., using the F-test); however, the model fits the same data with less free parameters, and it is therefore simpler than the one of Díaz Trigo et al. (2013). The model does include a moderately broad (σ = 183+108

−79 eV) emission line

at 7 keV. This line is consistent with a marginally significant line detected from this (Tomsick & Kaaret 2000; Tomsick et al. 2014; King et al. 2014) and other sources (e.g. Miller 2007), which is usually interpreted as due to emission from the hard (power-law) component reflected off the accretion disc, with the broadening being due to relativistic effects close to the black hole (e.g. Fabian et al. 2012). The fact that in the fits the best-fitting energy of this line pegs at the upper limit, 7 keV, that we imposed in the model may be partly due to the uncertainties in the energy calibration

University of Groningen Disc reflection in low-mass X-ray binaries Wang, Yanan

Disc reflection in low-mass X-ray binaries

Wang, Yanan

Disc reflection in low-mass X-ray binaries

Proefschrift

Yanan Wang

Contents

1

1

Thesis introduction

1.1 X-ray binaries

1

1.2 Components of binary systems

1

1.3 X-ray spectra

1

1

in-1

1

1.4 X-ray states

X-1

1.5 X-ray facilities

1

1

1

1

1

1.6 Thesis motivation and overview

1

1

2

2

The XMM-Newton spectra of the 2012 outburst of the

black-hole candidate 4U 1630–47 revisited

2

Abstract

2

2.1 Introduction

2

2

2.2 Observations and data reduction

2

2.3 Results

2

2

2

2

2

2

2

2

2.4 Discussion

2