Download High-mass X-ray Binaries: X-raying the winds Ángel Giménez García

High-mass X-ray Binaries:
X-raying the winds
Ángel Giménez García
High-mass X-ray Binaries:
X-raying the winds
Ángel Giménez García
Tesis Doctoral
Universidad de Alicante
Noviembre 2015
High-mass X-ray Binaries:
X-raying the winds
Memoria que presenta Ángel Giménez García para optar al título de Doctor en
Física Aplicada a las Ciencias y las Tecnologías
Dirigida por los Doctores
Jose Miguel Torrejón Vázquez
Guillermo Bernabéu Pastor
Instituto de Física Aplicada a la Ciencia y la Tecnología
Departamento de Física, Ingeniería de Sistemas y Teoría de la Señal
Universidad de Alicante
Esta tesis ha sido posible gracias a la financiación del Ministerio de Ciencia e
Innovación mediante la beca FPI BES-2011-050874, dentro del proyecto de
investigación AYA2010-15431.
c Ángel Giménez García
Copyright Resumen en Lengua Oficial
Comunidad Valenciana
Este capítulo es un resumen en castellano de la Tesis Doctoral. No se incluyen tablas ni
figuras ya que haremos referencia a las que se muestran más adelante en el cuerpo general
del documento.
El objetivo principal de esta tesis es estudiar el medio circunestelar de las estrellas masivas (región inmediatamente cercana a la estrella y que se extiende unos pocos radios
estelares), cuyo conocimiento es fundamental para explicar entre otras cosas la retroalimentación de material en las galaxias, su enriquecimiento químico, o el ritmo de formación
estelar. Este tipo de estrellas tiene una vida corta en términos astronómicos y termina en
forma de supernova. Esta manera de vivir y morir produce un enriquecimiento químico del
medio y desencadena la formación de nuevas estrellas a partir de la compresión de nubes
moleculares cercanas por la fuerza de la explosión. Es decir, la evolución de las estrellas
masivas es un elemento clave en el entendimiento del pasado, presente y futuro de nuestra
galaxia y cualquier otra.
Cuando las estrellas masivas se encuentran en sistemas binarios junto a un objeto compacto, es posible que éste último pueda atraer suficiente abundancia de material como para
producir una ingente cantidad de rayos X. Esta radiación altamente energética se propaga
por el medio y nos transmite importante información sobre sus propiedades. Por lo tanto,
para conocer las propiedades del medio circunestelar de estrellas masivas la emisión de
rayos X nos abre una nueva ventana a través de la cual adquirir nueva información y contrastar la visión que nos ofrecen otras longitudes de onda menos energéticas, cuyo espectro
está dominado por la radiación directamente emitida por la estrella masiva. Esta Tesis Doctoral intenta aprovechar todas estas ventanas para hacer un análisis multi-frecuencia del
medio circunestelar en Binarias de Rayos X de Alta Masa.
Introducción
Un sistema binario de rayos X está formado por un objeto compacto (una enana blanca,
una estrella de neutrones o un agujero negro), acretando material proveniente de una estrella
v
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Resumen en Castellano
compañera. Éstas suelen recibir el nombre de compañera óptica o normal, por emitir su energía alimentadas por las reacciones nucleares que se dan en su interior, lo que las suele
hacer brillantes en el óptico.
El tipo espectral de la componente normal condiciona fuertemente la transferencia de
masa hacia el objeto compacto, por lo que habitualmente los sistemas binarios de rayos X
se separan entre los de alta masa (HMXB, High Mass X-ray Binaries), donde la compañera
óptica es de tipo O o B; y los de baja masa (LMXB, Low Mass X-ray Binaries).
Las HMXBs eran tradicionalmente clasificadas a su vez según la clase espectral de la
compañera óptica. Por lo tanto, se dividían en dos grupos principales: sistemas albergando
una estrella supergigante o sistemas con una estrella Be de secuencia principal. Sin embargo, el descubrimiento de nuevas fuentes a lo largo de las últimas dos décadas ha propiciado que sea necesaria una revisión de esta clasificación tradicional, atendiendo a la amplia
gama de peculiaridades que se observan en los sistemas. A continuación, exponemos las
clases en las que actualmente podemos dividir las HMXBs y sus principales características.
Clasificación de las Binarias de Rayos X de Alta Masa
Las BeXBs
1
Aunque inicialmente se creyó que las BeXB eran casos atípicos dentro de las binarias de
rayos X, actualmente no hay duda de que es el tipo de sistema dominante entre las HMXB.
En ellas la estrella normal es una Be con clase de luminosidad entre III-V. La e significa
que en algún momento de su evolución estas estrellas han presentado líneas espectrales en
emisión, las cuales ponen de manifiesto la presencia de un disco ecuatorial cuyo mecanismo
de formación no está completamente entendido actualmente (Porter and Rivinius, 2003),
aunque se cree que está relacionado con el hecho de que este tipo de estrellas tienen sistemáticamente una rotación rápida.
Cuando las estrellas Be forman parte de un sistema BeXB, el objecto compacto puede
pasar cerca o a través del disco durante el periastro, causando su descomposición y propiciando la transferencia de masa que suministra energía para la emisión de rayos X.
Las γ Cassiopeae like
2
γ Cassiopeae fue descubierta por Pietro Angelo Sechi en 1867, convirtiéndose en la primera estrella de clase Be. Forma parte de un sistema binario con un periodo orbital de 204
días y baja excentricidad. De la componente secundaria se sabe poco, más allá de que tiene
una masa similar a la solar. En el infrarrojo, óptico y ultravioleta, es una típica Be. Sin embargo, su comportamiento en los rayos X no es nada habitual. Su espectro en altas energías
está dominado por emisión térmica, con temperaturas del plasma emisor de ∼ 1.4 × 108 K
1
2
Referencias: Reig (2011), Ziolkowski (2002).
Referencias: Lopes de Oliveira et al. (2010).
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vii
(∼ 12 keV), responsable del ∼ 80 − 90% del flujo total. Tiene una luminosidad variable
en cualquier escala ponderable de tiempo (desde algunos segundos a varios minutos), con
valores de ∼ 1032−33 erg/s, pero sin grandes estallidos como en las BeXB.
El descubrimiento, en los últimos años, de 6 nuevas fuentes con propiedades análogas a
γ Cass, ha motivado la definición de una nueva clase de estrellas Be con emisión de rayos
X: las γ Cassiopeae like, con los siguientes atributos:
• Emisión térmica de temperatura característica > 5 keV.
• Variabilidad en cualquier escala de tiempo
• Luminosidad de ∼ 1033 erg/s en rayos X (de 0.2 − 12 keV)
Las clásicas SGXBs
3
Las clásicas SGXBs (Supergiant X-ray Binaries) son fuentes con una compañera óptica
supergigante y un rango dinámico (ratio entre luminosidad máxima y mínima) más moderado que las fuentes transitorias descritas más abajo. Antes del lanzamiento de INTEGRAL,
este tipo de fuente era la única variedad conocida de HMXB con una compañera supergigante, y por lo tanto son comúnmente conocidas como clásicas.
En las SGXBs, el objecto compacto está inmerso en el fuerte viento estelar de una compañera OB supergigante, capturando todo lo que entra en su dominio gravitacional. El
objecto compacto se encuentra a una corta separación orbital de ∼ 1.5 − 2 R⋆ , por lo que es
capaz de atrapar suficiente materia para alimentar una emisión persistente en rayos X que
alcanza los ∼ 1036−37 erg/s. Asimismo, estas fuentes pueden sufrir eventualmente cambios
abruptos en su luminosidad, indicando fuertes cambios en el ritmo al que el objeto compacto
acreta. Estos cambios pueden estar producidos por variaciones repentinas en la densidad del
viento (Kreykenbohm et al., 2008; Martínez-Núñez et al., 2014) o por inestabilidades sobre
la magnetosfera de la estrella de neutrones (Shakura et al., 2012a).
Las SFXTs
4
El número de SGXBs antes del lanzamiento de INTEGRAL rondaba la docena. La
escasez de fuentes era generalmente atribuida a la corta duración de la fase de supergigante
en estrellas masivas. INTEGRAL triplicó el número de HMXBs conocidas con compañera
supergigante y reveló la existencia de nuevas fuentes muy distintas a los sistemas clásicos:
las SFXTs (Supergiant Fast X-ray Binaries).
Las SFXTs son sistemas con componente óptica supergigante, como las SGXBs, pero
están caracterizadas por su comportamiento extremadamente variable. El rango dinámico
3
4
Referencias: Walter et al. (2015).
Referencias: Negueruela et al. (2008a), Chaty (2011), Sidoli (2013), Walter et al. (2015).
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de las SGXBs es . 2 órdenes de magnitud, pero en las SFXTs puede alcanzar los 5-6 órdenes de magnitud en los casos más extremos.
La razón del comportamiento tan distinto entre SGXBs y SFXTs es todavía una incógnita. Negueruela et al. (2008b) sugirieron la posibilidad de que la inhomogeneidad propia
de los vientos de estrellas calientes, unida con distintas configuraciones orbitales, pudiera
ser la causa. Sin embargo, el corto periodo orbital en algunas SFXTs es contradictorio con
esta propuesta (Walter et al., 2015). Bozzo et al. (2008) propusieron que la extrema variabilidad en las SFXTs podría explicarse por su especial configuración de radios de corrotación,
de acreción y de Alfven. Dependiendo de la extensión de estos radios, los cuales dependen
de las condiciones del viento estelar, parámetros orbitales y campo magnético de la estrella
de neutrones, la acreción de materia podría verse parcialmente inhibida, reduciendo drásticamente la luminosidad de las fuentes. Por último, la teoría de acreción quasi-esférica
sobre estrellas de neutrones desarrollada por Shakura et al. (2012b) hace una descripción
detallada del proceso de acreción sobre estrellas de neutrones bajo determinadas circunstancias. En esta teoría, la luminosidad en rayos X viene determinada por la eficiencia en los
procesos de enfriamiento que se producen por encima de la magnetosfera.
Las HMGBs
5
Las HMGBs (High-mass γ-ray binary systems) son sistemas que tienen su pico de
emisión por encima de 1 MeV. Hoy en día se cree que esta emisión está producida por
partículas aceleradas en el choque que se da cuando el viento del pulsar colisiona con el
viento estelar. Según esta visión, estos sistemas estarían alimentados por la energía rotacional de la estrella de neutrones, en contraste con el resto de HMXBs, que están alimentadas por acreción. Hay actualmente 5 HMGBs confirmadas, todas ellas con una componente óptica de secuencia principal.
Otras fuentes
Finalmente, hay unas pocas fuentes conocidas que por una serie de razones no se pueden
catalogar dentro de ninguna de las anteriores clases de HMXB. Entre las que hemos estudiado en esta tesis, tenemos 4U 2206+54, Centaurus X-3 y Cygnus X-1. La estrella óptica
en 4U 2206+54 es una O9.5V (Blay et al., 2006), es decir, ni una Be ni una supergigante.
Las donantes en Centaurus X-3 y Cygnus X-1 son supergigantes, pero en ellas la acreción se
produce a través de un disco de acreción (Tjemkes et al., 1986; Shapiro et al., 1976), lo que
modifica completamente su comportamiento en los rayos X. Además, el objeto compacto
en Cygnus X-1 es un agujero negro, lo que conlleva particularidades adicionales.
La acreción como fuente de energía
6
La emisión de rayos X se produce a consecuencia del proceso de acreción, consistente
en la adición de cierta cantidad de masa, inicialmente situada en las inmediaciones del objeto compacto, sobre la superficie del mismo. La energía gravitatoria que inicialmente posee
5
6
Referencias: Dubus (2013).
Referencias: Frank et al. (2002)
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la masa sedimentada es transformada sucesivamente en energía cinética, energía térmica, y
finalmente liberada por radiación.
Para que este fenómeno tenga importancia deben darse las circunstancias especiales que
tenemos en los sistemas binarios de rayos X. Vamos a verlo con el caso más simple: supongamos que tenemos un objeto esférico de masa M y radio R, y una masa m inicialmente en
reposo cayendo desde el infinito, hasta que finalmente es depositada sobre la superficie del
objeto. En ese escenario tenemos que la partícula de masa m experimenta un cambio en su
energía mecánica:
Z R
Z R
GMm
GMm
~
~
∆E =
F · dr =
dr = −
= −∆Eac
(1)
2
R
r
∞
∞
Es decir, la partícula ha perdido energía mecánica, y hemos llamado a esta pérdida ∆Eac .
Para ver cuándo esta cantidad es significativa, podemos compararla con la energía que se
puede generar a partir de la misma masa m mediante reacciones nucleares. El máximo se
obtiene si, como es habitual en astrofísica, el material es inicialmente hidrógeno y se fusiona
en helio:
∆Enuc = 0.007mc2
donde c es la velocidad de la luz. Por tanto,
GM
M R⊙
∆Eac
=
= 3.03 10−4
2
∆Enuc 0.007c R
M⊙ R
Como podemos ver, para que la acreción sea importante como fuente de energía serán favorables las situaciones en que, o bien R ≪ R⊙ , o bien M ≫ M⊙ . Tal es el caso de los sistemas
binarios de rayos X, en que el cuerpo acretante es un objeto compacto como una estrella de
neutrones o un agujero negro, que típicamente pueden tener M > 1M⊙ , R 6 10−4 R⊙ .
Si en lugar de una partícula individual de masa, lo que tenemos es un ritmo constante de
materia Ṁ cayendo hacia el objeto esférico, y suponemos que toda la energía mecánica que
se pierde (ec. 1) es emitida, tendremos una luminosidad de acreción:
Lac =
GM Ṁ
R
(2)
Resultados
Análisis completo de FeKα en Binarias de Rayos X de Alta Masa con XMMNewton
Esta parte de la tesis está dedicada a un estudio de FeKα en toda la muestra de observaciones de HMXBs disponible en la base de datos del observatorio espacial XMM-Newton,
y su contenido principal ha sido publicado en la revista de arbitraje internacional Astronomy
& Astrophysics con la siguiente referencia:
• An XMM-Newton view of FeKα in HMXBs, Giménez-García, A., Torrejón, J.M., Eikmann, W., et al. 2015, A&A, 576, A108
A continuación exponemos los puntos fundamentales de este estudio.
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Resumen en Castellano
FeKα y el complejo del Fe
Las líneas espectrales que frecuentemente se observan en HMXBs alrededor de los 67 keV es lo que comúnmente se denomina complejo del Fe. Los rasgos espectrales en este
intervalo de energía se atribuyen a líneas de emisión del Fe, las cuales pueden estar producidas por fluorescencia o por recombinaciones radiativas de Fe.
La fluorescencia se produce cuando un átomo, expuesto al campo radiación originado
por la fuente de rayos X, absorbe un fotón con tanta energía que es capaz de hacerle perder
uno de sus electrones más interiores y fuertemente ligados (ver Fig. 3.1). En ese momento
el átomo se encontrará en un estado inestable, por lo que finalmente decaerá a un estado
en que el lugar del electrón desalojado es ocupado por un electrón de una capa superior.
Como la capas capas interiores están más fuertemente ligadas al núcleo, durante el proceso
se libera energía, ya sea expulsando otro electrón (efecto Auger), o bien emitiendo un fotón.
Cuando el átomo de que hablamos es Fe, el primer electrón expulsado es el más interior
(capa K), y el electrón que ocupa su lugar se encontraba en una de las dos capas inmediatamente superiores en el momento de la absorción (capas L y M), el fotón que se emite (si se
emite) tendrá una energía que dependerá del estado de ionización. Fotones producidos de
esta forma son los que forman la línea fluoresecente de Fe Kα (M → K) o de Fe Kβ (L → K).
El estudio del complejo del Fe, y de la línea de FeKα en particular, es una herramienta
fundamental para la comprensión del medio que rodea las fuentes de rayos X, ya que nos
proporciona información sobre su composición y estado físico (densidad, temperatura, estado de ionización...).
Las observaciones y su tratamiento
Para este estudio hemos recopilado toda la información disponible de HMXBs en la base
de datos del observatorio XMM-Newton anterior a Agosto de 2013. La instrumentación a
bordo de este observatorio es la más adecuada actualmente para realizar este análisis, debido al gran área efectiva de sus telescopios (permitiendo el estudio de un gran número de
fuentes) y a la moderada aunque suficiente resolución espectral de sus cámaras CCD (que
permiten resolver varias componentes dentro del complejo del Fe).
Siguiendo el catálogo de HMXBs elaborado por Liu et al. (2006), complementado con
la inclusión de nuevos descubrimientos o confirmaciones desde su publicación, hemos
recogido información de 46 fuentes, entre las cuales detectamos FeKα en 21 de ellas.
Mostramos la lista de fuentes en la Tabla 4.1. Algunas de las fuentes has sido observadas
por XMM-Newton en más de una ocasión, por lo que el total de espectros analizados es 108
(ver Apéndice A).
Para la reducción hemos utilizado SAS 11.0.0 (Science Analysis System), un conjunto
de tareas, scripts y librerías específicamente diseñadas para la reducción de datos recogidos
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xi
por XMM-Newton7 . Todos los ficheros relativos a cualquier observación pública pueden
descargarse del archivo científico de XMM-Newton (XSA) en http://xmm.esac.esa.int/xsa/.
Para el análisis de los espectros, una vez reducidos con SAS, hemos utilizado el programa
XSPEC, versión 12.8.08 , especializado en el ajuste de modelos a espectros de rayos X. En
la Tabla 4.2 presentamos todos los modelos del continuo que hemos probado en cada uno
de los espectros analizados. Los modelos has sido aceptados o rechazados en función del
χ2 -reducido asociado a cada ajuste espectral. En la Fig. 4.2 podemos ver que en la mayoría
de los ajustes tenemos χ2 -reducido ≃ 1, como se espera para un buen ajuste.
Para el ajuste de las líneas de emisión hemos utilizado perfiles Gaussianos. Entre las
líneas de emisión que hemos detectado, exigimos las siguientes condiciones para categorizar
una línea como FeKα:
1) El centroide de energía de la componente Gaussiana está dentro del intervalo
[6.3, 6.65] keV.
2) La significación estadística (σ sign ) de la componente Gaussiana es mayor que 2σ.
Calculamos σ sign a partir de χ2k1 − χ2k2 , asumiendo χ2k1 − χ2k2 ∼ χ2k1 −k2 9 , donde χ2k1
está calculado para un modelo incluyendo la componente Gaussiana, y χ2k2 para un
modelo sin esa componente.
Atlas de espectros
En el Apéndice A mostramos el conjunto completo de espectros analizados junto con el
modelo de mejor ajuste en cada uno de ellos. Asimismo, los parámetros derivados de estos
ajustes son mostrados en las Tablas A.1 y A.2.
Si observamos estos espectros, podemos ver tres tipos de patrones característicos en el
complejo del Fe, que hemos definido como Tipos I, II y III (ver Fig. 5.1). En el Tipo I
observamos las líneas de fluorescencia FeKα y FeKβ, pero no de recombinación Fe xxv y
Fe xxvi. En el Tipo II detectamos las líneas de fluorescencia y recombinación. Por último,
en el Tipo III no podemos detectar líneas de emisión. Cada uno de los grupos de HMXBs
muestra un tipo de patrón característico, que hemos recogido en la Tabla 5.1 junto con otras
características fundamentales de sus espectros como el tipo de modelo utilizado (térmico,
no térmico o una combinación de ambos) y nivel de absorpción de los rayos X.
Anchura de la línea
En la Tabla A.2 mostramos todos los parámetros de cada detección de FeKα, incluyendo
la anchura de la línea (σline ). Hemos hecho una distinción entre líneas estrechas (σline <
0.15 keV) y anchas (σline > 0.15 keV). Esta separación está motivada tanto físicamente
7
Referencias: http://xmm.esa.int/external/xmm_user_support/documentation/sas_usg/USG/
http://heasarc.gsfc.nasa.gov/xanadu/xspec/manual/manual.html
9
Esta hipótesis no es estrictamente cierta, ya que χ2k1 y χ2k2 no son independientes. Sin embargo, nos da una
estimación del impacto que la componente Gaussiana tiene en el modelo
8
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como observacionalmente. La gran mayoría de detecciones son líneas estrechas de FeKα
(ver Fig. 5.2). La poca anchura de estas líneas es incompatible con material rotando a altas
velocidades como las que esperaríamos en discos de acreción. De aquí en adelante, cuando
nos refiramos a FeKα será en relación a las líneas estrechas.
Centroide de energía
En la Fig. 5.3 mostramos un histograma con los centroides de energía de FeKα. El valor
medio está en 6.42±0.02 keV. No hay diferencias fundamentales entre los valores obtenidos
en las distintas clases de HMXBs o en distintos estados de luminosidad. Estos valores
implican que el estado de ionización del Fe es menor que Fe xviii (Kallman et al., 2004), en
consonancia con estudios previos en HMXBs (Torrejón et al., 2010b; Gottwald et al., 1995;
Nagase, 1989). En este sentido, el estudio de Torrejón et al. (2010b), usando un instrumento
con mejor resolución espectral, permite acotar el estado de ionización a Fe i-x. Nuestro
trabajo apoya este resultado aportando nuevas fuentes a la muestra.
Correlación de parámetros
Uno de los objetivos de estudiar una amplia muestra de fuentes es poder encontrar correlaciones entre parámetros que nos ayuden a comprender mejor lo que ocurre en el medio
circunestelar de los sistemas. A continuación exponemos las correlaciones que hemos detectado:
• El flujo de energía del continuo está directamente correlacionado con el flujo de
FeKα, como mostramos en la Fig. 5.4. Este resultado es el esperado, ya que una
mayor intensidad del campo de radiación produce un mayor número de transiciones
en los átomos de Fe y por lo tanto más fluorescencia.
• La anchura equivalente (EW) de FeKα, está inversamente correlacionada con la luminosidad del continuo (Fig. 5.5). Este fenómeno había sido observado anteriormente
en Binarias de Rayos X por Torrejón et al. (2010b), y recibe el nombre de efecto Baldwin debido a su similitud con la correlación inversa encontrada por Baldwin (1977)
en la EW de C iv y la luminosidad UV en AGNs. Curiosamente, las fuentes análogas
a γ Cassiopeae no siguen este patrón, lo que sugiere que el escenario de reprocesamiento de las líneas de fluorescencia en este tipo de fuentes es muy diferente al de
las SGXBs.
• Hemos encontrado correlación directa entre el centroide de energía de la línea y σline ,
como mostramos en la Fig. 5.6. Esto puede indicar cierto solapamiento de líneas
de FeKα de distinto grado de ionización y por lo tanto ligeramente desplazadas en
energía.
• Por último, hemos visto el crecimiento de la EW de FeKα con el aumento de la
absorpción en rayos X (NH ), como podemos comprobar en la Fig. 5.7. Esta correlación, también llamada curva de crecimiento, es similar a la la encontrada por
Torrejón et al. (2010b), así como a la que esperamos a partir de simulaciones de una
fuente de radiación de rayos X que se transmite a través de un medio isotrópicamente
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xiii
distribuido de elementos neutrales (Eikmann, 2012). La concordancia de las observaciones con las simulaciones indican que la fluorescencia se produce como consecuencia de la transmisión de rayos X por el medio circunestelar, y no por la reflexión
de esta radiación en un medio independiente como podría ser un disco de acreción
(ver Fig. 6.2).
NH : SGXBs and SFXTs
En la Fig. 5.8 podemos ver los histogramas de NH observados en SGXBs y SFXTs. Se
puede ver con claridad que las SGXBs presentan por lo general mayor absorpción que las
SFXTs, indicando una mayor densidad del medio circunestelar. Haciendo un simple análisis estadístico comprobamos que esta tendencia no se debe a una disposición aleatoria de
los valores de NH , sino que lo más probable es que haya alguna razón física que lo explique.
Entre las posibles razones está la interacción del objeto compacto con el viento estelar,
que como muestran las simulaciones hidrodinámicas (e.g. Manousakis and Walter, 2011;
Manousakis et al., 2012), pueden variar la densidad del medio en la línea de visión del observador y por lo tanto modificar el NH . También es posible que, asumiendo que la absorpción se produce principalmente en el viento estelar, la discrepancia en el NH indique ciertas
diferencias fundamentales en los vientos estelares de los dos grupos de sistemas. Este resultado dio pie a las investigaciones que llevamos a cabo más adelante en la Parte III de esta
tesis, donde hacemos un análisis de los vientos estelares en dos fuentes muy representativas:
Vela X-1 (SGXB) e IGR J17544-2619 (SFXT).
Análisis individuales: IGR J16320-4751 y 4U 1700-37
Dentro de la muestra de observaciones, hemos contado con dos fuentes con un mayor
número de apuntados, lo que nos ha permitido desarrollar un estudio más detallado de las
mismas. Estas fuentes son IGR J16320-4751 y 4U 1700-37.
Al disponer de varias observaciones cubriendo distintas fases orbitales, hemos podido
realizar un estudio de la modulación de NH con la fase orbital. Para ello, hemos calculado
la absorpción teóricamente esperada en un viento estelar homogéneo, asumiendo una ley
de velocidad β (Castor et al., 1975a). A pesar de la simplicidad del modelo, las observaciones son razonablemente compatibles con que una modulación producida por absorpción
del viento estelar (Fig. 5.10 y 5.12).
El análisis de estas fuentes está en concordancia con los resultados expuestos anteriormente de los análisis de la muestra completa. La curva de crecimiento es nuevamente
compatible con las simulaciones, como mostramos en las Fig. 5.9 y 5.11.
Análisis comparativo de dos donantes supergigantes en Binarias de Rayos X
de Alta Masa: la persistente Vela X-1 y la transitoria IGR J17544-2619
Como describimos anteriormente, las fuentes clásicas SGXBs y las SFXTs son sistemas
con una donante similar, pero al mismo tiempo tienen un comportamiento muy distinto en
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los rayos X. La razón de esta dicotomía es todavía una incógnita. Se han puesto sobre
la mesa distintas teorías, pero la mayoría de ellas necesitan asumir ciertos valores de los
parámetros del viento estelar, como la pérdida de masa y la velocidad terminal. Sólo mediante estudios empíricos dedicados al análisis de las atmósferas estelares de las estrellas
donantes podemos proveer a estas teorías con la información necesaria para ser debidamente
corroboradas. Sin embargo, estudios de este tipo en HMXBs son más bien escasos.
Con la intención de avanzar en esa dirección, hemos realizado un estudio comparativo
pormenorizado de las compañeras ópticas en dos fuentes muy representativas de sus respectivas clases: IGR J17544-2619 (SFXT) y Vela X-1 (SGXB). Hemos usado datos de archivo
en el infrarrojo, óptico y ultravioleta y los hemos analizado usando el código Potsdam WolfRayet (PoWR) de simulación de atmósferas de estrellas masivas, usando ecuaciones de noLTE y la influencia de los rayos X en las poblaciones atómicas.
En los siguiente apartados describimos las fuentes y exponemos los resultados fundamentales obtenidos y una discusión de los mismos. El contenido de este estudio está recogido
en un artículo enviado a la revista Astronomy & Astrophysics en Octubre de 2015.
Las fuentes
IGR J17544-2619 fue detectada en Septiembre de 2003 con el detector IBIS/ISGRI del
satélite INTEGRAL (Sunyaev et al., 2003). Su periodo orbital es de ∼4.9d (Clark et al.,
2009). El objeto compacto es una estrella de neutrones (in’t Zand, 2005). Pellizza et al.
(2006a) utilizó observaciones en óptico y NIR para clasificar la compañera óptica como
O9Ib. Observaciones con XMM-Newton y Suzaku muestran que el sistema tiene un rango
dinámico muy alto, cambiando su flujo en rayos X más de 4 órdenes de magnitud (Rampy et al.,
2009; González-Riestra et al., 2004).
Actualmente el periodo de rotación de la estrella de neutrones Pspin en IGR J17544-2619
es debatido ya que el análisis de distintas observaciones tomadas en momentos diferentes
dan lugar a resultados divergentes. Drave et al. (2012), analizando datos de RXTE en un
estado de luminosidad intermedio (∼ 1033−34 erg/s), encontró una pulsación en rayos X
de Pspin = 71.49s con una significación estadística de 4.37σ. Romano et al. (2015) utilizó observaciones de Swift durante un gran estallido de luminosidad (con un pico llegando
a los ∼ 1038 erg/s), y detectó pulsaciones de Pspin = 11.60s, también con una significación estadística de ∼ 4σ. Por otra parte, estos resultados contrastan con los análisis
de observaciones con XMM-Newton y NuSTAR llevadas a cabo por Drave et al. (2014) y
Bhalerao et al. (2015) respectivamente. Estos autores no encontraron señales de pulsación
en ninguna escala de tiempo entre 1-2000s.
Vela X-1 es una de las HMXBs más estudiadas, ya que es una fuente brillante descubierta en los primeros tiempos de la astronomía de rayos X (Chodil et al., 1967). El sistema
presenta una moderada excentricidad e= 0.09 (Bildsten et al., 1997), y un periodo orbital
Porb = 8.96 días (Kreykenbohm et al., 2008). El objeto compacto es una estrella de neutrones con un periodo de rotación Pspin = 283s (McClintock et al., 1976). La compañera
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óptica HD 77581 (B0.5Iae) fue identificada y clasificada por Vidal et al. (1973).
Resultados
Las Tablas 8.1 y 8.2 muestran el conjunto de observaciones que hemos utilizado. Los
espectros analizados y el mejor ajuste pueden verse en las Fig. B.1-B.4. El ajuste de la
distribución espectral de energía puede verse en las Fig. 10.4 y 10.10. Finalmente, los
parámetros estelares derivados de los ajustes son mostrados en la Tabla 10.1.
Los parámetros que hemos derivado del análisis está de acuerdo con la clasificación
espectral de las fuentes, a saber, O9I (IGR J17544-2619) y B0.5Iae (Vela X-1). La distancia a los sistemas ha sido revisada y también está en concordancia con las estimaciones
disponibles en la literatura. En IGR J17544-2619 hemos sido capaces de acotar más el error
en esta estimación, obteniendo una distancia d = 3.0 ± 0.2 kpc (ver Fig. 10.3). A partir
del radio estelar y su comportamiento en los rayos X, hemos visto que la solución orbital
obtenida por Clark et al. (2009) implica una excentricidad en IGR J17544-2619 e < 0.25.
Asimismo, las abundancias de algunos elementos químicos como N, C y O tanto en la
donante de IGR J17544-2619 como en la de Vela X-1 (ver Tabla 10.2) indican una cierta
evolución en la composición química de sus atmósferas. Esta evolución podría estar causada por episodios de transferencia de masa en el sistema binario (Langer, 2012), o por
procesos de mezcla con capas estelares interiores inducida rotacionalmente.
Discusión
Uno de los objetivos principales de este estudio es poder corroborar teorías propuestas
que puedan explicar el porqué de los diferentes rangos dinámicos en SGXBs y SFXTs. En
concreto, Bozzo et al. (2008) propone la posibilidad de que en las SFXTs la acreción esté
regularmente dificultada por mecanismos inhibidores como barreras centrífugas o magnéticas, lo que como hemos visto en la Ecuación 1 implicaría una menor cantidad de energía
disponible para ser finalmente irradiada como rayos X. Es decir, la luminosidad en estos regímenes de acreción es menor. En cambio, en las SGXBs la acreción se produciría
de manera directa, con una mayor luminosidad y mayor constancia. La acción de estos
mecanismos de inhibición dependería fundamentalmente de los parámetros del viento estelar (pérdida de masa Ṁ y velocidad relativa del viento respecto la estrella de neutrones 3rel ),
el momento magnético de la estrella de neutrones (µ), la separación orbital a y el periodo
de rotación de la estrella de neutrones Pspin .
Usando los valores de la Tabla 11.1 y las ecuaciones descritas en el trabajo de Bozzo et al.
(2008), podemos elaborar diagramas 3rel - Ṁ y 3rel -Pspin donde los distintos regímenes de
acreción ocupan distintas regiones de los diagramas.
En la Fig. 11.1 mostramos la posición de IGR J17544-2619 en el diagrama 3rel - Ṁ para
las dos estimaciones disponibles de Pspin . La fuente está situada en la región de acreción
directa para Pspin = 71.49s, y en la supersonic propeller regime para Pspin = 11.58s. En
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la Fig. 11.2 podemos ver que la inhibición de la acreción ocurre para periodos de rotación
más cortos que ∼ 30s. Por lo tanto, la estimación de Pspin = 11.58s coincide mejor con el
comportamiento en rayos X y su permanencia en regímenes de acreción inhibida.
Como podemos ver en la Fig. 11.1, variaciones relativamente modestas del viento de
IGR J17544-2619 pueden producir transiciones entre distintos regímenes de acreción. En
el caso de Pspin = 11.58s, necesitaríamos un incremento de un orden de magnitud en la
densidad del medio transitado por la estrella de neutrones. Estos saltos de densidad son
perfectamente factibles, como demuestran las simulaciones hidrodinámicas de vientos estelares impulsados radiativamente (e.g. Feldmeier et al., 1997). En fuentes como IGR J175442619, cambios abruptos en la densidad del viento podrían desencadenar la transición de un
régimen de acreción a otro.
En cuanto a Vela X-1, en la Fig. 11.3 podemos ver que la fuente se encuentra profundamente instalada en la zona de acreción directa. Por esa razón necesitaríamos cambios
mucho más extremos en la densidad o velocidad del viento estelar para que se produjera
una inhibición de la acreción. Estos cambios extremos son factibles, pero mucho menos
probables. Sin embargo, es plausible que la estrella de neutrones en Vela X-1 pueda transitar esporádicamente una región extraordinariamente poco densa que produzca un descenso
repentino en la luminosidad en rayos X.
Usando los parámetros expuestos en la Tabla 11.1 junto con la Ec. 11.1c, obtenemos que
la luminosidad de acreción en Vela X-1 es Lacc = 6.5 × 1036 erg/s. Este valor es muy cercano a la luminosidad media en rayos X: hLX i ≃ 0.7 × Lacc . Esta similitud demuestra que el
escenario de acreción directa es capaz de describir cualitativamente bien la manera en que
la materia es acretada en Vela X-1.
Recapitulando lo dicho en párrafos anteriores, podemos decir que el marco de trabajo
descrito por Bozzo et al. (2008) es capaz de explicar la tendencia de IGR J17544-2619 a
mostrar una gran variabilidad, en contraste con Vela X-1 y su persistente alta luminosidad en rayos X. Como se puede ver en las Fig. 11.1 y 11.3, la variabilidad necesaria en
el viento estelar para producir una transición de régimen de acreción es mucho menor en
IGR J17544-2619 que en Vela X-1. Los ingredientes principales de esta diferencia son el
Pspin (más corto en IGR J17544-2619), y la 3rel (mayor en IGR J17544-2619). Con esta
base, nos aventuramos a conjeturar si el comportamiento radicalmente diferente de SFXTs
y SGXBs se puede deber a su combinación de Pspin y 3rel . Es decir, si Pspin tiende a ser más
corto y 3rel mayor en las SFXTs que en las SGXBs.
En efecto, cierto número de evidencias parecen apuntar en esa dirección. A continuación
exponemos el conocimiento que actualmente se tiene del Pspin y la velocidad terminal del
viento estelar (3∞ ) en SGXBs y SFXTs:
• En SGXBs, tenemos que Pspin > 100s. La única excepción es OAO 1657-415 (Pspin =
38.2s). A pesar de que su corto Pspin difiulta la acreción directa, estos inconvenientes
son salvados por la extraordinariamente baja velocidad de su viento estelar (3∞ ∼
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250 km/s) de la donante Ofpe/WN9 (Mason et al., 2012). Como podemos ver en la
Fig. 11.4, con una velocidad tan baja del viento la acreción directa está prácticamente
garantizada con los niveles de pérdida de masa de este tipo de estrellas.
• Entre las SFXTs, hay dos sistemas que muestran Pspin < 100s: IGR J17544-2619 e
IGR J18483-0311 (Pspin = 21.05s, Sguera et al. 2007b). También existen detecciones
en SFXTs con Pspin > 100s. Sin embargo, según el reciente estudio de revisión de
Walter et al. (2015), en estos casos su definición como SFXTs es discutible, ya que
o bien su comportaniento es asimilable al de otros sistemas clásicos como GX 301-2
o Vela X-1, o bien carecen de observaciones de buena calidad. El Pspin del resto de
SFXTs permanece desconocido.
• El escaso número de estimaciones disponibles de 3∞ en SGXBs apunta a bajas velocidades del viento, similares a la obtenida en este trabajo para Vela X-1: 305 km/s en
GX 301-2 (Kaper et al., 2006), 500−600 km/s en IGR J17252-3616 (Manousakis et al.,
2012), 500 km/s en X1908+075 (Martínez-Núñez et al., 2015), y la ya mencionada
3∞ ∼ 250 km/s in OAO 1657-415. Otra SGXB, 4U1700-37, muestra una mayor
3∞ = 1700 km/s (van Loon et al., 2001), pero interesantemente, esta fuente es candidata a albergar un agujero negro debido a la imposibilidad de detectar pulsaciones
hasta el momento.
• El análisis que hemos llevado a cabo en la primera parte de la tesis indica que hay
una diferencia importante en la columna de absorción de rayos X entre SGXBs y
SFXTs, siendo en promedio significativamente menor en las SFXTs. Esta absorción
es fundamentalmente local en estas fuentes. Este resultado está de acuerdo con la
hipótesis de vientos más lentos y densos en las donantes de las SGXBs.
• En SFXTs, Lorenzo et al. (2014) asumió 3∞ = 1230 km/s en su análisis de IGR J112155952. No hay más estimaciones de la velocidad del viento en SFXTs aparte de este
trabajo.
En resumen, estamos aún lejos de un conocimiento completo de Pspin y 3∞ en SGXBs y
SFXTs. Sin embargo, hay razones para pensar que las SFXTs tienden a mostrar mayor 3rel
y más corto Pspin . Hacemos notar que un alto 3rel o un corto Pspin no son necesariamente
rasgos inherentes a las SFXTs por sí solos, sino que es la combinación de estos parámetros
lo que favorece cierto régimen de acreción particular en cada fuente (o la transición entre
regímenes). Además, la importancia relativa de Pspin y 3rel depende de manera importante
en el régimen de acreción considerado. Por ejemplo, la entrada en el régimen de subsonic
propeller desde la acreción directa es mucho más sensible a 3rel que la entrada al régimen de
supersonic propeller, como se puede ver en las líneas características de estas transiciones
en las Fig. 11.1-11.4. Dado que la intersección entre las líneas características de transición
a los regímenes subsonic y supersonic propeller tiene lugar para cierta velocidad relativa
3K , la velocidad del viento estelar es especialmente relevante siempre que 3rel > 3K . Cuando
3rel < 3K , la inhibición de acreción de tipo subsonic propeller no es posible y la acreción
directa sólo es interrumpida por una barrera centrífuga cuando el Pspin es suficientemente
corto.
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Asimismo, podemos hacer una comparación de 3∞ y la velocidad de escape (3esc ). En este
aspecto, Lamers et al. (1995) recogió una larga muestra de espectros de estrellas calientes,
concluyendo que el ratio 3∞ /3esc cae de manera pronunciada desde ∼ 2.6 a ∼ 1.3 aproxidamente alrededor del tipo espectral B1 cuando nos movemos de más a menos en temperaturas
efectiva (T eff ). Este descenso está explicado por un descenso en la aceleración de la línea
Fe iii en la parte subsónica del viento (Vink et al., 1999). En las fuentes aquí analizadas
tenemos:
• IGR J17544-2619 (O9.5I): 3∞ /3esc = 2.4+0.7
−0.5
• Vela X-1 (B0.5I): 3∞ /3esc = 1.6+0.8
−0.4
Estos valores siguen la tendencia observada y descrita por Lamers et al. (1995), por lo
que sugerimos que ésta puede ser la razón de los distintos 3∞ en las dos fuentes. La acción
de los rayos X puede tener una importante influencia en el viento, como demuestran los
estudios de Karino (2014). Sin embargo, este efecto es probablemente local, ya que no observamos diferencias significativas en 3∞ entre fases orbitales en eclipse y fuera de eclipse
en Vela X-1.
A partir de los valores de luminosidad y temperatura obtenidos, podemos posicionar las
fuentes en el Diagrama de Hertzprung-Russell (HRD) para comparar las masas resultantes
del ajuste espectral (masa espectroscópica) y las de los modelos de evolución estelar (masa
evolutiva). En la Fig. 11.5 vemos que nuestras dos donantes se sitúan en una región del HRD
compatible con la evolución de estrellas de masa inicial de ∼ 25 − 30 M⊙ . En IGR J175442619 la masa espectroscópica es compatible con la masa evolutiva. Vela X-1 muestra cierto
exceso de luminosidad, siendo la masa espectroscópica mayor que la evolutiva. Sin embargo, hacemos notar que la masa espectroscópica representa la masa estelar actual de la
estrella, la cual decrece con el tiempo debido al viento estelar y a posibles episodios de
transferencia de masa. Estos fenómenos han podido ser por tanto más intensos o más prolongados en Vela X-1 que en IGR J17544-2619.
Las abundancias de helio y nitrógeno son claramente mayores que las solares, lo que
puede producir un aumento de luminosidad según la relación de escala L ∝ µα , donde µ
es el peso molecular medio y α > 1 (Langer, 1992). Por lo tanto, esperamos cierta sobreluminosidad en ambas fuentes. Este efecto es más evidente en Vela X-1, como ya hemos
comentado. Teniendo todo ello en cuenta, las fuentes parecen haber tenido una historia
evolutiva diferente o están en una fase evolutiva distinta.
La evolución química de las donantes podría haber sido determinada por episodios de
importante transferencia de masa, dada la poca distancia orbital en los dos sistemas, lo que
conllevaría un incremento de helio y nitrógeno debido a la acreción de material enriquecido (Langer, 2012). Además, fases de desbordamiento del lóbulo de Roche producen la
aceleración angular de la estrella acretante (Packet, 1981), induciendo una mayor mezcla
de material por efecto rotacional. Este escenario esta apoyado por la observación de otras
HMXBs donde también se observan indicaciones de aumento de nitrógeno en la atmósfera
de la donante (González-Galán et al., 2014a).
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Uno de los fenómenos más sorprendentes que hemos encontrado en los espectros son las
asimetrías evidentes en muchas líneas espectrales intensas de Vela X-1 (ver Fig. 11.6). Son
especialmente evidentes en líneas de He i, pero también se observan en C, N, O y Si.
Asimetrías en las líneas espectrales fueron observadas por Martínez-Núñez et al. (2015)
en líneas de hidrógeno del espectro infrarrojo de X1908+75, una SGXB. Una explicación
natural a esta discrepancia entre modelos (esféricamente simétricos) y observaciones es la
desviación del donante y/o el medio circunestelar de la simetría esférica. Esta desviación
puede estar producida por efectos inducidos por fuerzas de marea (Koenigsberger et al.
2012) o por el impacto de los rayos X en el viento (Blondin et al., 1990; Krtička et al.,
2015).
Conclusiones
Las estrellas masivas son un elemento fundamental en la evolución de las galaxias debido al papel principal que juegan en la retroalimentación de material y su enriquecimiento
químico. El medio circunestelar de estas estrellas nos da información sobre muchas de sus
propiedades, por lo que su estudio es de un alto interés científico. En el caso de las HMXBs,
tenemos además la posibilidad de utilizar la fuente de rayos X como una auténtica sonda a
través de ese medio circunestelar, lo que abre vías adicionales de investigación. En este trabajo hemos tratado de aprovechar tanto las posibilidades que nos ofrecen las observaciones
en rayos X como las del infrarrojo, óptico y ultravioleta, para intentar describir y explicar
las propiedades del medio circunestelar de estos objetos.
Primero, hemos realizado un análisis espectral de la muestra completa de observaciones
de HMXBs disponibles con XMM-Newton hasta Agosto de 2013, con el fin de caracterizar
FeKα, la principal línea de emisión en los rayos X. En total, el estudio incluye 46 HMXBs,
21 de las cuales muestran emisión significativa de FeKα, lo que supone el estudio más
completo de la línea de FeKα en HMXBs hecho hasta la fecha. Como se esperaba, hemos
encontrado un grupo muy heterogéneo de objetos y estados de luminosidad, que ha sido
debidamente organizado. Finalmente tenemos un conjunto de 108 espectros, cuyo análisis
ha conducido a las siguientes conclusiones:
• El atlas espectral del complejo del Fe (ver Apéndice A) nos da una descripción cualitativa de los distintos grupos de HMXBs. Especialmente reconocibles son los patrones encontrados en SGXBs (líneas de fluorescencia pero no de recombinación), y
los análogos de γ Cass (modelados con modelos mekal que incluyen líneas de recombinación, con el añadido de líneas de fluorescencia). FeKα es muy probablemente un
rasgo ubicuo de las HMXBs, pero su detección depende de la calidad de las observaciones. SGXBs y SFXTs, que muestran un mayor NH entre las HMXBs, tienden a
exhibir una fluorescencia más prominente.
• Los flujos de energía del continuo y FeKα están directamente correlacionados, como
esperamos de la emisión fluorescente por la iluminación de una fuente de rayos X.
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Los diferentes coeficientes de correlación entre observaciones en eclipse y fuera de
eclipse indican que FeKα se produce en una región extensa que abarca desde las
proximidades de la fuente de rayos X hasta distancias cercanas a la del radio estelar.
• Confirmamos una correlación inversa entre la luminosidad en rayos X y la EW de
FeKα (X-ray Baldwin effect). Los análogos de γ Cass no siguen esta correlación. Este
hecho sugiere que el escenario de formación de la fluorescencia es fundamentalmente
distinto en SGXBs y análogos de γ Cass.
• La anchura de FeKα es predominantemente menor que 0.15 keV y puede ser explicada por procesos de solapamiento de líneas, ensachamiento Compton y desplazamiento Doppler moderado (∼ 1000 km/s).
• La curva de crecimiento en SGXBs muestra una clara correlación entre la EW de
FeKα y NH , indicando un fuerte vínculo entre el material que absorbe la radiación
de rayos X y la que produce la fluorescencia. A partir de simulaciones numéricas
vemos que este material está distribuido de manera aproximadamente isotrópica en la
mayoría de SGXBs.
• El NH en SGXBs es sistemáticamente mayor que en SFXTs, lo que apunta o bien a
una diferente interacción objeto compacto - viento estelar, o bien diferencias en los
parámetros orbitales o bien distintos vientos estelares.
• La modulación orbital de NH en IGR J16320-4751 y 4U 1700-37, junto con los resultados mencionados anteriormente, demuestran que el viento estelar en donantes
de supergigantes contribuye de manera fundamental a la absorpción de rayos X y la
emisión de FeKα.
El estudio de FeKα nos ha dado pie a invetigar en más detalle las características de los
vientos estelares en los dos grupos de HMXBs en los que encontramos una compañera supergigante: las SGXBs y las SFXTs.
En concreto, hemos hecho un análisis espectral detallado de las compañeras ópticas en
Vela X-1 e IGR J17544-2619, dos de los miembros más representativos de las SGXBs y
SFXTs, respectivamente. Para ello hemos utilizado el código Potsdam Wolf-Rayet (PoWR),
originalmente ideado para modelar la atmósfera de estrellas Wolf-Rayet, pero actualmente
aplicable a la de cualquier estrella caliente OB. Este análisis nos ha permitido obtener una
estimación de los siguientes parámetros de las donantes: luminosidad, extinción, masa estelar, radio estelar, temperatura efectiva, gravedad superficial, velocidad terminal del viento,
pérdida de masa, factor de clumping, velocidad de micro y macro-turbulencia, velocidad de
rotación proyectada y abundancias químicas. A partir de estos parámetros, hemos podido
derivar otros igualmente importantes a partir de trabajos anteriores de otros autores: el radio estelar de IGR J17544-2619 implica acotar la excentricidad del sistema a e < 0.25. La
velocidad rotacional de Vela X-1 implica que la masa de la estrella de neutrones puede ser
NS
∼ 1.5 M⊙ , cercana al valor canónico (1.4 M⊙ ).
MVela
X-1
En este estudio hemos visto que los parámetros encontrados en IGR J17544-2619 y
Vela X-1 no son particularmente peculiares, sino que se adaptan bien a lo que esperamos de
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su tipo espectral. Además, las diferencias entre sus propiedades físicas son menores, y sus
parámetros orbitales son igualmente comparables, ya que los dos sistemas presentan una órbita casi circular y bastante cerrada. Sin embargo, en el contexto del marco teórico descrito
por Bozzo et al. (2008), sus moderadas diferencias en el viento estelar, combinadas con el
periodo de rotación de la estrella de neutrones, pueden conducir a regímenes de acreción
muy distintos, lo que cualitativamente explica sus grandes diferencias de comportamiento
en rayos X. Debido a la poca cantidad de estudios de este tipo no podemos asegurar que lo
que encontramos en Vela X-1 e IGR J17544-2619 sea extrapolable al conjunto general de
SGXBs y SFXTs, pero hay razones para pensar que de hecho podría ser el caso. Para ello
serán necesarias más investigaciones del viento estelar de donantes en este tipo de sistemas,
así como el descubrimiento del Pspin en un mayor número de SFXTs.
En resumen, el trabajo realizado a lo largo de esta tesis nos ha permitido descubrir las
propiedades del medio circunestelar de las HMXBs desde distintas perspectivas. En particular, el estudio de FeKα aporta importante información del medio que se encuentra alrededor
del objeto compacto. Esta información puede ser utilizada para comprender mejor el medio
circunestelar de las estrellas masivas, lo que tiene un gran interés en varios campos de la astrofísica. Al mismo tiempo, estudios detallados de las compańeras ópticas usando modelos
sofisticados como PoWR nos dan un conocimiento más profundo de las condiciones físicas
presentes en las atmósferas y vientos estelares de estas estrellas. Estos estudios indudablemente ayudan a interpretar más correctamente el comportamiento que observamos en los
rayos X y por lo tanto mejora nuestro entendimiento del conjunto de HMXBs.
Contents
Resumen en Lengua Oficial Comunidad Valenciana
v
I Introduction
3
1 High-mass X-ray Binaries
1.1 Classification . . . . . . . . . . . . . . . . . . . . . . . . . .
1.1.1 The Be X-ray Binaries . . . . . . . . . . . . . . . . .
1.1.2 The γ Cassiopeae analogs . . . . . . . . . . . . . . .
1.1.3 The Classical and Persistent Supergiant X-ray Binaries
1.1.4 The Supergiant Fast X-ray Transients . . . . . . . . .
1.1.5 High-mass γ-ray Binaries . . . . . . . . . . . . . . .
1.1.6 Other sources . . . . . . . . . . . . . . . . . . . . . .
1.2 The circumstellar environment of massive stars . . . . . . . .
1.2.1 Be stars . . . . . . . . . . . . . . . . . . . . . . . . .
1.2.2 OB supergiants . . . . . . . . . . . . . . . . . . . . .
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2 Acretion as a source of energy
2.1 Accretion onto neutron stars . . . . . . . . . .
2.1.1 Low luminosity and collisionless shock
2.1.2 Low luminosity and absence of shocks
2.1.3 High luminosity . . . . . . . . . . . .
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II An XMM-Newton view of FeKα in High-mass X-ray Binaries
21
3
The Fe complex with XMM-Newton
3.1 The Fe complex . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2 The XMM-Newton observatory . . . . . . . . . . . . . . . . . . . . . . .
25
25
26
4
The sources and data treatment
29
xxiii
xxiv
4.1
4.2
4.3
Table of Contents
The sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Data reduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Spectral fitting . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5 Results
5.1 Spectral atlas . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2 FeKα width . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3 Centroid energy . . . . . . . . . . . . . . . . . . . . . . . . .
5.4 Correlated parameters . . . . . . . . . . . . . . . . . . . . . .
5.4.1 Continuum flux vs FeKα Flux . . . . . . . . . . . . .
5.4.2 FeKα Width vs centroid energy . . . . . . . . . . . .
5.4.3 Curve of growth . . . . . . . . . . . . . . . . . . . .
5.5 NH : SGXBs and SFXTs . . . . . . . . . . . . . . . . . . . .
5.6 Individual sources analysis:IGR J16320-4751 and 4U 1700-37
5.6.1 IGR J16320-4751 . . . . . . . . . . . . . . . . . . . .
5.6.2 4U 1700-37 . . . . . . . . . . . . . . . . . . . . . . .
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31
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6 Discussion
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7 Summary
59
III A comparative analysis: Vela X-1 and IGR J17544-2619
61
8 The observations
65
9 The PoWR code
69
10 The fitting procedure
10.1 IGR J17544-2619 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
10.2 Vela X-1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
71
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11 Discussion
11.1 Wind-fed accretion . . . . . . . . . . . . . . . . . .
11.1.1 Outside the magnetospheric radius: RM > Ra
11.1.2 Inside the magnetospheric radius: RM < Ra .
11.1.3 Transitions across regimes . . . . . . . . . .
11.1.4 IGR J17544-2619 . . . . . . . . . . . . . . .
11.2 Evolutionary tracks . . . . . . . . . . . . . . . . . .
11.3 Asymmetries in spectral lines of Vela X-1 . . . . . .
12 Summary
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93
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. 94
. 95
. 96
. 96
. 101
. 102
107
Table of Contents
xxv
IV Conclusions
109
A The fits of the FeKα survey
113
B The fits of the supergiant donors
135
B.1 IGR J17544-2619 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 136
B.2 Vela X-1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 139
Bibliography
143
List of Figures
1.1
1.2
1.3
The BeXBs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
The HMXBs with supergiant donors . . . . . . . . . . . . . . . . . . . . .
The principle of radiatively driven stellar winds . . . . . . . . . . . . . . .
6
8
13
2.1
2.2
2.3
2.4
2.5
Graphic representation of a pulsar . . . . . . . . . . . . . . . . . . . .
Accretion column at the state of low luminosity and collisionless shock
Accretion column at the state of low luminosity and absence of shocks .
Accretion column at high luminosity state . . . . . . . . . . . . . . . .
Comparison observations-theory in X Persei . . . . . . . . . . . . . . .
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15
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19
20
3.1
3.2
Production of FeKα in a neutral iron atom . . . . . . . . . . . . . . . . . .
Inner structure of the XMM-Newton observatory . . . . . . . . . . . . . .
26
27
4.1
4.2
Selection of time intervals . . . . . . . . . . . . . . . . . . . . . . . . . .
Number of accepted models depending on the reduced-χ2 value. . . . . . .
30
32
5.1
5.2
5.3
5.4
5.5
5.6
5.7
5.8
5.9
5.10
5.11
5.12
Patterns found in the Fe complex of HMXBs . .
Histogram of the FeKα width . . . . . . . . . .
Centroid energy of FeKα . . . . . . . . . . . .
F1−10keV versus F FeKα . . . . . . . . . . . . .
EW of FeKα against L1−10 keV . . . . . . . . .
Width of FeKα versus the centroid energy . . .
Curve of growth for FeKα . . . . . . . . . . .
NH : SGXBs versus SFXTs . . . . . . . . . . .
Curve of growth in IGR J16320-4751 . . . . .
Orbital modulation of NH in IGR J16320-4751
Curve of growth in 4U 1700-37 . . . . . . . . .
Orbital modulation of NH in 4U 1700-37 . . . .
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6.1
6.2
NH against the width of FeKα . . . . . . . . . . . . . . . . . . . . . . . .
Sketch of configurations of circumstellar matter in HMXBs . . . . . . . . .
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xxvii
xxviii
List of Figures
10.1 Estimation of T ⋆ in IGR J17544-2619 . . . . . . . . . . . . .
10.2 Surface gravity estimation in IGR J17544-2619 . . . . . . . .
10.3 Extinction curve in the galactic direction of IGR J17544-2619
10.4 Fit of the SED of IGR J17544-2619 . . . . . . . . . . . . . .
10.5 Estimation of 3∞ √in IGR J17544-2619 . . . . . . . . . . . . .
10.6 Estimation of Ṁ D in IGR J17544-2619 . . . . . . . . . . .
10.7 Hα in IGR J17544-2619 at different orbital phases . . . . . . .
10.8 Estimation of Ṁ in IGR J17544-2619 . . . . . . . . . . . . .
10.9 Estimation of ξph in IGR J17544-2619 . . . . . . . . . . . . .
10.10Fit of the SED of Vela X-1 . . . . . . . . . . . . . . . . . . .
10.11Estimation of 3∞ in Vela X-1 . . . . . . . . . . . . . . . . . .
10.12Si iv and C iv resonance
lines in Vela X-1 . . . . . . . . . . .
√
10.13Estimation of Ṁ D in Vela X-1 . . . . . . . . . . . . . . . .
10.14 Ṁ estimation in Vela X-1 . . . . . . . . . . . . . . . . . . . .
10.15Desaturation of Si iv and N v resonance lines . . . . . . . . . .
10.16Estimation of the parameter β in Vela X-1 . . . . . . . . . . .
10.17Rotational velocity in Vela X-1 . . . . . . . . . . . . . . . . .
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11.1
11.2
11.3
11.4
11.5
11.6
11.7
Position of IGR J17544-2619 in the 3rel - Ṁ diagram. .
Position of IGR J17544-2619 in the 3rel -Pspin . . . .
Position of Vela X-1 in the 3rel - Ṁ diagram. . . . . .
Position of Vela X-1 in the 3rel -Pspin diagram. . . . .
Evolutionary tracks from the Geneva Stellar Models.
Spectral lines showing notable asymmetries . . . . .
Blue absorption in important lines. . . . . . . . . . .
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98
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103
105
A.1
A.2
A.3
A.4
A.5
A.6
A.7
BeXBs fits . . . . . .
SGXBs fits . . . . .
SFXTs fits . . . . . .
γ Cass-analogs fits .
HMGBs fits . . . . .
Peculiar sources fits .
AX J1749.1-2733 fit
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114
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120
121
122
B.1
B.1
B.2
B.3
B.4
B.5
Optical spectrum of IGR J17544-2619 (blue), and the best fit model (red).
continued. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Infrared spectrum of IGR J17544-2619 (blue), and the best fit model (red).
Ultraviolet spectrum of Vela X-1 . . . . . . . . . . . . . . . . . . . . . .
Optical spectrum of Vela X-1 (blue), and the best fit model (red). . . . .
continued. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
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138
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. 141
List of Tables
1.1
Features in the X-rays of γ Cass, OB stars and BeXBs . . . . . . . . . . . .
7
4.1
4.2
Table of sources included in the sample of HMXBs . . . . . . . . . . . . .
List of models used to fit the continuum . . . . . . . . . . . . . . . . . . .
34
35
5.1
Description of features of the HMXBs atlas . . . . . . . . . . . . . . . . .
39
8.1
8.2
Table of observations of IGR J17544-2619 . . . . . . . . . . . . . . . . . .
Table of observations of Vela X-1 . . . . . . . . . . . . . . . . . . . . . .
66
67
10.1 Stellar parameters obtained from the best fit . . . . . . . . . . . . . . . . .
10.2 Chemical abundances derived from the best fit . . . . . . . . . . . . . . . .
72
73
11.1 Parameters used in Chapter 11. . . . . . . . . . . . . . . . . . . . . . . . .
93
A.1 Parameters of the continuum . . . . . . . . . . . . . . . . . . . . . . . . . 123
A.2 FeKα parameters . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 130
xxix
Abstract
X-ray Binary systems are relevant for the contemporary Physics thanks to their unique
conditions as astrophysical laboratories, which are far beyond the capabilities of any conceivable laboratory on our planet. In these systems we can witness huge accretion processes
transforming energy more efficiently than nuclear reactions. They permanently test our current knowledge on magnetohydrodynamics and radiative processes. They involve the study
of nuclear physics from the study of the properties of the neutron stars. Moreover, the
X-ray emission is an outstanding scan of the interstellar medium and supplies invaluable
information about any structure that interacts with it: stellar winds, circumstellar disks, etc.
It is beyond any doubt that the astrophysical properties of these systems are scientifically
appealing from many different perspectives.
Furthermore, massive stars play a crucial role in our understanding of the past, present
and future of our own galaxy, and probably of any other galaxy too. They are actual cosmic
engines triggering stellar formation due to their particular way of existing and capitulating:
they live fast and explode as supernovae. The interstellar medium is consequently enriched
with new chemical elements. This enrichment, after several generations of stars, enables
the eventual formation of rocky planets such as the Earth, which are generally assumed as
an essential block for the development of life throughout the Universe.
When massive stars are confined in a binary system with a compact object (a black hole,
a neutron star or a white dwarf) and the conditions are propitious, abundant amounts of
X-rays are produced in the very surroundings of the compact star. This highly energetic
radiation conveys information about the medium where it has been formed and has been
propagated (the circumstellar medium). On the other hand, the emission of the systems in
less energetic wavelength domains, from the infrared to the ultraviolet, is dominated by the
direct radiation of the massive star. This light is the carrier of information about its photospheric and stellar wind properties.
This thesis is a multi-wavelength analysis of X-ray Binaries hosting a massive star. It
is structured in two main blocks. In the first one we use the X-rays in order to study the
circumstellar medium of a large set of systems. In the second block, we use infrared, optical
and ultraviolet observations in order to characterise the physical properties of the donors in
two relevant members of High-mass X-ray Binaries. The combination of both analyses
gives a comprehensive view and sheds light in the physical interpretations of this important
class of astrophysical objects.
1
Part I
Introduction
In the first part of the thesis we present the topic of this work: the High-mass X-ray Binaries and its circumstellar medium. We describe these objects and expose their principal
observational properties. Furthermore, we expand on the standard classification that arises
from these observational features and introduce the theoretical frameworks that have been
put forward in order to explain the different properties of the systems. Then, we outline
the fundamental properties of the circumstellar medium in the two main classes of donors
that are present in High-mass X-ray Binaries: the OB supergiants and the Be stars. We
also explain how the gravitational energy of the matter in the circumstellar medium is transformed into radiation by means of accretion processes. Finally, we describe three possible
configurations of accretion column on magnetic neutron stars, given that it is the most common scenario. In all, we sketch the current understanding of High-mass X-ray Binaries and
highlight the fundamental physical aspects that will be useful throughout our research.
Chapter 1
High-mass X-ray Binaries
An X-ray Binary system (XRB) consists of a compact object (a white dwarf, neutron
star o black hole) accreting matter from a companion star. When this companion is more
massive than ∼ 10 M⊙ , the system is called a High-mass X-ray Binary (HMXB). The gravitational energy from the accreted material is eventually transformed into X-ray radiation
in a very efficient fashion. The donor star is usually called optical or normal companion,
provided that their luminosity is still powered by the nuclear reactions occurring in their
interior, which generally implies a high luminosity in the optical domain. In the catalogue
of Liu et al. (2006) there are 114 candidates of HMXB in the Milky Way, although some
of them are not confirmed so far. Nowadays, the number of confirmed HMXBs is everincreasing and the most recent review by Walter et al. (2015) included 87 of them.
1.1 Classification
High-mass X-ray Binaries were traditionally classified according to the luminosity class
of the donor within two main groups: the systems containing a supergiant star powered by
the stellar wind or Roche overflow of the donor; and the systems hosting a Be star with a
circumstellar disk providing the energy budget for the X-rays production. However, in the
last decade and a half new discoveries have led to a more complex picture of the HMXBs
zoo. New groups of sources have emerged with their own peculiarities, stressing the necessity of further study from several perspectives: the geometry of the system, properties of the
compact object and optical star, the stellar wind, etc. Next, we expose the different classes
of HMXBs and their most important features.
1.1.1 The Be X-ray Binaries
1
Even though they were initially tagged as atypical cases of HMXBs, nowadays Be
X-ray Binaries (BeXBs) are the most numerous group among the HMXBs. The optical
companion in BeXBs is a Be star of luminosity class III-V. The characteristic e stands for
the detection of hydrogen emission lines at some stage of their evolution, which indicate the
1
References: Ziolkowski (2002), Reig (2011).
5
6
Chapter 1. High-mass X-ray Binaries
presence of an equatorial disk around the donor star.
Figure 1.1: Graphic representation of a BeXB. The mass-loss from the Be star is mainly
c
driven by the equatorial disk, from where the compact object accretes ESA.
Every known BeXB harbours a neutron star, except the recently discovered black hole
orbiting HD 215227 (Casares et al., 2014). When the compact object is close to the periastron and passes near the circumstellar disk, the latter is disrupted and the mass transfer is
enhanced (Figure 1.1). The observational result is an increase in luminosity up to ∼ 1037
erg/s and a non-thermal spectrum (Reig, 2011).
BeXBs are generally characterized by the transient nature of the enhanced luminosity
episodes. These outbursts can be divided in two types. Type I outbursts are short, regular
and periodic (or quasi-periodic). They normally peak about the periastron passage. Type II
outbursts are longer and more intense, and they do not show any obvious preference with
regard to the orbital phase. In addition, they are usually preceded by an intensification of
the emission lines in the Be spectrum. According to Martin et al. (2014), Type II outbursts
can be explained as a result of a growth of eccentricity in the equatorial disk. This growth
arises from the interaction with the compact object describing an orbit with an inclination
misaligned with the equatorial disk. After several orbital periods the compact object is able
to capture a large amount of material forming an accretion disk and producing a giant flare.
The majority of BeXBs are transient, but there are also persistent sources (Reig and Roche,
1.1. Classification
7
1999). They are characterised by the absence of large outbursts, a luminosity LX . 1035 erg/s,
high spin period of the neutron star Pspin > 200 s and long orbital period Porb > 200 d (Reig,
2011).
1.1.2 The γ Cassiopeae analogs
2
γ Cassiopeae is a Be star discovered by Pietro Angelo Sechi in 1867. Indeed, it is the
first discovery of a Be star. It is part of a binary system with an orbital period of 204 days
and low eccentricity. The secondary component is mostly unknown, apart from having a
mass ∼ 1 M⊙ . When observed in the infrared, optic and ultraviolet wavelengths, it seems
a normal Be. However, its X-rays behavior is far from being coherent with the depicted
picture of BeXBs. The high-energy emission is characterised by thermal emission with
temperature T ∼ 1.4 × 108 K (∼ 12 keV), responsible of ∼ 80 − 90% of the total flux. The
luminosity (∼ 1032−33 erg/s) is variable considering any timescale (from several seconds to
minutes), but, in contrast with the BeXBs, it does not show large outbursts.
γ Cass
Emission
Temperature (keV)
Variability
Luminosity (erg/s)
Thermal
∼ 12
High
∼ 1032−33
isolated OBs
Thermal
∼ 0.5
Low
. 1033
BeXB
persistents
transients
Non-thermal
High
∼ 1033
Non-thermal
High
∼ 1036−37
Table 1.1: Main observational features in the X-rays of γ Cass, isolated OB
stars and BeXBs.
The discovery, in the last years, of new sources showing properties similar to γ Cass, motivated the definition of a new class of Be stars with X-ray emission: the γ Cassiopeae like
sources (or γ Cassiopeae analogs). They show the following features:
• Thermal emission with characteristic temperature > 5 keV.
• Variability within any timescale.
• X-rays luminosity (0.2 − 12 keV) ∼ 1033 erg/s.
Nowadays, the origin of the X-rays in these systems is not clear. Two major possibilities
have been put forward. The first one is that the X-rays are powered via accretion of matter
onto a white dwarf or a neutron star. The second one appeals to a magneto-rotational instability triggered by the interaction of the stellar surface and the magnetic field of the disk.
Any of those explanations, if confirmed, would have important astrophysical implications.
The former would imply the discovery of the first white dwarf in a HMXB, as expected from
evolutionary models of binaries, or the presence of a neutron star in an accretion regime
2
References: Lopes de Oliveira et al. (2010).
8
Chapter 1. High-mass X-ray Binaries
hitherto unknown. The latter would imply that the magnetic field is very significant and is
able to produce magneto-rotational instabilities (Balbus and Hawley, 1991), indicating the
plausibility of γ Cassiopeae analogs as magnetar progenitors.
1.1.3 The Classical and Persistent Supergiant X-ray Binaries
3
Among the sources with a supergiant optical star, we can separate the persistent sources
showing a moderate dynamic range (ratio between peak and bottom luminosities) from the
transient sources (described below in the next subsection). Prior to the launch of INTEGRAL, persistent sources were the only known variety of HMXBs with supergiant companion, and therefore are commonly called classical. Henceforth, we use the acronym
SGXB when referring to the classical and persistent Supergiant X-ray Binaries.
In SGXBs, the compact object is embedded in the powerful stellar wind of a OB supergiant companion, swallowing everything that enters its gravitational domain. The compact
object is usually found at a distance of ∼ 1.5−2 R⋆ . In such a close orbit, the captured matter
is able to fuel persistent X-ray emission up to ∼ 1036−37 erg/s. Flares and off-states are often observed in SGXBs, indicating an abrupt transition in the accretion rate. They might be
produced either by sudden variations in density within the medium transited by the compact
object (Kreykenbohm et al., 2008; Martínez-Núñez et al., 2014) or by instabilities above the
magnetosphere of the neutron star, as proposed in the quasi-spherical accretion theory by
Shakura et al. (2012a).
Figure 1.2: Graphic representation of a HMXB with a supergiant companion. The compact
object orbits a dense region that provides the energy budget for the X-ray emission.
1.1.4 The Supergiant Fast X-ray Transients
4
The number of SGXBs known before the launch of INTEGRAL was around a dozen.
The scarcity of sources was generally imputed to the very short duration of the supergiant
phase in the lifetime of a massive star and the consequent serendipity required to catch
3
4
References: Walter et al. (2015).
References: Negueruela et al. (2008a), Chaty (2011), Sidoli (2013), Walter et al. (2015).
1.1. Classification
9
the sight of such an ephemeral stage. However, INTEGRAL tripled the number of known
HMXBs with supergiant companion in the Milky Way and revealed the existence of new
types of systems which could not be assimilated as classical systems. They were the Supergiant Fast X-ray Transients (SFXT).
Supergiant Fast X-ray Transients are systems with a supergiant optical star, as in SGXBs,
but they are defined by extremely transient behaviour. The dynamic range in SGXBs is . 2
orders of magnitude. In contrast, the dynamic range in SFXTs can reach up to five orders of
magnitude in the most extreme cases (e.g. IGR J17544-2619). During quiescence, SFXTs
exhibit a low X-ray luminosity (LX ∼ 1032 erg/s), but they spend most of their time in an
intermediate level of emission (∼ 1033−34 erg/s). They display short outbursts (∼few hours),
reaching luminosities up to 1036−37 erg/s (Sidoli, 2011; Sidoli et al., 2009). We note that
there are also "intermediate SFXTs" with a dynamic range of & 2 orders of magnitude,
somewhere between SGXBs and SFXTs. Hence, there is no sharp border that clearly separates SGXBs and SFXTs.
The dichotomy between SGXBs and SFXTs
Negueruela et al. (2008b) suggested that the intrinsic clumpiness of the wind of hot supergiant donors, together with different orbital configurations, may explain the different
dynamic ranges between SGXBs and SFXTs. If the eccentricity of SFXTs is high enough,
the compact object swings between dense regions with a high probability of accreting a
wind clump and flare up, and diffuse regions where this probability is low and the source is
consequently faint in the X-rays. In SGXBs, the compact object would orbit in a closer and
more circular trajectory, accreting matter incessantly. However, the short orbital period of
some SFXTs is contradictory with this scheme (Walter et al., 2015).
Other ingredients, such as the magnetic field and/or the spin period of the neutron star,
might be important. This is supported by the monitoring of SFXTs. Tracing SFXTs for a
long period, Lutovinov et al. (2013) conclude that, in SFXTs, the accretion is notably inhibited most of the time. Bozzo et al. (2008) invoked the different possible configurations of
accretion, co-rotation and magnetospheric radius in order to relax the extremely sharp density contrast required for the interpretations discussed above. These configurations depend
on the wind, orbital, and neutron star parameters. For instance, if the magnetospheric radius
is larger than the accretion radius (Bondi, 1952), the inflow of matter is significantly inhibited by a magnetic barrier, resulting in a relatively low X-ray emission from the source. In
the interpretation of Bozzo et al. (2008), the physical conditions in SFXTs make them prone
to regime transitions as a response to relatively modest variations in the wind properties of
the donor, which cause abrupt changes in X-ray luminosity.
The theory of quasi-spherical accretion onto slowly rotating magnetized neutron stars developed by Shakura et al. (2012b) describes the so-called subsonic settling accretion regime
in more detail. In slowly rotating neutron stars, the penetration of matter into the magnetosphere is driven predominantly by Rayleigh-Taylor instabilities (Elsner and Lamb, 1976).
10
Chapter 1. High-mass X-ray Binaries
When the cooling of the plasma in the boundary of the magnetosphere is not sufficiently
efficient, the accretion of matter is highly inhibited and consequently the X-ray luminosity
is low. On the other hand, when the cooling time is much smaller than the characteristic
free-fall time (tcool ≪ tff ), the instability conditions are fulfilled and the plasma easily enters
the magnetosphere, triggering high X-ray luminosity. The last is achieved when the X-ray
luminosity is LX & 4·1036 erg/s, and the rapid Compton cooling dominates over the radiative
cooling. For the brightest flares (LX > 1036 ), Shakura et al. (2014) proposed that a magnetized wind of the donor might induce magnetic reconnection, enhancing the accretion up to
the critical X-ray luminosity and triggering the suction of the whole shell by the neutron star.
1.1.5 High-mass γ-ray Binaries
5
High-mass γ-ray binary systems (HMGBs) are HMXBs where the emission peaks
above 1 MeV. Nowadays, it is thought that the emission is caused by accelerated particles in
the shock that is produced when the pulsar wind collides the massive star wind. Therefore,
they are powered by the rotational energy of the neutron star, in opposition to the rest of
HMXBs, which are accretion fed. There are currently five confirmed HMGBs, all of them
with a main sequence optical star.
1.1.6 Other sources
Finally, there are a few sources that, for a number of reasons, cannot be classified in any
of the already mentioned classes of HMXBs. Among the set of sources studied in this thesis,
they are 4U 2206+54, Centaurus X-3, and Cygnus X-1. The optical star in 4U 2206+54 is
a O9.5V (Blay et al., 2006), which is neither a supergiant nor a Be star. The system may be
part of a new group of wind-fed HMXBs with a main sequence donor (Ribó et al., 2006).
The donors in Centaurus X-3 and Cygnus X-1 are supergiants, but the accretion is driven by
an accretion disk (Tjemkes et al., 1986; Shapiro et al., 1976). This fact completely changes
their behavior in the X-rays compared to other classes of HMXBs. In addition, the compact
object in Cygnus X-1 is a black hole, which adds further peculiarities to the spectrum of the
source.
1.2 The circumstellar environment of massive stars
We refer to the circumstellar environment of massive stars as the stellar medium immediately surrounding the object at a distance r . 2 R⋆ . This is the region that makes the highest
impact on the X-ray behavior of HMXBs.
As already mentioned, the optical companion in HMXBs are, with very few exceptions,
OB supergiants or Be III-V stars. Next, we describe the properties of the circumstellar
medium in these type of objects.
5
References: Dubus (2013).
1.2. The circumstellar environment of massive stars
11
1.2.1 Be stars
6
Be stars are not uncommon. They are the 15-20% of B stars in the local environment.
Moreover, in the earlier universe the proportion of Be’s must be higher, given that lower
metallicity implies weaker radiatively driven stellar wind and less loss of angular momentum.
The main property of the circumstellar medium in Be stars is obviously the disk, which
becomes detectable thanks to the presence of Balmer emission lines in the spectrum. It is
usually known as decretion disk because the mass is transported outwards, in contrast to the
accretion disks, where the mass flows inwards. The disks are Keplerian, as deduced from
the observed very small radial velocities and azimuthal velocity 3φ ∝ r−1/2 in the reference
frame of the star (Rivinius et al., 1999; Meilland et al., 2012).
The geometry and structure of the disk has been studied in the past by many authors
using interferometry and spectropolarimetry. Quirrenbach et al. (1997) analysed seven Be
stars and ruled out envelope models that are both optically and geometrically thick. Instead,
these authors suggested that the disks are geometrically thin, and constrained the vertical
extent of the disk to a half-opening angle . 20◦ in ζ Tau. The estimation of the size of the
disk is highly dependent on the employed model. Moreover, observations can only give insight on the size of the emitting region of a certain spectral line, e.g. Hα. Quirrenbach et al.
(1997) estimated the size of Hα emitting regions in their sample to be 3 − 12 R⋆ . Regarding the density structure of the disk, models assuming a gaussian vertical profile and a
power-law radial profile are able to successfully reproduce observational properties of large
samples of Be stars (Silaj et al., 2010; Touhami et al., 2011).
The formation mechanism of the disk is not fully understood so far. It is very likely
related to the fact that these stars are systematically fast rotators, but it is necessary to appeal to further mechanisms in order to explain the properties of a Keplerian decretion disk,
given that it would dissipate unless it is constantly provided with angular momentum from
the inner boundary. Some of the mechanisms hitherto proposed for the supply of mass and
angular momentum are non-radial pulsations, explosive ejections and Wind-Compressed
Disk models (Bjorkman, 2000). These mechanisms are also susceptible to be enhanced by
additional phenomena such as binary tidal torques or magnetic fields. Nevertheless, none
of them, considered isolated, are able to successfully explain all the observable features in
the Be’s disks.
Apart from the disk, the circumstellar medium is partially shaped by the stellar wind.
Ultraviolet observations confirm that there are important signatures of stellar winds in Be
stars in the form of additional discrete absorption components (Grady et al., 1987). A bimodal behavior was observed depending on the projected rotational velocity 3 sin i: low
3 sin i objects (≤ 150 km/s) show similar winds to non Be stars; intermediate to high 3 sin i
sources exhibit stronger winds when they are in Be phases. In any case, the stellar wind
in Be stars is not the leading component driving the mass loss of the star, since this role is
6
References: (Rivinius et al., 2013).
12
Chapter 1. High-mass X-ray Binaries
mainly played by the decretion disk.
1.2.2 OB supergiants
7
In massive supergiants, the main phenomenon that shapes the circumstellar medium is
the stellar wind. These stars lose mass at a high rate Ṁ & 10−7 M⊙ yr−1 , accelerating the
wind up to terminal velocities 3∞ ≃ 500 − 3000 km/s. The basic mechanism for such a
high acceleration is the absorption of radiation by the atmospheric atoms. In every transition, the momentum of the absorbed photon is transmitted to the absorber. Given that
the radiation field points outwards, the net result is the acceleration of particles outwards
(Figure 1.3). Moreover, Coulomb interactions drag additional material and enhance the
mechanism. Eventually, the gravitational potential of the star is overcome and the stellar
wind is formed (Lucy and Solomon, 1970; Castor et al., 1975b).
A particular property of line-driven stellar winds is that they are unstable to velocity
perturbations. As already pointed out by (Lucy and Solomon, 1970), any increase in the
velocity of a fluid element tends to grow and consolidate. This is due to the exposure of
such a fluid element to a less absorbed radiation field in the specific wavelengths where
it absorbs, simply because the surrounding gas moves with a different velocity and therefore is differently Doppler-shifted with respect to the star. Shocks are formed as a natural
consequence, and the expected velocity field in the framework of an homogeneous wind is
disrupted. The shocks heat the atmosphere and are detectable from its X-ray emission and
the superionization of some elements like nitrogen and oxygen. In addition, the stellar wind
becomes clumped, namely, it has regions where the medium is denser than expected in an
homogeneous wind. This fact modifies the optical depth of recombination lines that are important for the mass-loss rate diagnostics, introducing an important source of uncertainty.
Therefore, the consideration of clumps is crucial for the characterization of the stellar winds
in hot supergiants.
Hence, in HMXBs with a supergiant donor, the medium transited by the compact object is not smooth because velocity and density inhomogeneities (clumps) are present as
an intrinsic feature of the radiatively driven winds of hot stars. On top of that, the action
of the compact object may induce further inhomogeneities. Hydrodynamical simulations
show that the X-ray radiation and the gravity field of the compact object disturb the wind of
the donor, inducing the formation of denser structures such as filaments, bow shocks, and
wakes (Blondin et al., 1990, 1991).
Considering the complexity of the circumstellar medium in HMXBs with a supergiant
companion, it seems natural to observe such a rich gamut of observational features in these
type of systems. Moreover, the compact object properties might vary from one system to
another. All these properties extend the scientific relevance of the systems.
7
References: Lucy and Solomon (1970), Castor et al. (1975b), Puls et al. (2008).
1.2. The circumstellar environment of massive stars
13
Figure 1.3:
The principle of radiatively driven stellar winds.
The atoms in
the atmosphere absorb radiation and the momentum carried by the photons. The
c Image from
net result is an acceleration outwards that boosts the wind.
http://www.ifa.hawaii.edu/users/kud/windsfromhotstars/hotwinds.html
Chapter 2
Acretion as a source of energy
1
The X-ray emission in HMXBs is produced as a consequence of an accretion process,
consisting on the addition of certain amount of matter, initially located in the neighborhood
of the compact object, onto its surface. The gravitational energy is first transformed into
kinetic energy when the material is accelerated towards the compact object. Then, the
kinetic energy is transformed into thermal energy by the interaction with additional accreted
material or the surface of the compact object. Finally, this energy is radiated.
c
Figure 2.1: Graphic representation of magnetic accretion onto a neutron star. NASA
The amount of energy released in the accretion processes strongly depends on the mass
and radius of the object that accretes the matter. Let us illustrate it as simple as possible: we
assume an spherical body of mass M and radius R, and a particle of mass m initially at rest
1
References: Frank et al. (2002) and lectures notes from the course Procesos de acreción, given in the
Master’s degree of Astrophysics at the University of La Laguna by Dr. Ignacio González Martínez-País.
15
16
Chapter 2. Acretion as a source of energy
falling from the infinite, until it finally settles down onto the surface of the spherical object.
Then, the particle experiences a change in its mechanical energy:
∆E =
Z
R
∞
~ =
F~ · dr
Z
R
∞
GMm
GMm
dr = −
= −∆Eacc
2
R
r
(2.1)
That is to say, the particle loses mechanical energy, and we label this loss as ∆Eacc . In
order to see the significance of ∆Eacc , we compare it with the energy that is generated from
the same mass m in nuclear reactions. The maximum is obtained if the material is hydrogen
and fuse into helium:
∆Enuc = 0.007mc2
where c is the speed of light. Hence,
M R⊙
∆Eacc
GM
=
= 3.03 10−4
2
∆Enuc 0.007c R
M⊙ R
As we can see, the most favorable situation for the accretion energy to become important
is R ≪ R⊙ , and/or M ≫ M⊙ . This is the case of HMXBs, where the accreting object is a
neutron star or a black hole.
In case we have, instead of an individual particle, a constant flux of matter at rate Ṁ
flowing towards the spherical object, and we assume that all the mechanical energy is finally
radiated, the resultant luminosity is known as the accretion luminosity and is defined as:
Lacc =
GM Ṁ
R
(2.2)
2.1 Accretion onto neutron stars
By far, the most common type of compact object in HMXBs is the neutron star. Next,
we describe the basics of accretion onto its surface.
Among its extreme physical properties, the intense magnetic field is one of the more crucial and characteristic of neutron stars. Typically, the surface magnetic field is as high as
B & 1012 G. This fact strongly affects the way the matter is accreted: the ionized particles
are trapped by the field lines when they hit the magnetosphere and they are channeled towards the poles. These particles form what is known as accretion columns (Fig. 2.1).
The physics of accretion columns involves relativistic free-fall velocities 3ff ∼ c/2, extreme magnetic fields and the action of strong radiation fields whose pressure may overcome the gravitational force at high accretion rates. Given the difficulties those physical
conditions entail, there is not a comprehensive and self-consistent description in the literature enclosing the entire interval of observed X-ray luminosities. Nevertheless, the issue
has been addressed by several authors who found solutions within certain X-ray luminosity
restrictions, as summarised below.
2.1. Accretion onto neutron stars
17
2.1.1 Low luminosity and collisionless shock
In the accretion columns the gas falls at supersonic speed. Assuming that the gas is decelerated to subsonic velocities before settling down, some kind of shock must be present
somewhere inside the accretion column. A shock dominated by collisions is not possible in
this case because the mean free path of the particles is larger than the stellar radius. Therefore, it must be of other type, also known as a collisionless shock (Stockem Novo et al.,
2015). This kind of shocks are also found in very different scenarios like the region where
the stellar wind of the Sun hits the Earth’s magnetosphere.
Under the assumption that no energy is transferred from the ions to the electrons within
the shock, the post-shock temperature of the ions is much higher than the electrons due
to its much higher kinetic energy and the shock jump conditions (Langer and Rappaport,
1982). Then, the ions heat the electrons in order to reach the thermal equilibrium, and the
electrons release energy by means of free-free, Inverse Compton and cyclotron radiation.
The emerging spectrum is dominated by cyclotron emission lines at the local cyclotron
frequency for every height of the column (see Figure 2.2).
Figure 2.2: Left panel: Geometry of the accretion column on the pole of a neutron star
at the state of low luminosity and collisionless shock. Right panel: Emergent spectrum,
with energy units in cyclotron photons on the surface of the neutron star (h̄w = hνcyc ).
(Langer and Rappaport, 1982).
2.1.2 Low luminosity and absence of shocks
In this case, the accretion column is heated by Coulomb collisions and compressed by
the ram pressure (ρ32 ) of the accreted matter. The two main cooling mechanisms are freefree and Inverse Compton. Meszaros et al. (1983) calculated the physical structure of an
18
Chapter 2. Acretion as a source of energy
accretion column under these hypothesis. These authors found that the emerging spectrum
is similar to a power-law energy profile (KE −α ) as shown in the left panel of Figure 2.3. For
comparison, in the right panel of the same figure we show the profile of a power-law model
modified by photoelectric absorption. That is to say, a phenomenological model such as
a simple power-law can be physically identified as Inverse Compton emission. We stress
this point because this type of model has been frequently employed throughout this work in
order to characterise and study many HMXBs’ spectra.
Figure 2.3: Left panel: emergent spectrum computed from the model of an accretion
column heated by Coulomb interactions (Meszaros et al., 1983). Right panel: Absorbed
power-law energy profile. Even though this function is usually adopted as a phenomenological model in order to describe the spectrum of many HMXBs, it can be physically assimilated to the transfer of energy from very hot electrons to the radiation field via Inverse
Compton processes.
2.1.3 High luminosity
The previous models do not consider the radiation pressure and consequently they cannot be applied to the most luminous X-ray sources. According to Davidson (1973), the
radiation pressure decelerates the material producing a very steep transition of the velocity
and density of the accreting flux commonly called radiative shock. It leads to the formation
of a very dense mound in the base of the column where the accreted matter settles down
(Figure 2.4, left panel).
Becker and Wolff (2007) proposed a model of an accretion column with a radiative shock
where the emergent spectrum is the result of the Inverse Comptonization of seed photons
formerly emitted throughout the column via black-body, free-free and cyclotron emission
(Figure 2.4, right panel). The final outcome of energy has a power-law distribution with
an exponential cut-off at high energies. The black-body radiation comes from the mound in
the base of the column, where thermodynamic equilibrium prevails. Cyclotron and free-free
2.1. Accretion onto neutron stars
19
Figure 2.4: Left panel: base of the accretion column at high luminosity state. The plot
2
shows the constant velocity profiles in terms of Q = 332 for the cases (a) Q = 0.9, (b)
ff
Q = 0.5, (c) Q = 0.1. Horizontal and vertical axes are scaled with variables κSc r y κSc z ,
where κ = σT /mp , S = ρ|3|, and r, z are the usual cylindrical coordinates (Davidson, 1973).
Right panel: Scheme of an accretion column at high luminosity state. The seed photons
are generated throughout the entire column via free-free and cyclotron emission, and on the
surface of the mound via black-body radiation (Becker and Wolff, 2007).
radiation are produced along the optically thin region located above the mound.
This model is able to successfully reproduce the spectral continuum of very luminous
sources: Her X-1, LMC X-4 and Cen X-3. It also describes the spectrum of the HMXB
4U 0115+63 (Ferrigno et al., 2009). In addition, as we show in Figure 2.5, the previous
version of the model by Becker and Wolff (2005) is able to describe the emergent spectrum
of a relatively faint source like X Persei. It implies that the radiation pressure might be
important even in the sources that are not the brightest ones. A possible explanation is the
increase of the scattering cross-section of electrons with photons at frequencies similar to
the resonance cyclotron ones. It would lead to a significant enhancement in the effective
pressure exerted by the radiation field.
20
Chapter 2. Acretion as a source of energy
Figure 2.5: Comparison of the spectral energy distribution between observations (points)
and theoretical models (lines) for the source 4U 0352+30 (X-Persei). The authors
(Becker and Wolff, 2005) used column densities NH = 0 (solid line), 3 × 1021 (dashed
line), 9 × 1021 cm−2 (dot-dashed line).
Part II
An XMM-Newton view of FeKα in
High-mass X-ray Binaries
In this part we present a comprehensive analysis of the whole sample of available XMMNewton observations of High-mass X-ray Binaries until August 2013, focusing on the FeKα
emission line. This line is key to better understanding the physical properties of the material
surrounding the X-ray source within a few stellar radii (the circumstellar medium).
We collected observations from 46 HMXBs and detected FeKα in 21 of them. We used
the standard classification of HMXBs to divide the sample into different groups. We find
that:
1. Different classes of HMXBs display different qualitative behaviours in the FeKα
spectral region. This is visible especially in SGXBs (showing ubiquitous Fe fluorescence but not recombination Fe lines) and in γ Cass analogues (showing both
fluorescent and recombination Fe lines).
2. FeKα is centred at a mean value of 6.42 keV. Considering the instrumental and fits uncertainties, this value is compatible with ionization states that are lower than Fe xviii.
3. The flux of the continuum is well correlated with the flux of the line, as expected.
Eclipse observations show that the Fe fluorescence emission comes from an extended
region surrounding the X-ray source.
4. We observe an inverse correlation between the X-ray luminosity and the equivalent
width of FeKα (EW). This phenomenon is known as the X-ray Baldwin effect.
5. FeKα is narrow (σline < 0.15 keV), reflecting that the reprocessing material does not
move at high speeds. We attempt to explain the broadness of the line in terms of three
possible broadening phenomena: line blending, Compton scattering, and Doppler
shifts (with velocities of the reprocessing material V ∼ 1000 km/s).
22
6. The equivalent hydrogen column (NH ) directly correlates to the EW of FeKα, displaying clear similarities to numerical simulations. It highlights the strong link between
the absorbing and the fluorescent matter.
7. The observed NH in supergiant X-ray binaries (SGXBs) is in general higher than
in supergiant fast X-ray transients (SFXTs). We suggest two possible explanations:
different orbital configurations or a different interaction compact object - wind.
8. Finally, we analysed the sources IGR J16320-4751 and 4U 1700-37 in more detail,
covering several orbital phases. The observed variation in NH between phases is
compatible with the absorption produced by the wind of their optical companions.
These results clearly point to a very important contribution of the donor’s wind in the
FeKα emission and the absorption when the donor is a supergiant massive star.
Part II is structured as follows. In Chapter 3 we describe the Fe complex and the XMMNewton observatory. In Chapter 4 we present the set of sources, the reduction process and
the more important details concerning the spectral fits. In Chapter 5 we show our results: a
spectral atlas that includes every fit and different plots relating fit parameters. In Chapter 6
we interpret the obtained results and summarize the most important conclusions in Chapter 7. In Appendix A we show the spectral atlas, which contains the plot of every spectrum
that we have analysed in this survey. We show the observations and the models, together
with the ratio between them. We also present in Appendix A a set of tables that describe the
obtained parameters from the spectral fits.
The main content of this part has been published in the peer-reviewed journal Astronomy
& Astrophysics with the following reference:
• An XMM-Newton view of FeKα in HMXBs, Giménez-García, A., Torrejón, J.M., Eikmann, W., et al. 2015, A&A, 576, A108
Chapter 3
The Fe complex with XMM-Newton
3.1 The Fe complex
Since the early stages of X-ray astronomy, iron lines in the spectral region of ∼6-7 keV
(the Fe complex) have been studied in a large number of X-ray sources given its fruitfulness
as a tool for plasma diagnostics. They were reported for the first time in the supernova remnant Cas A (Serlemitsos et al., 1973), and only two years later in a high-mass X-ray binary
(HMXB) using the Ariel 5 satellite (Sanford et al., 1975). The most recent X-ray space
missions (Swift, Suzaku, Chandra, and XMM-Newton) have triggered a notable improvement in the attainable spectral resolution and effective area, permitting between different
emission features in the Fe complex to be distinguished: narrow and broad fluorescence
lines (FeKα and FeKβ), Compton shoulders and recombination lines (Fe xxv and Fe xxvi)
(Torrejón et al., 2010b). This improvement has given a remarkable impetus to the study of
the Fe complex, and it justifies a comprehensive analysis in HMXBs.
In particular, FeKα has been proven to be a fundamental tool in the study of HMXBs
(Martínez-Núñez et al., 2014; Rodes-Roca et al., 2011; van der Meer et al., 2005). The origin of the fluorescence-emitting region has been discussed by many authors in the past.
Nagase (1989) considered accretion disks and the matter stagnated in the accretion and ionization wakes in the stellar wind as plausible areas of FeKα production. Watanabe et al.
(2006a) analysed the classical HMXB Vela X-1 and proposed the extended stellar wind,
reflection by the stellar photosphere, and an accretion wake as the most likely candidates
for fluorescence-reprocessing regions. In any case, FeKα is very sensitive to the physical
conditions of the vicinity of the X-ray source, so it provides remarkable information that
must be analysed.
Fluorescence is produced as a consequence of the X-ray illumination of matter. When
an Fe atom absorbs a photon carrying sufficient energy to remove an electron from its Kshell (E>7.2 keV), the vacancy can be occupied by another electron from an outer shell
(Figure 3.1). If the electron comes from the L shell, the transition produces FeKα emission.
This emission is produced when the vacancy is filled by a former M-shell electron. When
Fe is more ionized than Fe xix, the fluorescence yield starts to decrease with the ionization
25
26
Chapter 3. The Fe complex with XMM-Newton
Figure 3.1: Production of FeKα in a neutral iron atom. An electron from the K-shell absorbs
a highly energetic photon and triggers a cascade of transitions by the electrons from outer
shells, in order to fill the vacancies.
state (Kallman et al., 2004). Therefore, FeKα is a footprint of not extremely ionized Fe
(less than Fe xx). On the other hand, recombination lines Fe xxv and Fe xxvi unveil the
presence of very hot gas, where Fe atoms are almost completely stripped.
Previous comprehensive surveys of the Fe complex in HMXBs were carried out by
Gottwald et al. (1995) using EXOSAT and Torrejón et al. (2010b) using the High-Energy
Transmission Gratings (HETGS) onboard Chandra. The high spectral resolution provided
by Chandra gratings proved to be instrumental in disentangling the different ionization
species present in the Fe complex. However, the relatively low throughput of the instrument only allowed studying the brightest binaries. In this study we increase the previous
sample significantly by using the high throughput of XMM-Newton EPN. This has allowed
us to include fainter systems (such as Be X-ray binaries (BeXBs) or SFXTs in quiescence),
while the moderate resolution of the EPN CCDs has allowed us to test previous correlations
based on a small sample.
3.2 The XMM-Newton observatory
The XMM-Newton observatory (Lumb et al., 2012) is fitted with three X-ray telescopes
of 1500 cm2 and a coalignated optical telescope. Spectroscopy and photometry are done by
the six instruments on board: three X-ray imaging cameras EPIC (European Photon Imaging Camera), two grating X ray spectrometers RGS (Reflection Grating Spectrometer), and
an optical monitor (OM). EPIC cameras (0.1-15 keV) are the only instruments at XMMNewton that cover the energy range of the Fe complex. Among EPIC, one camera uses PN
CCDs, and the other two use MOS CCDs. EPIC PN cameras (EPN) surpass the effective
area of the MOS cameras at 6-7 keV by a factor ∼3, making EPN more suitable for our
purposes. Compared to other missions, the High Energy Transmission Grating Spectrometer (HETGS) onboard Chandra provides better energy resolution in the energy range of the
3.2. The XMM-Newton observatory
27
Figure 3.2: View of the inner structure of the XMM-Newton observatory. We can see the
Focal Plan Assembly (FPA), consisting of the Focal Plane Platform with the EPIC and RGS
detectors. The Mirror Support Platform (MSP) carries the mirrors, the gratings, the optical
monitor and two star trackers. Finally, the Service Module (SVM) carries associated units
necessary for the operational tasks.
Fe complex, but the effective area available with EPN is significantly higher. EPN provides
the adequate conditions for performing the study presented here, on account of the moderate (but sufficient) spectral resolution (∆E/E ∼ 40) and large effective area (∼ 1000 cm2 ),
enabling us to analyse a large amount of sources in a homogeneous and consistent way.
Chapter 4
The sources and data treatment
4.1 The sources
We followed the catalogue of Liu et al. (2006), in addition to later discoveries or confirmations, to identify the currently known HMXBs, and used every available XMM-Newton
public observation1 . The sources not included in the Liu catalogue, but considered here are
HD 119682, SS 397, IGR J16328-4726, HD 45314, HD 157832, Swift J045106.8-694803,
IGR J16207-5129, and XTE J1743-363. We show the list of the sources in Table 4.1, giving
the class where we have grouped them and the reference for such a classification.
HMXBs are usually variable sources and we often observe a dramatic change in luminosity even in the same observation, thus remarkably affecting the spectral parameters. In
these cases, an averaged spectrum does not reproduce the actual emission of the source, and
it is advisable to split the observation into more than one time interval. We have considered
five different states2 of the systems in order to define the time intervals: dips, quiescence,
flares, eclipse ingress/egress, and eclipse. We used the following criteria. When luminosity
drops a factor &2 on the timescale of .1 hour, we tagged the time interval as a dip. Analogously, when luminosity rises &2 on the timescale of .1 hour, we labelled the time interval
as a flare. For observations covering eclipsing phases, we defined time intervals for eclipse
ingress/egress and eclipse. The rest of time intervals are tagged as quiescent states.
In Figure 4.1 we see the light curve of an observation of 4U 1538-522 as an example of
how we have split the time intervals in the observations. The source was observed during
the ingress in an X-ray eclipse, which is clearly noticeable in the light curve. We separated
the observation into two time intervals, one covering the ingress in eclipse and another one
covering the eclipse.
In summary, we have collected data from 46 HMXBs. Twenty-one of them exhibit FeKα
emission. We note that some sources have more than one available observation. Taking
1
http://xmm.esac.esa.int/xsa/
The states considered in this work and those also called states in black hole binary systems must not be
confused.
2
29
30
Chapter 4. The sources and data treatment
Figure 4.1: Light curve of the observation of 4U 1538-522 (ObsID:0152780201). We have
split the observation in two parts, one for the ingress in the eclipse and another one for the
eclipse.
everything into account (46 sources, temporal splitting depending on the state of the source,
and more than one observation per source in some cases), we end up with a total number of
108 spectra that we have analysed.
4.2 Data reduction
We have reduced the data using Science Analysis System (SAS), version 12.0.1. Since
the sample of observations contain a heterogenous group of HMXBs, we found different observation modes (timing and imaging) to account for the different properties of the sources.
In the brightest systems, the observations were usually performed using the timing mode,
while the faintest sources were observed using imaging modes.
Timing modes permit the arrival of photons to be processed at a high rate, since only
one CCD operates, and the information is collapsed into one dimension, allowing a fast
read out. The time resolution is as high as 30 µs (7 µs in burst mode, Kirsch et al. (2006)).
Even with the high timing resolution reached with these observation modes, pile-up is still
present in several cases, especially when the count rate is & 800 counts s−1 . We checked in
every observation whether pile-up is affecting the data, using the SAS task epatplot, and we
excised the core of the source’s point spread function in the pertinent cases. The size of the
excised region has been chosen wide enough to remove the unwanted pile-up effects (see
examples of the use of epatplot in Ng et al. (2010)).
The background-subtraction process also depends on the brightness of the source. In
4.3. Spectral fitting
31
the EPN timing mode, the PSF of the sources displaying & 200 counts s−1 will span the
whole CCD. Therefore, any area selected as a background region will be contaminated by
source photons. Since this effect is strongly energy dependent, for the brightest sources we
have chosen a method of background subtraction that is similar to the one performed in the
analysis of Vela X-1 by Martínez-Núñez et al. (2014), where a blank sky spectrum taken in
timing mode is used as the real background for energies below 2.5 keV, while the rest of the
spectrum corresponds to the outermost pixels of the CCD. Meanwhile, for common observations, we have used source-free regions to extract a background spectrum and subtract it
from the former source plus background energy distribution.
Ancillary response files were generated using the SAS task arfgen. For observations
taken in timing mode affected by pile-up, we followed the recommendations of the XMMNewton SAS User Guide in order to generate the appropriate ancillary response files. Response matrices were created using the SAS task rmfgen.
4.3 Spectral fitting
For the spectral analysis we used XSPEC, version 12.8.03 . We rebinned the spectra to
have a minimum of 20 counts per bin and a bin size of at least one-third of the FWHM of
the intrinsic energy resolution, in order to be allowed to apply χ2 statistics in the fitting of a
set of Poissonian data (Cash, 1979).
In Table 4.2 we present the sample of models employed for the continuum in the fits.
Every model is a combination of additive and multiplicative models. An additive model
stands for a source of X-rays (e.g. bremsstrahlung radiation), and a multiplicative model
represents a energy-dependent change of an additive model (e.g. photoelectric absorption).
The models presented in Table 4.2 were tested in every observation and accepted deχ2
, with n the number of bins and m the number of fitted
pending on the reduced-χ2 ( n−m
parameters). Each observation has particular characteristics, and therefore the decision of
which reduced-χ2 value is acceptable has been taken one by one. In Figure 4.2 we can see
that most of the fits result in a reduced-χ2 ≃ 1, as expected for a suitable fit. The highest
value of reduced-χ2 for an accepted model has been 1.82. The parameters arising from the
fits are listed in Tables A.1 and A.2.
We can classify the additive components of the models as thermal or non-thermal. A
component is called thermal when radiation is produced as a consequence of the thermal
motion of the plasma particles (e.g. blackbody radiation). Otherwise, the emitted radiation
is non-thermal (e.g. non-thermal Inverse Compton emission). If all the additive components
of a model are thermal, we classify the model as thermal (analogously for non-thermal).
We also used hybrid models, combining thermal and non-thermal components. The thermal
components used in this work are the following:
3
http://heasarc.nasa.gov/xanadu/xspec/
32
Chapter 4. The sources and data treatment
Figure 4.2: Number of accepted models depending on the reduced-χ2 value.
• bbody: blackbody emission.
• diskbb: model of an accretion disk emission made of multiple blackbody components.
• bremss: thermal bremsstrahlung emission (electrons distributed according to the MaxwellBoltzmann distribution).
• mekal: emission from optically thin hot gas, including spectral lines from several
elements (Mewe et al., 1985).
• cemekl: built from the mekal model, incorporating multi-temperature emission.
On the other hand, the only non-thermal component used in this work is:
• powerlaw: phenomenological model consisting of a simple inverse power law profile
(∝ E −Γ ). This profile is a footprint of inverse-Compton scattering by hot electrons
(non-thermally distributed) of a seed radiation field.
For the photoelectric absorption, we used tbnew4 , the improved version of the Wilms et al.
(2000) model tbabs, setting the cross sections to the Verner et al. (1996) ones and the abundances according to Wilms et al. (2000). The most important parameter of this model is the
total equivalent hydrogen column NH , which is the integrated amount of hydrogen atoms in
the line of sight from the observer to the source, per cm2 . We also added the model cabs
to account for the Compton scattering, which is not comprised in the tbnew model and is
especially significant for NH & 1024 cm−2 .
The emission lines are fitted using Gaussian profiles. We have categorized any emission
line that fulfils the following conditions as FeKα:
4
http://pulsar.sternwarte.uni-erlangen.de/wilms/research/tbabs/
4.3. Spectral fitting
33
1) The centroid energy of the Gaussian component lies in the interval [6.3, 6.65] keV.
The interval includes the expected energy of FeKα emission from Fe ii (∼ 6.395 keV)
to Fe xxiii (∼ 6.63 keV) (Kallman et al., 2004). This condition excludes the detection of any hypothetical fluorescent emission from Fe xxiv-xxv at ∼ 6.67 − 6.7 keV,
thereby excluding any confusion between FeKα and the recombination line Fe xxv
at similar wavelength. The fluorescence yields of Fe xxiv-xxv are low compared to
lower ionization states.
2) The statistical significance (σ sign ) of the Gaussian component is greater than 2σ. We
calculated σ sign from χ2k1 − χ2k2 , assuming χ2k1 − χ2k2 ∼ χ2k1 −k2 5 , where χ2k1 arises from
a fit using certain model with the Gaussian component included, and χ2k2 arising from
a fit using the same model without the Gaussian component.
In some cases, FeKα line is clearly noticeable, but FeKβ is not prominent enough to
permit erroneous estimation of its parameters. In these cases, we have constrained the
centroid energy and the norm of FeKβ according to Kallman et al. (2004):
• Energy(FeKβ) = Energy(FeKα) + 0.652 keV
• Norm(FeKβ) = Norm(FeKα) × 0.13 photons/cm2 /s.
The estimated parameters, like the EW, are very sensitive to the fit of the continuum.
Therefore, although the Fe complex appears in the ∼6-7 keV energies, we broadened the
spectral scope to an energy range of 1-10 keV to perform the analysis. It also allows us
to consider possible calibration inaccuracies in the charge transfer inefficiency (CTI) and
the X-ray loading (XRL), an issue reported in previous analysis of EPN observations (see
Martínez-Núñez et al. (2014) and Fürst et al. (2011)). In the few cases of possible CTI or
Eold
+ o f f set (see Table A.1).
XRL, we applied an artificial gain Enew = slope
The estimation of the parameter confidence regions (at 90% level) have been calculated
with a Markov Chain Monte Carlo (MCMC) technique, implemented in XSPEC, where N
generations of the set of free parameters are used to determine the best-fit values and the
confidence regions. We set N = 1.5 × 104 in our calculations. These chains are also valid
for estimating fluxes and equivalent widths.
5
This assumption is not strictly true, since χ2k1 and χ2k2 are not independent. However, it provides an estimation of the impact of the Gaussian component in the model.
34
Class
BeXB
SGXB
SFXT
γ-Cass-like
Chapter 4. The sources and data treatment
Source
Distance
(kpc)
Reference
(class)
Reference
(distance)
1A 0535+26
2S 1845-024
X Persei
AX J1820.5-1434
RX J0146.9+6121
RX J0440.9+4431
RX J1037.5-5647
SAX J2103.5+4545
Swift J045106.8-694803
V0332+53
2.00 ± 0.70
10.00
0.80 ± 0.14
8.20 ± 3.50
1.95
3.30 ± 0.50
5.00
3.20 ± 0.80
50.60 ± 2.10
7.50 ± 1.50
Shenavrin et al. (2011)
Bodaghee et al. (2007)
Liu et al. (2006)
Liu et al. (2006)
Liu et al. (2006)
Liu et al. (2006)
Liu et al. (2006)
Liu et al. (2006)
Beardmore et al. (2009)
Liu et al. (2006)
Steele et al. (1998)
Grimm et al. (2002)
Megier et al. (2009)
Kinugasa et al. (1998)
Kharchenko et al. (2005)
Reig et al. (2005)
Grimm et al. (2002)
Baykal et al. (2002)
Bartlett et al. (2013)
Negueruela et al. (1999)
6.40 ± 1.00
2.12 ± 0.34
5.00
8.00 ± 2.50
3.04
6.10
3.60 ± 2.60
3.50
12.50
12.40
1.90 ± 0.20
5.00 ± 4.00
Liu et al. (2006)
Liu et al. (2006)
Liu et al. (2006)
Liu et al. (2006)
Liu et al. (2006)
Nespoli et al. (2008)
Filliatre and Chaty (2004)
Coleiro et al. (2013)
Walter and Zurita Heras (2007)
Torrejón et al. (2010a)
Liu et al. (2006)
Liu et al. (2006)
Reynolds et al. (1992)
Megier et al. (2009)
Cox et al. (2005)
Mason et al. (2009)
Kaper et al. (2006)
Nespoli et al. (2008)
Filliatre and Chaty (2004)
Rahoui et al. (2008)
Smith (2004)
Torrejón et al. (2010a)
Sadakane et al. (1985b)
-
Romano et al. (2011b)
González-Galán et al. (2014b)
Sguera et al. (2006)
Fiocchi et al. (2013)
Sguera et al. (2006)
Romano et al. (2011a)
Smith et al. (2006)
Pellizza et al. (2006b)
Sguera et al. (2007a)
Rahoui and Chaty (2008)
Nespoli et al. (2007)
Negueruela and Reig (2004)
Masetti et al. (2006)
Fiocchi et al. (2013)
Drave et al. (2013)
Chaty et al. (2008)
Negueruela et al. (2006)
Pellizza et al. (2006b)
Coe et al. (1996)
Rahoui and Chaty (2008)
0.12 ± 0.01
0.39 ± 0.05
1.11
0.53
1.50 ± 0.50
1.50 ± 0.26
2.20 ± 0.60
1.50 ± 0.20
Lopes de Oliveira et al. (2010)
Smith and Balona (2006)
Rakowski et al. (2006)
Lopes de Oliveira and Motch (2011)
Lopes de Oliveira et al. (2006)
Rauw et al. (2013)
Lopes de Oliveira et al. (2010)
Lopes de Oliveira (2007)
Megier et al. (2009)
Megier et al. (2009)
Kharchenko et al. (2005)
Lopes de Oliveira and Motch (2011)
Lopes de Oliveira et al. (2006)
Megier et al. (2009)
Lopes de Oliveira et al. (2006)
Riquelme et al. (2012)
Dubus (2013)
Dubus (2013)
Grimm et al. (2002)
Casares et al. (2005)
2.90 ± 0.20
10.00 ± 1.00
2.10 ± 0.25
16.00 ± 3.50
Blay et al. (2006)
Liu et al. (2006)
Liu et al. (2006)
Karasev et al. (2010)
Riquelme et al. (2012)
Hutchings et al. (1979)
Ziółkowski (2005)
Karasev et al. (2010)
4U 1538-522
4U 1700-37
4U 1907+09
EXO1722-363
GX 301-2
IGR J16207-5129
IGR J16318-4848
IGR J16320-4751
IGR J16465-4507
SAX J1802.7-2017
Vela X-1
XTE J0421+560
AXJ1841.0-0536
IGR J00370+6122
IGR J11215-5952
IGR J16328-4726
IGR J16418-4532
IGRJ16479-4514
XTE J1739-302
IGR J17544-2619
IGR J18450-0435
IGR J18483-0311
γ Cassiopeiae
HD 110432
HD 119682
HD 157832
HD 161103
HD 45314
SAO 49725
SS397
HMGB
LS I +61 303
LS 5039
O9.5V+NS
disk-fed
BH-SGXB
B1-2
4U 2206+54
Cen X-3
Cygnus X-1
AX J1749.1-2733
7.80 ± 0.74
3.00
6.20
6.50 ± 3.50
13.00
7.50 ± 2.50
2.30 ± 0.60
3.20 ± 1.00
3.60
3.50 ± 0.50
2.50
2.50 ± 0.10
Table 4.1: Table of sources included in the sample of HMXBs.
Continuum models
Non-thermal
N1
N2
N3
powerlaw × tbnew × cabs
powerlaw1 × tbnew1 × cabs1 + powerlaw2 × tbnew2 × cabs2
powerlaw1 × tbnew1 × cabs1 + powerlaw2 × tbnew2 × cabs2 + powerlaw3 × tbnew3 × cabs3
Thermal
T1
T2
T3
T4
T5
T6
T7
T8
T9
T 10
T 11
T 12
T 13
mekal × tbnew × cabs
(mekal + mekal) × tbnew × cabs
(mekal + mekal + mekal) × tbnew × cabs
(cemekl) × tbnew × cabs
bbody × tbnew × cabs
(bbody + bbody) × tbnew × cabs
(bbody1 + bbody2 ) × tbnew × cabs
bbody1 × tbnew1 × cabs1 + bbody2 × tbnew2 × cabs2
diskbb × tbnew × cabs
(diskbb + bbody) × tbnew × cabs
bremss × tbnew × cabs
bbody × tbnew × cabs + bremss × tbnew × cabs
(bremss + bbody) × tbnew × cabs
Both
T N1
T N2
T N3
T N4
T N5
4.3. Spectral fitting
Model
(powerlaw + bbody) × tbnew × cabs
powerlaw × tbnew1 × cabs1 + bbody × tbnew2 × cabs2
(powerlaw + diskbb) × tbnew × cabs
powerlaw × tbnew1 × cabs1 + diskbb × tbnew2 × cabs2
(powerlaw + mekal) × tbnew × cabs
35
Table 4.2: List of models used to fit the continuum, described in XSPEC notation. The basic components are powerlaw, bbody, diskbb, bremss,
mekal, and cemekl, together with tbnew, to account for the absorption and cabs for the non-relativistic Compton scattering. The employed
models are a combination of these components, in addition to Gaussian profiles modelling emission lines. We divide the models into three
types: N# for non-thermal, T # for thermal, and T N# for models containing both thermal and non-thermal components.
Chapter 5
Results
5.1 Spectral atlas
In Appendix A, we present the full sample of analysed spectra. We show the set of analysed observations (cross points), the model employed (solid line), the components of the
model (dotted line), and the ratio between observation and model (lower box in each spectrum plot).
We can see that the different classes of HMXBs behave qualitatively differently in the
region of the Fe complex (∼6-7 keV), reflecting the distinct accretion regimes that characterize them. We have observed three patterns in the Fe complex, which we define as Types I,
II, and III (see Figure 5.1). We define Type I, when fluorescence lines FeKα and FeKβ are
observed, but not recombination lines Fe xxv and Fe xxvi. We define Type II, when fluorescence lines are detected, together with recombination lines Fe xxv and Fe xxvi. Finally, we
define Type III, when Fe lines are not detected.
The general features observed in this work for the different groups of HMXB are summarized in Table 5.1, and explained below in more details:
• BeXBs. We collected data from ten sources. All the observations were performed in
quiescence. We have detected FeKα emission in only one BeXB (SAX J2103.5+4545).
The upper limit of the FeKα EW in the rest of BeXBs is in general higher than the
observed value in SAX J2103.5+4545, implying that the lack of detections might be
due to a poor signal-to-noise. The spectra can be modelled by thermal or a combination of thermal and non-thermal components, except for Swift J045106.8-694803
(fitted using an absorbed power law). Seven sources accept a thermal model, and six
a combination model (4 of them accept both).
• SGXBs. We have gathered data from 12 sources. Ten of them show detectable Fe
fluorescence emission. The only exceptions are IGR J16465-4507 and SAX J1802.72017, the most distant SGXBs included in this work, at 12.5 and 12.4 kpc, respectively. The EW upper limits in these two sources are high, implying that their faintness is very likely the reason we do not detect FeKα. The 12 SGXBs can be well fitted
37
38
Chapter 5. Results
using non-thermal models, although thermal components are also plausible in some
sources. In general, SGXBs are characterized by high absorption and the presence of
Fe fluorescence emission lines.
• SFXTs. We have collected data from ten sources. Three of them show FeKα:
AX J1841.0-0536, IGR J11215-5952, and IGR J16479-4514. The EW upper limit
in the rest of sources is high. Therefore, FeKα would probably be detectable with a
better signal-to-noise. The models employed for fitting the SFXT systems are very
heterogeneous, with no preference for thermal or non-thermal, or for a combination
of both kinds of models.
• γ Cassiopeae analogues. We have gathered observations from eight sources. Five of
them exhibit FeKα. The EW upper limit in the three other sources is very high. Again,
it implies a very likely presence of fluorescence in the case of better signal-to-noise.
In addition, recombination lines of Fe xxv and Fe xxvi are always present in the set
of γ Cassiopeae analogues. These lines are included in the XSPEC model mekal. For
most of the observations we have achieved a good fit using a combination of mekal
components. In a few cases we used other components: diskbb and powerlaw, but
mekal is by far the one employed most in γ Cassiopeae-like systems, in agreement
with previous X-ray analyses (Lopes de Oliveira et al., 2010, 2006).
• HMGBs. We collected data from two HMGBs: LS I+61 303 and LS 5039. None
of them show Fe features. However, the signal-to-noise in these observations is poor
and the upper limits of the FeKα EW are high enough to not rule out the presence of
the line. We have used both thermal and non-thermal components in the fits.
• Peculiars. These are a set of sources that do not fit into any of the aforementioned
classes of HMXBs, as explained in the introduction. We collected data from three
such systems: 4U 2206+54, Centaurus X-3, and Cygnus X-1:
– 4U 2206+54 does not show any detectable Fe emission line, and the upper limit
in the EW of FeKα is low (comparable to the upper limits in the BeXBs). It can
be fitted by means of an hybrid model (thermal plus non-thermal components).
– Centaurus X-3 presents a rich emission-line spectrum. Concretely, in the Fe
complex we are able to identify FeKα, FeKβ, Fe xxv, and Fe xxvi. We used
either an hybrid model either a non-thermal model.
– Cygnus X-1 exhibits a broad Fe feature, sometimes combined with a faint and
narrow, but statistically significant, FeKα line. We have mostly used nonthermal models, occasionally combined with a bbody or a diskbb component.
• AX J1749.1-2733. In this system, the optical member has been classified as a B1-2
(Karasev et al., 2010), but the luminosity class remains unknown, preventing us to
incorporate the source in a defined group. Although it does not exhibit detectable Fe
emission, the high absorption clearly points to a supergiant companion. In addition,
the EW upper limit of FeKα is compatible with the values observed in SGXBs and
SFXTs. It can be well fitted using an absorbed powerlaw or a blackbody.
5.2. FeKα width
39
Group
# Sources
Fe complex
Models
NH
BeXB
10
Type III
T, TN
Low
SGXB
12
Type I
N
High
SFXT
10
Type III
(in quiescence)
T, N, TN
High
γ Cass like
8
Type II
T
Low
HMGB
2
Type III
T, N, TN
Low
Table 5.1: Description of the features observed in this work for the different groups of
HMXBs. We have analyzed data of 46 sources. However, those classified as peculiars (3
sources) or unclassified (AX J1749.1-2733) are not included in this table. We define NH of a
group as high, when the typical value observed is well over the estimations of the interstellar
NH in the line of sight of the sources following Willingale et al. (2013). That is, we say that
the NH of a group is high when the absorption is typically intrinsic to the systems.
Fe complex: Type I
Fe complex: Type II
Fe complex: Type III
15
10
5
0
1.2
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1.5
1
0.5
0
0.6
0.4
0.2
0
1.2
1.05
1
ratio
1.1
ratio
ratio
1.1
1
0.9
1
0.9
0.95
0.8
0.8
6
7
Energy (keV)
7
6
Energy (keV)
6
7
Energy (keV)
Figure 5.1: Patterns found in the Fe complex of HMXBs: Type I (left panel), Type II (central
panel), and Type III (right panel).
In summary, it is very likely that FeKα is a ubiquitous feature in HMXBs, and its detection strongly depends on the quality of the observations. In this regard, the EW of the
line is very affected by the level of intrinsic absorption present on the sources (see also
Section 5.4.3). SGXBs and SFXTs, which show higher absorption, tend to exhibit a more
prominent Fe fluorescence.
5.2 FeKα width
In Table A.2 we show the parameters of every detected FeKα, including the width of
the line (σline ). We made a distinction between narrow and broad FeKα. We defined narrow FeKα as when σline < 0.15 keV, and broad FeKα as when σline > 0.15 keV. This
separation is both physically and observationally motivated.
Chapter 5. Results
Narrow FeK
40
Broad Fe features
Figure 5.2: Histogram of the FeKα width. The bulk of the detections are grouped showing
σline < 0.15 keV. We define them as narrow FeKα. The rest are defined as broad Fe features.
Even though we have detected 60 narrow FeKα, this plot only includes 38. The reason is
that 22 of them are very narrow (or the signal-to-noise very low) and they have been treated
in the fits as delta functions. In Table A.2 we present their width as σline = 0.
The origin of broad Fe features in X-ray binaries is still an open question, but the most
likely alternatives are related to the presence of an accretion disk (see e.g. Hanke et al.
(2009); Ng et al. (2010); Duro et al. (2011)). However, narrow features are not compatible with material rotating at high velocities or being relativistically broadened. Given that
broad and narrow FeKα have clearly different origins, they must be analysed in different
ways. Then, it raises the question of how to define the separation mark between them. In
Figure 5.2 we can see that in our sample the number of detected sources decreases when
increasing σline . Moreover, most of the detections are grouped at σline < 0.15 keV. As a
result, σline < 0.15 keV seems a natural threshold for the definition of narrow lines in the
sample. In addition, we must pay attention to the plausible contamination of FeKα with
Fe xxv, which is located at ∼ 6.7 keV. The chosen criterion separates the sources where
it is very unlikely that FeKα is contaminated by Fe xxv (narrow lines), from the sources
that probably suffer from this problem (broad lines). A more detailed analysis of broad Fe
features in HMXBs is beyond the scope of this paper and it will be discussed fully in a
forthcoming work.
Hereafter, when FeKα is mentioned, we refer to the narrow feature. From the total
number of 108 analysed observations we find detected (narrow) FeKα in 60 of them.
5.3. Centroid energy
41
5.3 Centroid energy
In Figure 5.3 we can see a histogram that presents the centroid energy of FeKα. A Gaussian fit of the data reveals a mean value for the centroid energy of 6.42 keV. There are
no significant differences in the averaged values obtained for different classes of HMXBs
or for different states. The standard deviation is 0.02 keV, comparable to the error that
we typically obtain in the estimation of the centroid energy in the fits (see Table A.2).
After taking the standard deviation and the uncertainties in the CTI corrections in EPN1
into account, the centroid energy of FeKα constrains the ionization state of Fe to less ionized than Fe xviii (Kallman et al., 2004), in agreement with previous studies in HMXBs
(Torrejón et al., 2010b; Gottwald et al., 1995; Nagase, 1989). In this regard, the study of
Torrejón et al. (2010b) using HETGS (more accurate in wavelength than EPN) gives a narrower constraint on the ionization state (Fe i-x). Our present work supports this result and
adds more sources to the sample.
On the right-hand side of Figure 5.3 we can see seven FeKα detections emerging from
the Gaussian profile. Four of them are unlikely to be described by such a Gaussian profile,
since they lie more than three times the standard deviation away from the mean energy. All
four belong to Cygnus X-1.
5.4 Correlated parameters
One of the goals of this work is to study plausible correlations that involve the parameters
of FeKα (position, width, intensity, and EW, and other parameters in the fits, such as the
absorbing column and the intensity of the continuum. Even when a good fit is reached, the
confidence region of a parameter might be occasionally difficult to find owing to the dependence of the parameter on other parameters of the model. In each of the following sections,
we specify the number of cases where a successful estimation of the 90% confidence region
has been done.
5.4.1 Continuum flux vs FeKα Flux
In Figure 5.4 we represent the unabsorbed flux of the continuum between 1-10 keV cancelling FeKα emission (F1−10keV ), against the flux of FeKα (F FeKα ). We have successfully
found a 90% confidence region of the flux of FeKα in 56 cases.
On a logarithmic scale, we identify two different patterns of correlation. First, for a
subset including all the eclipse observations and IGR J16318-4848, we find a correlation
with Pearson coefficient (PC) of 0.98. Second, for the rest of the observations, we find a
correlation with PC=0.89.
1
Please find more information about long-term CTI correction in the release note EPICpn Long-Term CTI, by M.J.S Smith et al. (2014), at http://xmm2.esac.esa.int/docs/documents/CALSRN-0309-1-0.ps.gz; and EPIC status of calibration and data analysis by Guainazzi (2008), at
http://xmm2.esac.esa.int/external/xmm_sw_cal/calib/CAL-TN-0018.pdf.
42
Chapter 5. Results
Figure 5.3: Centroid energy of FeKα with a Gaussian fit overplotted (blue profile). The
mean value is 6.42±0.02 keV, compatible with Fe i-xviii. Even though we have 60 detections
of FeKα, in this plot we only see 55. Five cases fulfil the requirements of detection, but the
low signal-to-noise rations do not permit finding an accurate centroid energy. They have
not been included in this plot. In these five cases we set the centroid energy of FeKα to
∼6.4 keV.
Figure 5.4: F1−10keV versus F FeKα . Blue dashed line marks the correlation observed for
IGR J16318-4848 jointly with eclipse observations, and the red solid line follows the bulk
of the observations. The colour map indicates the σ sign of the line (defined in Section 4.3).
5.4. Correlated parameters
43
A linear fit of the parameters (on logarithmic scale) in the out-of-eclipse observations
(Figure 5.4) gives the following dependence:
1.00±0.08
F1−10keV (erg/s/cm2 ) = F FeKα
(erg/s/cm2 ) × 102.18±0.87 .
(5.1)
The errors show the standard deviation of the parameters in the linear fit.
The observed divergence amongst eclipse (plus IGR J16318-4848) and out-of-eclipse
observations suggests that the companion star blocks the continuum and the FeKα emission
in different proportions. Therefore, an important contribution of the fluorescence emission
is produced in an extended region of R & R⋆ . This is consistent with previous analysis of
eclipse observations of HMXBs (e.g. Rodes-Roca et al. (2011) and Audley et al. (2006)).
In particular, Audley et al. (2006) estimate that 20% of FeKα in OAO 1657-415 is emitted
from 19 light seconds off the X-ray source.
We also have the luminosity of the continuum to compare with the EW of FeKα. For the
flux-to-luminosity conversion, we used the estimations of the distance shown in Table 4.1.
We have excluded eclipse and IGR J16318-4848 from this analysis, given that the EW
is strongly affected by the high obscuration of the continuum that they suffer from. In
Figure 5.5 we plot the EW of FeKα against the unabsorbed luminosity of the continuum
between 1-10 keV cancelling FeKα emission (L1−10keV ). We observe two different groups
of sources: 1) γ Cassiopeae analogues lying at low luminosities (L1−10keV < 1033 erg/s);
2) the rest of sources that exhibit FeKα. The γ Cassiopeae analogues do not show any
evident correlation (there are very few points), while the rest present a moderate inverse
correlation (PC=-0.25, and PC=-0.39 using only the sources with an available estimation
of distance with error, marked as filled diamonds in Figure 5.5). A linear fit in Figure 5.5
leads to
17.45±9.83
EW = L−0.52±0.27
.
(5.2)
1−10keV (erg/s) × 10
Baldwin (1977) observed an inverse correlation in the EW of CIV and the UV luminosity
in AGNs. Analogously, the decrease in the EW of FeKα when increasing the X-ray luminosity is called the X-ray Baldwin effect. The dependence that we observe is compatible
within the error with the one observed by Torrejón et al. (2010b) in X-ray binaries using a
narrower energy range: EW ∝ L−0.29 .
1.6−2.5Å
5.4.2 FeKα Width vs centroid energy
In Figure 5.6 we present the centroid energy of this feature versus its width (σline ). We
have successfully found a 90% confidence region of σline in 20 cases. We can see a moderate correlation (PC=0.55), indicating a possible blending of lines. Two observations (uppermost side of Fig. 5.6) do not follow the correlation. They correspond to observations of
4U 1700-37 (Obs.ID 0083280201) and EXO 1722-363 (Obs.ID 0405640201) where the Fe
complex is hardly resolved, and therefore it is very likely that a contribution of Fe xxv and
Fe xxvi in the model of FeKα is increasing the measured width of the FeKα line.
44
Chapter 5. Results
Figure 5.5: EW of FeKα against L1−10 keV . γ Cassiopeae analogues (circles) lie at
L1−10 keV < 1033 erg/s. Open symbols indicate that either the distance or the error in the
estimation of the distance is unknown. The solid line corresponds to a linear fit on logarithmic scale of the filled diamonds, that is, the sources with available distance with error
estimation and L1−10 keV > 1033 erg/s.
Coloured squares correspond to the expected width from the contribution of three different broadening phenomena: line blending, Doppler shifts, and Compton broadening. A
discussion of the different broadening contributions is given in Chapter 6.
5.4.3 Curve of growth
In Figure 5.7 we show, for out-of-eclipse observations, the NH versus the EW of FeKα
(what is generally known as the curve of growth). We have successfully found a 90% confidence region of both NH and EW in 46 cases. We want to take observations where NH
reflects the intrinsic absorption of the system into account, so we set NH > 2 as a condition to safely exceed the typical NH of the interstellar medium for the sources here studied
(checked using the online application following Willingale et al. (2013)). The use of this
criterion excludes the BeXB SAX J2103.5+4545, the γ Cassiopeae analogs: γ Cassiopeae
and HD 110432; and the SFXT IGR J11215-5952. Moreover, eclipse observations show
higher EW, and they are not comparable to out-of-eclipse observations. Therefore, eclipse
observations have not been plotted in Figure 5.7. As a consequence of the chosen criteria,
we end up with a set of 36 observations, where all the donors are supergiants.
Both NH and the EW of FeKα are expected to correlate in HMXBs (Torrejón et al.,
2010b), since the spectral lines are usually stronger when the optical depth increases. Our
sample confirms these expectations, showing a notable correlation (PC=0.85).
5.5. NH : SGXBs and SFXTs
45
Figure 5.6: Width of FeKα (σline ) versus the centroid energy (black squares). The black
solid line traces a linear fit. We have marked in colour the expected width from the contribution of three different broadening phenomena: line blending, Compton broadening, and
Doppler shifts, considering velocities of V(km/s) = 1000 (red) and V(km/s) = 2000 (green).
Every single black square has an associated single red square and a single green square
corresponding to the expected values of σline for that observation (see Chapter 6).
We have determined the theoretical curve of growth using numerical simulations. In this
simulations there is an input of X-ray radiation with a power law profile, which is transmitted through a cloud of spherically distributed neutral matter (Eikmann, 2012).
We took the power law index (Γ) in the simulations into account, since steeper profiles
entail less photons available above the Fe K-shell threshold energy, thus decreasing the EW;
that is to say, for the same NH , the higher Γ, the lower the EW of FeKα. In Figure 5.7 the
turquoise band traces the results from the simulations with Γ ∈ [0.5, 2], which is the typical
range where we find Γ in our fits.
5.5 NH : SGXBs and SFXTs
In Figure 5.8 we have plotted histograms for the NH values observed in SGXBs and
SFXTs. Where we have more than one observation for the same source, we averaged the
values to obtain one NH that is representative of each system. The orbital phase critically
affects the observed NH , and therefore ingress/egress and eclipse phases have not been taken
into account.
We find that SGXBs are in general more absorbed than SFXTs. We performed a permutation test to quantify whether the observed disparity in the NH is a random effect. We have
ten NH values for SGXBs and nine for SFXTs. We merged them in a set of 19 elements
46
Chapter 5. Results
Figure 5.7: Curve of growth observed for FeKα, that is, EW against NH . The turquoise band
marks the expected correlation using numerical simulations. The sources are identified by
different symbols when more than one observation is included: 4U 1700-37 (open circle),
4U 1907+09 (open upward triangle), Cygnus X-1 (open downward triangle), EXO1722363 (open diamond), IGR J16318-4848 (open square), and IGR J16320-4751 (plus). Only
one observation for Centaurus X-3, GX 301-2, Vela X-1, and XTE J0421+560 (all four a
star symbol).
5.6. Individual sources analysis:
IGR J16320-4751 and 4U 1700-37
47
Figure 5.8: Histograms showing a comparison of the NH values observed in SGXBs (filled
red) and SFXTs (empty blue).
and considered every possible combination of two groups of ten and nine elements (92378
possibilities). We compared the median of the two subsets and calculated the absolute difference: 99.7% of the cases have produced a lower absolute difference than the observed
one. If using the mean instead of the median, the percentage is also very high (98.8%). In
conclusion, it is very likely that the discrepancies in the observed NH values for SGXBs and
SFXTs are produced by physical reasons rather than arising by chance.
5.6 Individual sources analysis:
IGR J16320-4751 and 4U 1700-37
5.6.1 IGR J16320-4751
IGR J16320-4751 was detected by ASCA in 1994 and 1997 (corresponding to AX J1631.94752), and by INTEGRAL in 2003 (Tomsick et al., 2003). It is a HMXB composed of an
O8I optical star and a neutron star (Rahoui et al., 2008). It shows a modulation of 8.96 days,
which is considered its orbital period (Corbet et al., 2005), and a pulsation period of ∼1300s
(Lutovinov et al., 2005). The ESA archives permitted eleven observations of IGR J163204751 to be collected, enabling us to study the curve of growth in more detail, as well as to
track the absorption variation during the orbital phase.
In Fig. 5.9 we can see the curve of growth, as shown in Fig. 5.7, restricted to IGR J163204751. We clearly see the dependence between NH and EW of FeKα, as stated for the bulk
of the sources in Section 5.4.3, and the general trend following the numerical simulations.
However, the agreement with the simulations is not completely fulfilled, given that the spectral fits of IGR J16320-4751 have led to a power law index Γ ∼ 0.5 (see Table A.1). Since
48
Chapter 5. Results
Figure 5.9: EW of FeKα against NH , in IGR J16320-4751. The turquoise band marks the
numerical simulations results, with Γ ∈ [0.5, 2].
we expect more EW of FeKα from a lower power law index, the points for IGR J163204751 are expected to be located in the upper edge of the turquoise band, corresponding to
the simulations with Γ = 0.5. We consider that the general trend is correct, but there are
still some uncertainties in the fits and/or the theoretical hypothesis (spherical geometry and
neutral matter).
From 14 August to 17 September 2008, there was a campaign of nine observations of
IGR J16320-4751 by XMM-Newton. We used this set of data to plot the NH modulation
depending on the orbital phase (Fig. 5.10). We set φ = 0 at the NH maximum. We also
calculated the theoretical absorption expected from a smooth wind in a non-eccentric orbit using a β velocity law (Castor et al., 1975a) and the motion equation, considering the
variations in the orbital inclination i, orbital separation a, star radius R⋆ , mass loss rate Ṁ,
parameter β, hydrogen mass fraction XH and the terminal velocity of the wind 3∞ . This
absorption NH can be calculated from the integration along the line of sight of hydrogen
number density at distance r from the donor:
nH (r) =
XH Ṁ
mH 3∞ (1 − R⋆ /r)β 4 π r2
(5.3)
Indeed, Ṁ and 3∞ cannot be distinguished in this simple model, so the actual parameter
used is Ṁ/3∞ . However, hereafter we give values of Ṁ and 3∞ as if they were free variables, since they are much more commonly used than Ṁ/3∞ in the literature. This way, we
constrain our parameters to the observed and predicted range of values in O supergiants:
Ṁ = 10−7 − 10−5 M⊙ /yr and 3∞ = 500 − 3000 km/s (Kudritzki and Puls, 2000; Vink et al.,
2001).
5.6. Individual sources analysis:
IGR J16320-4751 and 4U 1700-37
49
Figure 5.10: Orbital modulation of NH in the system IGR J16320-4751. Solid lines correspond to the expected absorption from a smooth wind and a non-eccentric orbit, assuming
an orbital separation a = 1.6 R⋆ , R⋆ = 20R⊙ , Ṁ = 10−5 M⊙ /yr, 3∞ = 700km/s, β = 0.8;
and different orbital inclinations i = 0, π/10, π/6 rad (red, green, and blue).
For a null orbital inclination, we obviously obtain a flat NH modulation (Fig. 5.10), which
describes the observed NH (except at φ = 0), assuming a = 1.6 R⋆ , R⋆ = 20R⊙ , Ṁ =
10−5 M⊙ /yr, 3∞ = 700 km/s, and β = 0.8. Gradually increasing the orbital inclination
(i = π/10, π/6 rad ), we are able to describe a high NH at φ = 0, but losing similarities
around other orbital phases. Given the simplicity of the model, the obtained parameters are
certainly just indicative.
5.6.2 4U 1700-37
The source 4U 1700-37 was detected for the first time by Uhuru in 1970 (Jones et al.,
1973). The optical star is HD 153919, an O6.5Iaf located at 1.9 kpc (Ankay et al., 2001).
The orbital period is 3.41 days. Since X-ray pulsations have not been detected so far, the
compact object can either be a neutron star or a black hole. The database of ESA contains
five observations of 4U 1700-37, which we split into nine spectra to distinguish different
states of the source.
In Figure 5.11 we can see the curve of growth for 4U 1700-37. Although seven of the
nine spectra show FeKα, we were able to constrain the boundaries of NH in only five of
the analyses. One of them corresponds to an eclipse observation. It shows much more EW
because the continuum flux is blocked by the optical star, whereas FeKα comes from a more
extended region that is not completely hidden during eclipse. We do not see any obvious
dependence between NH and EW, although the points lie in a region close to the expected
values (turquoise band). Either way, a set of four observations (excluding the eclipse) is too
small to perform a statistical analysis.
50
Chapter 5. Results
Figure 5.11: EW of FeKα against NH , in 4U 1700-37. The filled circle corresponds to an
eclipse observation. The turquoise band traces the numerical calculations with Γ ∈ [0.5, 2].
From 17 to 20f February 2001, 4U 1700-37 was observed by XMM-Newton four times in
a campaign covering different orbital phases. We can therefore study the orbital modulation
of NH in the same way as we did with IGR J16320-4751, but including more constraints
coming from the non-LTE analysis of Clark et al. (2002), where the following parameters
are derived: R⋆ = 21.9R⊙ , Ṁ = 9.5 × 10−6 M⊙ /yr, 3∞ = 1750km/s, and β = 1.3. Considering that it is an eclipsing binary, we assume i ∼ π/2. Therefore the only free parameter in
our toy model is the orbital separation a. The best agreement is achieved when a = 1.4 R⋆
(see Figure 5.12). This orbital separation is consistent (in absolute units) with previous estimations of Conti and Cowley (1975) (R⋆ = 20 R⊙ , a = 1.35 R⋆ ) and Heap and Corcoran
(1992) (R⋆ = 18 ± 3 R⊙ , a = 2.0 ± 0.4 R⋆ ).
5.6. Individual sources analysis:
IGR J16320-4751 and 4U 1700-37
51
Figure 5.12: Orbital modulation of NH in 4U 1700-37. The solid line corresponds to the
expected absorption from a smooth wind and a non-eccentric orbit, assuming the stellar
values obtained in Clark et al. (2002), an orbital inclination i ∼ π/2 rad, and an orbital
separation a = 1.4 R⋆ .
Chapter 6
Discussion
In Figure 5.3 we have shown the centroid energies of FeKα. The distribution of the histogram is roughly Gaussian with a standard deviation that reflects the uncertainties in the
fits. However, four values are too high to be compatible with this distribution, and all of
them belong to Cygnus X-1. It can be caused either by an inadequate fit or for a physical
reason. As stated before, in Cygnus X-1 we detect a broad Fe feature, interpreted as a relativistically broadened fluorescence line. However, we modelled the relativistically broadened feature with a Gaussian profile, which gives an acceptable fit, but might be inadequate,
thereby affecting the parameters of the narrow FeKα arising from the fits. Alternatively, as
a plausible physical explanation, in Cygnus X-1 the matter is accreted via an accretion disk,
in contrast to the wind-fed accretion of most of the sources showing FeKα in this study.
Therefore, the physical properties of the region emitting fluorescence might be different in
Cygnus X-1 and the rest of systems. If this region is hotter in Cygnus X-1, the centroid
energy of FeKα would be shifted to higher energies, as we observe.
IGR J16318-4848 is one of the most absorbed systems and it presents a special configuration of matter in its surroundings where dust and cold gas distribute in a non-spherical
manner, forming a disk-like structure of matter up to ∼ 100R⋆ (Chaty and Rahoui, 2012).
A likely high inclination of the system would produce the extreme X-ray absorption and the
eclipse-like correlation between FeKα and continuum fluxes.
In Figure 5.6 we see that the centroid energy of FeKα is higher when the line is broader.
When more ionized Fe goes along with more variety in the Fe ions involved in the total
emission, the width resulting from the blending of lines must depend on the centroid energy
of the line, as observed. We have estimated the broadening produced in the lines by line
blending by σB ≈ E − 6.4 (keV), with E the centroid energy of the line in keV. We note
that it is also plausible that unresolved Fe xxv and Fe xxvi actually shift and broaden FeKα,
producing an equivalent effect.
More processes are also able to significantly broaden FeKα. We have considered Compton broadening and Doppler shifts as plausible candidates. Compton scattering has been
proposed as a possible broadening mechanism of emission lines in neutron star LMXBs
53
54
Chapter 6. Discussion
(Díaz Trigo et al. (2012), for GX 13+1). For HMXBs, Compton broadening might also be
significant, given the high NH values observed (and the consequent high number of free
electrons). However, if this process is what determines the width of the lines, we should
observe a direct correlation between the absorption column and the line width. In Figure 6.1
we can see that such a direct correlation is not present. Moreover, an inverse correlation is
plausible. Cackett and Miller (2013) have analysed three neutron-star LMXBs and arrived
at a similar result. Therefore, Compton broadening cannot be considered as mainly responsible for the observed width in HMXBs, although it is not ruled out as a modest contributor.
We assign σC to the contribution of Compton scattering in the line broadening.
To estimate σC , we used an empirical formula accurate to within 30%, derived from
Kallman (1989) and corrected in Brandt and Matt (1994):
σC = 0.019 E K τT h (1 + 0.78kT e ) ≃ 0.12 τT h (1 + 0.78kT e )
where E K ≃ 6.4 keV is the energy of FeKα, τT h is Rthe Thomson optical depth, and kT e the
electronic temperature in keV. We use τT h = σT h ne (s) ds = σT h Ne , where σT h is the
Thomson cross section, ne the electron number density, and where the integral is calculated
along the line of sight. Assuming solar abundances, a temperature kT e ≪ 1 keV and an
almost completely ionized matter, as reasonable for galactic massive stars atmospheres, we
obtain (see e.g. equation 3.61 in Novotny (1973)):
Z
X (1 + X)
ρ(s) ds
(1 + X) =
NH ≈ 0.7 NH ⇒
Ne =
2mH
2
.
⇒ σC ≈ 0.5 NH (1022 cm2 ) eV
Doppler shifts must be taken into account, since a velocity of more than 500 km/s (a very
feasible speed, either in the wind or in the accretion flow) would broaden the lines by more
than 10 eV. We assign σD to any broadening produced by Doppler effects.
Line blending, Compton scattering, and Doppler shifts produce a resultant width of
q
σtotal = σ2B + σC2 + σ2D
.
We adopted σD = 20, 40 (eV) corresponding to velocities of V ≈ 1000, 2000 km/s, which
are very plausible either in the wind of the supergiant or in the accretion flow. We overplotted the corresponding values of σtotal in Figure 5.6. For each observation, we computed the
expected value (σD = 20 (eV) and σD = 40 (eV)). The vast majority of the line widths can
be described in this way.
In IGR J16318-4848, the high absorption measured (above 2 × 1024 cm2 ) and the consequent expected Compton broadening of more than 100 eV are not congruent with the
55
Figure 6.1: Total equivalent hydrogen column (NH ) against the width of FeKα.
measured width of ∼ 35 eV. This is another indication that the absorbing matter in this system is cold and not ionized, as already stated by Chaty and Rahoui (2012). The employed
expression for describing σc therefore cannot apply here, since there are not enough free
electrons to broaden the line by means of Compton scattering.
In Figure 5.7, we show the curve of growth of FeKα. We require that NH > 2 (intrinsic
absorption rather than interstellar). In our sample, this criterion constrains the systems in
Figure 5.7 to those with supergiant donors alone. We observe a direct correlation between
NH and the EW. This correlation highlights that the X-ray absorption is strongly linked to
the matter that emits FeKα, since it is produced by matter in the line of sight, where the
X-rays are absorbed, and not in other plausible regions such as an accretion disk. (see a
sketch in Figure 6.2). In the systems included in Figure 5.7 (all with a supergiant optical
star), FeKα is produced from the transmission of X-rays through the circumstellar medium,
that is, either through the strong wind of the supergiant donor or through any structure in the
line of sight, such as ionization or accretion wakes. The hypothetical reflection of X-rays
in an independent medium might produce an additional amount of FeKα, as observed in
the BeXB GRO J1008−57 by Kühnel et al. (2013), which is not noticeable in the systems
shown in Figure 5.7.
As stated before, the region where fluorescence is emitted must be more extended than
R > R⋆ , and consequently the wind of the companion star, which is illuminated by the
X-ray source, is an obvious contributor to both the absorption and the FeKα emission. The
orbital modulation of NH shown in Figures 5.10 and 5.12 also support this interpretation.
Moreover, most of the observations track the numerical simulations assuming a characteristic range of Γ values, indicating that an isotropic distribution of absorbing (and FeKα
56
Chapter 6. Discussion
Figure 6.2: Simple sketch of two plausible configurations of circumstellar matter in
HMXBs. On the left side, X-rays are transmitted through a dense medium (e.g. the strong
wind of the donor), producing high NH directly correlated with the EW of FeKα. On the
right side, X-rays are reflected in an accretion disk producing fluorescence and also transmitted through a more diffuse medium. In this case NH is not necessarily correlated with
the EW of FeKα.
emitting) matter is not far from reality. We do not ignore the variability and heterogeneous
properties of the HMXBs environment, that might be reflected in the observed dispersion
of the curve of growth and in the moderate discrepancies regarding the simplified view of
spherically distributed neutral matter.
In Figure 5.8 we have compared the observed values of NH for SGXBs and SFXTs. We
observed that SGXBs are in general more absorbed sources than SFXTs. This implies that,
in SGXBs, either the compact object orbits a denser region of the donor wind, or else the
interaction compact object - wind modifies the environment, producing an enhancement of
density in its surroundings.
We took a look at the orbital parameters of SFXTs (see e.g. Table 2 in Romano et al.
(2014)) and find that their orbital periods lie in a wide range of values, from around three days
for IGR J16479-4514, up to 164 days for IGR J11215-5952. Some of them show high eccentricity. However, currently there is no complete description of the orbital parameters in
SFXTs. Therefore, we cannot rule out the possibility that, in this sample, SFXTs are less
absorbed than SGXBs because of the distance of the compact object to the donor star. In
this regard, further studies of orbital parameters of SFXTs will be useful.
Regarding the interaction compact object - wind, hydrodynamic simulations show that
the gravitational potential of the compact object, and the X-ray radiation field, can significantly modify the observed value of NH (Manousakis and Walter, 2011; Manousakis et al.,
57
2012). In SGXBs, where the X-ray emission is more persistent, these effects might be
stronger than in SFXTs, so notably increasing the absorption.
Chapter 7
Summary
We performed the spectral analysis of the whole sample of publicly available EPN XMMNewton observations of HMXBs until August 2013, in order to describe its FeKα emission.
In total, the study involves 46 HMXBs, 21 of them showing significant FeKα emission. As
expected, we dealt with a very heterogenous set of objects and states of the sources, which
must be properly organized. We classified the systems in the following groups: BeXBs,
SGXBs, SFXTs, γ Cass analogues, HMGBs, and peculiar sources. Furthermore, we divided
the observations depending on the source behaviour in the following states: quiescence,
flare, eclipse ingress/egress, and eclipse. With these criteria, we finally had a set of 108
spectra for our analysis, which led to the following conclusions:
• The spectral atlas gives a qualitative description of the different groups of HMXBs,
especially recognizable for SGXBs (fluorescence but not recombination Fe lines),
and γ Cass analogues (modelled by mekal models and presenting fluorescence and
recombination Fe lines). FeKα is very likely a ubiquitous feature in HMXBs, but its
detection strongly depends on the quality of the observations. SGXBs and SFXTs,
which show the higher NH among the HMXBs, tend to exhibit a more prominent Fe
fluorescence.
• The value of the centroid energy of FeKα constrains the ionization state of the reprocessing material to be below Fe xviii.
• The FeKα and continuum fluxes are well correlated, as expected for the fluorescence
emission of matter illuminated from an X-ray source. The different coefficients of
correlation for eclipse and out-of-eclipse observations agrees with previous eclipse
observations of HMXBs, in the sense of showing that the FeKα is produced in a
region that ranges from the vicinity of the X-ray source to distances that are close to
or greater than the stellar radius.
• We confirm an inverse correlation between the X-ray luminosity and the EW of FeKα
X-ray Baldwin effect. The γ Cass analogues do not follow this correlation. This
suggests that the Fe Kα reprocessing scenario is fundamentally different in SGXBs
and in γ Cass analogues.
59
60
Chapter 7. Summary
• The width of FeKα is predominantly below 0.15 keV and can be widely explained
by appealing to line blending, Compton broadening, and moderate Doppler shifts
(∼1000 km/s).
• The curve of growth in SGXBs shows a clear correlation between FeKα EW and
NH , indicating a strong link between the absorbing and the fluorescent matter. From
numerical simulations, the assumption of spherically distributed absorbing matter is
roughly correct for most of the SGXBs.
• The NH values observed in SGXBs are higher than in SFXTs. The disparity is hardly
produced by chance, as shown by a permutation test of the sample, denoting a fundamental physical reason beneath. Systematic differences in the orbital parameters
or different interaction compact object - stellar wind are plausible candidates for explaining such a discrepancy.
• The orbital modulation of NH in IGR J16320-4751 and 4U 1700-37, together with
the aforementioned results, points to the stellar wind as the main contributor to both
continuum absorption and FeKα emission in the case of supergiant donors.
In summary, we present the most comprehensive study of FeKα in HMXBs to date,
complementing previous surveys at high resolution (Torrejón et al., 2010b). We have significantly increased the number of sources and extended the study to all major classes of
massive binaries.
Part III
A comparative analysis of
two supergiant donors in
High-mass X-ray Binaries:
the persistent Vela X-1 and
the transient IGR J17544-2619
As described in Part I, classical and persistent Supergiant X-ray Binaries (SGXBs) and
Supergiant Fast X-ray Transients (SFXTs) are two types of High-mass X-ray Binaries
(HMXBs) that present similar donors but, at the same time, show very different behavior in the X-rays. The reason for this dichotomy of HMXBs is not yet known. As we have
shown in Part II of this thesis, the X-rays absorption is systematically different in SFXTs
and SGXBs. Given that the stellar wind plays a key role in the X-rays absorption, it seems
natural to investigate whether the stellar winds of the donors are also systematically different. Moreover, among the several theoretical explanations that have been proposed in order
to explain the SFXT-SGXB puzzle, some of them invoke specific stellar wind properties of
the donor stars. Only dedicated empiric analysis of the donors’ stellar wind can provide
the required information to accomplish an adequate test of these theories. However, such
analyses are scarce.
To close this gap, in this part of the thesis we perform a comparative analysis of the optical companion in two important systems: IGR J17544-2619 (SFXT) and Vela X-1 (SGXB).
We analyse the spectra of each star in detail and derive their stellar and wind properties. As
a next step, we compare the wind parameters, giving us an excellent chance of recognizing
key differences between donor winds in SFXTs and SGXBs.
We use archival infrared, optical and ultraviolet observations, and analyse them with the
non-LTE Potsdam Wolf-Rayet model atmosphere code. We derive the physical properties
62
of the stars and their stellar winds, accounting for the influence of X-rays on the stellar
winds.
We find that the stellar parameters derived from the analysis generally agree well with
the spectral types of the two donors: O9I (IGR J17544-2619) and B0.5Iae (Vela X-1). The
distance to the sources have been revised and also agrees well with the estimations already
available in the literature. In IGR J17544-2619 we are able to narrow the uncertainty to
d = 3.0 ± 0.2 kpc. From the stellar radius of the donor and its X-ray behavior, the eccentricity of IGR J17544-2619 is constrained to e < 0.25. The derived chemical abundances
point to certain chemical mixing during the lifetime of the donors. An important difference
between the stellar winds of the two stars is their terminal velocities (3∞ = 1500 km/s in
IGR J17544-2619 and 3∞ = 700 km/s in Vela X-1), which has critical consequences on the
X-ray luminosity of these sources.
In summary, the donors of IGR J17544-2619 and Vela X-1 have similar spectral types
as well as similar parameters that physically characterise them and their spectra. In addition, the orbital parameters of the systems are similar too, with a nearly circular orbit and
short orbital period. However, they show moderate differences in their stellar wind velocity
and spin period of their neutron stars that have a strong impact on the X-ray luminosity
of the sources. This specific combination of wind speed and pulsar spin favours an accretion regime with a persistently high luminosity in Vela X-1, while it favours an inhibiting
accretion mechanism in IGR J17544-2619. Our study demonstrates that the wind relative
velocity is critical in the determination of the class of HMXBs hosting a supergiant donor,
given that it may shift the accretion mechanism from direct accretion to propeller regimes
when combined with other parameters.
Part III is structured as follows. In Chapter 8 we describe the set of observations used in
this work. In Chapter 9 we explain the main features of PoWR, the code employed in the
fits. In Chapter 10 we detail the fit process and give the obtained results. In Chapter 11 we
discuss several consequences arising from our results. Finally, in Chapter 12 we enumerate
the conclusions that we find from this work.
The principal content of this part of the thesis has been submitted to the peer-reviewed
journal Astronomy & Astrophysics in October 2015.
Chapter 8
The observations
In this study we used data from International Ultraviolet Explorer (IUE)1 , the fiber-fed
extended range optical spectrograph (FEROS)2 operated at the European Southern Observatory (ESO) in La Silla, Chile; and the infrared (IR) spectrograph SpeX in the NASA
Infrared Telescope Facility (IRTF) in Mauna Kea, Hawaii.
The IUE is provided with two spectrographs (long-wavelength in the range 1850 −
3300 Å and short- wavelength in 1150 − 2000 Å) and four cameras (prime and redundant
camera, for each spectrograph). Each spectrograph can be used with either large aperture (a
slot 10x20 arcsec), or small aperture (a circle 3 arcsec diameter). In addition, each spectrograph has two dispersion modes: high resolution and low resolution. High resolution mode
(∼ 0.2 Å) utilizes an echelle grating plus a cross-disperser. Low resolution mode (∼ 6 Å)
utilizes only the cross-disperser. IUE provides flux calibrated data. This is an important
advantage due to two main reasons: first, we used these observations to fit the spectral energy distribution from the models, as explained below in Section 10.2; and second, we did
not have to normalize the UV spectrum. As we can see in Figures B.3 and 10.10, it is not
straightforward to see the actual flux level of the UV continuum, since this spectral range
is almost completely covered by spectral lines. Therefore, any normalization by visual inspection would lead to significant errors. Instead, we rectified the IUE spectra using the
PoWR model continuum.
FEROS is a spectrograph that yields high resolution echelle spectroscopy (R ∼ 48000)
and high efficiency (∼ 20%) in the optical wavelength range (3600 − 9200 Å) (Kaufer et al.,
1999). SpeX is an infrared spectrograph in the 0.8 − 5.5 µm range. Among the different
modes available in this instrument, we used the 0.8 − 2.4 µm cross-dispersed mode (SXD),
which yields moderate spectral resolution (R ∼ 2000) (Rayner et al., 2003).
In Table 8.1 we present the set of observations of IGR J17544-2619. We used an observation from SpeX taken on August 8, 2004. In the ESO archive there are 14 FEROS
observations of IGR J17544-2619 taken on four different dates during September 2005.
1
2
available at https://archive.stsci.edu/iue/
available at http://archive.eso.org/
65
66
Chapter 8. The observations
Instrument
Phase
Date
MJD
Exposure
(s)
(YYYY-MM-DD)
SpeX
FEROS
0.65
2004-08-15
53232.29
60
0.01
0.01
0.01
0.02
0.61
0.61
0.61
0.62
0.74
0.75
0.76
0.76
0.97
0.98
2005-09-30
2005-09-30
2005-09-30
2005-09-30
2005-09-28
2005-09-28
2005-09-28
2005-09-28
2005-09-09
2005-09-09
2005-09-09
2005-09-09
2005-09-15
2005-09-15
53643.05
53643.03
53643.01
53643.07
53641.08
53641.06
53641.04
53641.10
53622.01
53622.02
53622.10
53622.08
53628.06
53628.08
1470
1470
1470
1470
1470
1470
1470
1470
1470
1470
1470
1470
1470
1470
Table 8.1: Table of observations of IGR J17544-2619. We used T 90 = T 0 = T φ=0 =
55924.271 (MJD) and orbital period Porb = 4.9272 d (Drave et al., 2014).
There are not IUE available public observations of IGR J17544-2619.
In Table 8.2 we present the set of observations of Vela X-1. In the ESO archive there are
six consecutive FEROS observations of 700s taken on April 22, 2006. For the IUE data, we
used the high dispersion and large aperture observations using the short-wavelength spectrograph (1150-2000 Å) and the prime camera (SWP). There are 49 observations in the
public database of the IUE following these criteria.
For each instrument, we averaged over all the available observations taking into account
the exposure time in order to improve the signal-to-noise ratio. We did not take the variability of the UV spectral lines depending on the orbital phase into account, that has been
reported for Vela X-1 (Sadakane et al., 1985a). The variability consists on the presence of an
extra absorption component in several spectral lines, specially ones belonging to Al iii and
Fe iii , mainly at phases φ > 0.5. This variability must be taken into account to interpret
the full picture of the stellar wind of Vela X-1. However, in this work, we prioritized a
signal-to-noise ratio as high as possible over fitting a number of phase dependent spectra
with significantly lower signal-to-noise. This permits us to estimate the stellar parameters
of Vela X-1 more accurately, while not affecting any of the conclusions derived in this work,
as we have carefully examined.
67
Instrument
Phase
Date
MJD
Exposure
(s)
(YYYY-MM-DD)
SWP
FEROS
0.05
0.07
0.08
0.09
0.09
0.10
0.10
0.10
0.14
0.17
0.22
0.28
0.28
0.29
0.29
0.33
0.40
0.40
0.41
0.45
0.46
0.46
0.49
0.50
0.51
0.52
0.52
0.53
0.53
0.55
0.60
0.61
0.66
0.71
0.73
0.74
0.75
0.76
0.77
0.77
0.79
0.84
0.85
0.90
0.97
0.97
0.97
0.98
0.99
1978-05-05
1984-02-19
1985-05-03
1985-05-03
1985-05-03
1985-05-03
1985-05-03
1993-11-08
1978-12-07
1992-11-06
1993-11-09
1983-01-22
1992-11-07
1983-01-22
1984-02-21
1993-11-10
1984-02-22
1988-02-22
1992-11-08
1978-04-30
1982-12-19
1993-11-11
1985-05-07
1985-05-07
1985-05-07
1988-02-23
1988-03-12
1978-12-20
1983-01-07
1993-11-03
1978-12-02
1983-01-16
1993-11-04
1978-12-03
1984-02-16
1983-01-09
1985-04-21
1985-04-21
1979-03-21
1985-04-21
1993-11-05
1984-02-17
1978-07-23
1993-11-06
1983-01-11
1983-01-20
1984-02-18
1983-01-20
1993-11-07
43633.62
45749.32
46188.67
46188.75
46188.81
46188.86
46188.92
49299.55
43849.51
48932.56
49300.55
45356.80
48933.57
45356.91
45751.31
49301.55
45752.36
47213.55
48934.72
43628.21
45322.52
49302.71
46192.36
46192.47
46192.58
47214.54
47232.54
43862.03
45341.09
49294.56
43844.71
45350.77
49295.55
43845.69
45746.31
45343.01
46176.77
46176.86
43953.77
46176.99
49296.71
45747.32
43712.49
49297.73
45345.10
45354.07
45748.49
45354.13
49298.54
9000
8280
4500
4500
3300
1020
6000
8400
8400
10800
8100
10800
9600
4500
9000
9000
9000
8460
9900
10800
9000
8400
7200
7200
7200
8460
7826
7800
10800
6000
5400
10800
8400
8400
9000
10800
7200
4500
9000
6900
7500
9000
7500
6600
10800
5400
7500
3300
9600
0.68
0.68
0.68
0.68
0.68
0.68
2005-04-22
2005-04-22
2005-04-22
2005-04-22
2005-04-22
2005-04-22
53482.05
53482.06
53482.07
53482.07
53482.09
53482.10
700
700
700
700
700
700
Table 8.2: Table of observations of Vela X-1. We used T 90 = T 0 = T φ=0 =
52974.001 (MJD) and orbital period Porb = 8.964357 d (Kreykenbohm et al., 2008).
Chapter 9
The PoWR code
PoWR computes models of hot stellar atmospheres assuming spherical symmetry and
stationary outflow. The non-LTE population numbers are calculated using the equations of
statistical equilibrium and radiative transfer in the co-moving frame. Since these equations
are coupled, the solution is iteratively found. Once convergence is reached, the synthetic
spectrum is calculated integrating along the emergent radiation rays. The main features of
the code have been described by Gräfener et al. (2002) and Hamann and Gräfener (2003).
The basic input parameters in PoWR are the following: stellar temperature (T ⋆ ), luminosity (L⋆ ), mass-loss rate ( Ṁ), surface gravity (g⋆ ) and chemical abundances. The chemical elements taken into account are detailed in Table 10.2. The stellar radius (R⋆ ) follows
from T ⋆ and L⋆ using the Stefan-Boltzmann law: L⋆ = 4πσT ⋆4 R2⋆ , where σ is the StefanBoltzmann constant. We note that, in PoWR, R⋆ refers to the layer where the Rosseland
continuum optical depth τmax = 20, and not to the definition of stellar radius (or photospheric radius), where τRoss = 2/3. Nevertheless, we will give the stellar parameters in the
next sections referring to both τmax = 20 and the τmax = 2/3, in order to avoid any confusion (e.g., we will use R⋆ for the radius at τmax = 20 and R2/3 for the radius at τmax = 2/3).
The surface gravity g⋆ and R⋆ imply the stellar mass (M⋆ ) via g⋆ = GM⋆ R−2
⋆ . Instead
of g⋆ , one may specify the effective surface gravity geff , which accurately accounts for the
outward force exerted by the radiation field, as thoroughly described by Sander et al. (2015).
The density stratification in the stellar atmosphere, ρ(r), is calculated from the continuity equation Ṁ = 4πr2 3(r)ρ(r), given Ṁ and the radial velocity stratification 3(r). For
3(r), PoWR distinguish between two different regimes: the quasi-hydrostatic domain and
the wind domain. A detailed description of the quasi-hydrostatic domain can be found in
Sander et al. (2015). In the wind domain, the β-law is adopted (Castor et al., 1975b):
r0 β
3(r) = 3∞ 1 −
(9.1)
r
where 3∞ is the terminal velocity of the wind, r0 ≈ R⋆ (depending on the precise location of
the connection point) and β is an input parameter typically ranging between β = 0.6 − 2.0
(Puls et al., 2008). The connection point is chosen in order to ensure a smooth transition
between the two domains. The temperature stratification is calculated from the condition of
69
70
Chapter 9. The PoWR code
radiative equilibrium (Hamann and Gräfener, 2003).
The code also permits to account for density inhomogeneities and additional X-rays from
a spherically-symmetric, shock heated plasma. Density inhomogeneities are described in
PoWR by means of an optional radial-dependent input parameter: the density contrast
D(r) = ρcl /ρ̄, where ρcl is the density of the clumped medium and ρ̄ is the average density. The inter-clump medium is assumed to be empty. During the analysis, D(r) is assumed
to grow from D(rsonic ) = 1 (smooth plasma) to a maximum value D, which is reached at
the layer where the stellar wind velocity is fmax × 3∞ . D is a free parameters derived in the
analysis. fmax has a modest influence on the spectra. We assumed fmax ∼ 0.6 on the basis of
this moderate effect. The X-rays are described using three parameters: the X-ray temperature T X , the filling factor XF (i.e. the ratio between shocked to unshocked plasma), and the
onset radius RX , as described in Baum et al. (1992). In this work, we assumed T X = 107 K,
RX = 1.2 R⋆ and XF = 0.05. The main influence of X-rays in the model is via Auger ionization, which is responsible for the appearance of resonance lines belonging to high ions such
as Nv and Ovi in the spectra of O stars (Cassinelli and Olson, 1979; Oskinova et al., 2011).
Any changes in these parameters barely affect the spectrum, as long as they they produce a
similar X-ray luminosity.
During the iterative calculation of the population numbers, the spectral lines are taken to
be Gaussian with a constant Doppler width of 3Dop = 40 km/s; the effect of 3Dop on the spectrum is negligible for most lines (see discussion by Shenar et al., 2015). During the formal
integration, the line profiles include natural broadening, pressure broadening, and Doppler
broadening. The Doppler width is decomposed per element to a depth dependent thermal
motion and a microturbulent velocity ξ(r). The photospheric microturbulence, ξph , is derived in the analysis, and beyond the photosphere we assumed that it grows from ξ = ξph to
ξ = 100 km/s at the layer where the stellar wind velocity is 500 km/s. Rotational broadening
is simulated via convolution with rotational profiles with a width corresponding to the projected rotational velocity 3rot sin i (denoted by 3rot hereafter for simplicity), except for important wind lines, for which the convolution is no longer valid (see e.g. Hillier et al., 2012),
and where an explicit angle-integration would be required (as described by Shenar et al.,
2014). The so-called macroturbulence 3mac is accounted for by convolving the spectra with
so-called Radial-Tangential profiles (Gray, 1975; Simón-Díaz and Herrero, 2007).
Chapter 10
The fitting procedure
We used the PoWR code to calculate synthetic spectra and a Spectral Energy Distribution (SED) which best match the observations. The large number of free parameters,
together with the long computation time for each model, do not permit the construction of
a grid of models that covers the full parameter space. Instead, we attempted to identify
the best-fitting model by visual inspection and systematic variation of the parameters. As
an initial step, we calculate models using typical parameters of late O / early B stars. We
then use specific spectral lines for each parameter as a guideline for the fit. Generally, the
effective gravity geff is derived from the pressure-broadened wings of the Balmer lines and
He ii lines. The temperature T ⋆ is derived based on line ratios belonging to ions of the same
element. The mass-loss rate Ṁ, 3∞ and D are derived from "wind-lines", with D adjusted
so that a simultaneous fit is obtained for both resonance lines (which scale as ρ) as not as
recombination lines such as Hα (which scale as ρ2 ). The luminosity L and the reddening
E B−V are derived by fitting the SED to photometry and flux-calibrated spectra. We apply the
reddening law by Fitzpatrick (1999). Abundances are estimated from the overall strengths
of the spectral lines. The photospheric microturbulence ξph is found from the strength and
shape of helium lines. Finally, the parameters β, 3rot and 3mac are adopted on the basis of
the shape and depth of the spectral lines, together with previous estimations found in the
literature, when available. Upon adjusting the model, the whole spectral domain was examined to iteratively improve the fit. Overall, we managed to find models which satisfactorily
reproduce the observed spectra and SEDs of the donors of the two systems analysed here.
We show the complete fits in Appendix B. The details about the fitting procedure for
the two objects are given in the following subsections. The obtained parameters are summarised in Table 10.1 and the chemical abundances in Table 10.2. The parameters that do
not include an error estimation in the tables are adopted following the above mentioned criteria.
Even though the optical companion in Vela X-1 is usually known as HD 77581, for the
sake of simplicity we will refer to the donors with the name that is used for the X-rays
sources, namely, IGR J17544-2619 and Vela X-1. Depending on the context, the reader
should easily recognize whether it is the donor or the X-ray source which is being referred
71
72
Chapter 10. The fitting procedure
Parameters
J17544-2619
Vela X-1
log (L/L⊙ )
M⋆ /M⊙
R⋆ /R⊙
R2/3 /R⋆
T ⋆ (kK)
T 2/3 (kK)
log (g⋆ (cgs))
log (geff (cgs))
log (geff 2/3 (cgs))
3∞ (km/s)
3esc (km/s)
log ( Ṁ/(M⊙ /yr))
D
ξph (km/s)
β
3mac (km/s)
3rot (km/s)
EB-V
RV
d (kpc)
5.4 ± 0.1
25.9 ± 2.0b
20+4
−3
1.04
29.0 ± 1.0
28.5 ± 1.0
3.25 ± 0.20
2.80 ± 0.20
2.77 ± 0.20
1500 ± 200
618 ± 75
−5.8 ± 0.2
4
25 ± 10
0.8
60
160
2.14 ± 0.10
2.9
3.0 ± 0.2
5.5 ± 0.1
21.5 ± 4
28.4a
1.09
25.5 ± 1.0
24.4 ± 1.0
2.86 ± 0.10
2.35 ± 0.10
2.27 ± 0.10
700+200
−100
436 ± 65
−6.2 ± 0.2
11 ± 5
30 ± 10
1.0
80
56c
0.77 ± 0.05
3.1
2.0 ± 0.2
Table 10.1: Stellar parameters obtained from the best fit. (a) Joss and Rappaport (1984).
(b) Pellizza et al. (2006a). (c) Fraser et al. (2010)
to.
10.1 IGR J17544-2619
IGR J17544-2619 was first detected on September 2003 with the IBIS/ISGRI detector on
board INTEGRAL (Sunyaev et al., 2003). It is located in the direction of the galactic center,
at galactic coordinates l = 3.24◦ , b = −0.34◦ . The orbital period is ∼4.9d (Clark et al.,
2009). According to Chandra observations, the compact object is a neutron star (in’t Zand,
2005). Pellizza et al. (2006a) used optical and NIR observations in order to classify the optical companion as a O9Ib. XMM-Newton and Suzaku observations showed that the system
exhibits a high dynamic range in its X-ray variability, changing X-ray flux by four orders of
magnitude (González-Riestra et al., 2004; Rampy et al., 2009).
Nowadays, the spin period Pspin of the hypothetical neutron star in IGR J17544-2619
is a matter of debate, given the results arising from observations taken at different times,
different luminosities and different instruments. Drave et al. (2012) analysed RXTE data of
the source at intermediate X-ray luminosity (∼ 1033−34 erg/s), and reported the detection of
an X-ray pulsation with Pspin = 71.49s at a statistical significance of 4.37σ. Romano et al.
10.1. IGR J17544-2619
73
IGR J17544-2619
Vela X-1
Quemical
Element
Mass Fraction
Rel.
Ab.
Mass Fraction
Rel.
Ab.
H
He
C
N
O
Si
S
P
Al
Mg
Fea
(6.2 ± 0.5)E − 01
(3.7 ± 0.5)E − 01
(5.0 ± 3.0)E − 04
(2.2 ± 0.6)E − 03
(6.0 ± 2.0)E − 03
(7.3 ± 2.0)E − 04
5.0E − 04
6.4E − 06
5.8E − 05
7.0E − 04
1.4E − 03
0.85
1.47
0.17
2.58
0.76
1.00
1.00
1.00
1.00
1.00
1.00
(6.5 ± 0.5)E − 01
(3.4 ± 0.5)E − 01
(5.0 ± 3.0)E − 04
(1.8 ± 0.6)E − 03
(7.0 ± 0.2)E − 03
(5.5 ± 2.0)E − 04
5.0E − 04
6.4E − 06
7.0E − 05
7.0E − 04
1.4E − 03
0.89
1.35
0.17
2.11
0.88
0.75
1.00
1.00
1.00
1.00
1.00
Table 10.2: Chemical abundances derived from the best fit, in mass fraction and relative to
solar abundances from Asplund et al. (2009). (a) The notation of Fe actually stands for a
generic atom including iron group elements Sc, Ti, V, Cr, Mn, Co and Ni. For more details
see Gräfener et al. (2002).
(2015) inspected Swift observations of the source experiencing an extraordinarily bright
outburst (peak luminosity ∼ 1038 erg/s), and reported the detection of X-ray pulsations with
Pspin = 11.60s at a statistical significance of about 4σ too. However, these results contrast
with the analyses of XMM-Newton and NuSTAR observations performed by Drave et al.
(2014) and Bhalerao et al. (2015) respectively. These authors do not find any evidence of
pulsations on time scales of 1-2000s.
We have adjusted T ⋆ of IGR J17544-2619 using different ions, mainly He i-He ii and
Si iii-Si iv. In Figure 10.1 we show an example of four helium lines of which the best-fit
model provides a good description. Higher (lower) temperatures yield more (less) absorption than observed in the He ii lines. We have used other lines of helium, silicon, nitrogen
and oxygen. The vast majority of them are well described by the best-fit model, within
the errors. The obtained effective temperature is compatible with the donor’s spectral class
O9 Ib (Martins et al., 2005).
The effective gravity geff was found using the hydrogen Balmer lines Hγ and Hδ. We did
not use Hβ and Hα because these lines are notably affected by the stellar wind. Figure 10.2
shows a comparison of the observations with the best-fitting model for these two Balmer
lines. We show that the observations are compatible with a relatively wide range of values,
as also reflected in the errors given in Table 10.1.
The distance to IGR J17544-2619 is not well known, with an estimate of 2-4 kpc Pellizza et al.
(2006a), based on the extinction and the calibration of the absolute magnitude for O9Ib
74
Chapter 10. The fitting procedure
1.00
Normalized flux
Normalized flux
HeI
HeI
1.00
0.95
0.95
0.90
0.90
0.85
0.85
4910
4920
o
λ/A
4930
6670
6690
1.00
Normalized flux
HeII
HeII
1.00
Normalized flux
6680
o
λ/A
0.95
0.90
0.98
0.85
5400
5410
o
λ/A
5420
0.96
6400
6410
o
λ/A
Figure 10.1: Example of four helium lines in IGR J17544-2619, used for the estimation
of T ⋆ . We show the observation (solid blue line), the best-fit model (red dashed line), a
model with lower temperature of T ⋆ = 28 kK (green dashed line), and a model with higher
temperature of T ⋆ = 30 kK (pink dashed line).
10.1. IGR J17544-2619
75
1.0
0.9
0.8
0.7
0.6
Hγ
Normalized flux
Normalized flux
Hδ
1.0
0.9
0.8
0.7
4100
o
λ/A
0.6
4330
4340
o
λ/A
4350
Figure 10.2: Hγ and Hδ in IGR J17544-2619, used for the surface gravity estimation. We
show the observation (solid blue line), the best-fit model (red dashed line) and a model with
larger effective gravity of log (geff ) = 3.0 in cgs units (green dashed line).
76
Chapter 10. The fitting procedure
Figure 10.3: Extinction curve in the galactic direction of IGR J17544-2619. The shaded area
reflects the error in the distance estimation from the errors of estimation of the extinction
and the errors in the calculation of the extinction curve.
stars. In this work we improve this estimation. As a first step, we fitted the SED to
photometry from the 2MASS catalogue (Cutri et al., 2003), Zacharias et al. (2012) and
Rahoui and Chaty (2008) assuming the distance to be 3 kpc. Then, we derived initial values
for the luminosity of the donor and the reddening to the system.
As a second step, in order to provide more constrains on the distance, we employed a
method based on the well constrained luminosity of Red Clump Giant stars (RCG). These
stars can be isolated in a NIR colour-magnitude diagram and permit the estimation of the
interstellar extinction along the line of sight (López-Corredoira et al., 2002). Due to their
narrow luminosity function, the apparent magnitude of RCGs provides an estimation of the
distance. Then, given a certain line of sight, a diagram of the extinction versus the distance
can be derived (for more details see González-Fernández et al. (2014)). For IGR J175442619 we employed the derived E J−K from the SED fit to obtain an estimate of the distance.
We note that this method is only applicable to stars in the direction of the galactic center
like IGR J17544-2619, where the medium is more homogeneous and the density of RCGs
is higher. Using this method, we obtain a distance of 3.0 ± 0.2 kpc (Figure 10.3). Revised
values for the luminosity and reddening are then derived. The final results of the SED fit
are shown in Figure 10.4.
10.1. IGR J17544-2619
77
-12
IGR J17544-2619
8.310
12.94 12.10
o
log f λ [erg s -1 cm -2 A -1 ]
8.791
-14
8.018
Rahoui 2008
14.71
B
V
R
J
H
K
PAH1
PAH2
7.619
-16
Zacharias 2012
2MASS
7.933
3.6
3.8
4.0
4.2
4.4
o
log λ / A
4.6
4.8
5.0
5.2
Figure 10.4: Fit of the SED of IGR J17544-2619. In red we plot the best-fit model. We
indicate the photometry values for each band in blue. The employed references are cited.
The values of extinction, distance and luminosity are shown in Table 10.1.
From the luminosity and temperature we derive R⋆ , which provides an upper limit to the
eccentricity of the system. For the lower limit R⋆ = 17 R⊙ , we find e < 0.25. For higher
eccentricities, periodic Roche-lobe overflow is expected from the orbital solution of the system (Clark et al., 2009), at odds with the X-ray behavior of the source. Given the radius of
the source and the derived surface gravity, we find M⋆ = 25.9 M⊙ . This value matches very
well with the estimation of M⋆ = 25 − 28 M⊙ done by Pellizza et al. (2006a) based on the
mass calibration with its spectral type.
The terminal velocity of the stellar wind 3∞ was derived using the P-Cygni profile of
He i λ10833 Å (see Figure 10.5). The blue wing in He i λ10833 Å is a very good indicator due to its strong sensitivity to 3∞ . It is reasonably well fitted when assuming
3∞ ≃ 1500 km/s.
Ṁ and D were simultaneously adjusted by means of Hα and the P-Cygni profile of
HeI λ10833 Å.√Provided that the strength of emission in these recombination spectral lines
varies with Ṁ D (Gräfener et al., 2002), we cannot estimate Ṁ and D independently using
these lines. As it is shown in Figure 10.6, we were not able to fit all the lines at the same
time. The best-fit model provides an acceptable description of Hα, but yields insufficient
emission for HeI λ10833 Å. We choose the best description of Hα as the best-fit because it
provides a better fit to the overall spectrum. We note that the optical and infrared spectra
were not taken at the same time, and therefore any kind of variability in the lines might
produce a disagreement. However, Hα does not show such a large variability within the
observations we have analysed (see Figure 10.7).
Without available resonance scattering lines in the observations at hand, we cannot compare P-Cygni lines with recombination lines to deduce the clumping factor D. However,
our calculations show that changing Ṁ dramatically
affects the absorption spectrum in a
√
fashion which is not related to the product Ṁ D. An example is shown in Figure 10.8,
where we show three models calculated with different values of Ṁ and D, but with a fixed
78
Chapter 10. The fitting procedure
1.15
1.05
1.00
HeI
Normalized flux
1.10
0.95
10800
10850
o
λ/A
10900
Figure 10.5: He i at λ 10833 Å, used for the estimation of 3∞ in IGR J17544-2619 fitting
the blue wing of the P-Cygni profile. We show the observation (solid blue line), the best-fit
model (red dashed line), a model with 3∞ = 1300 km/s (green dashed line), and a model
with 3∞ = 1700 km/s (pink dashed line).
10.1. IGR J17544-2619
79
1.25
1.15
1.20
Hα
1.10
Normalized flux
1.10
1.05
1.00
1.05
1.00
HeI
Normalized flux
1.15
0.95
0.90
0.85
6525
0.95
6550
6575
o
λ/A
6600
10800
10850
o
λ/A
10900
√
Figure 10.6: Hα and He i λ10833Å lines for the estimation of Ṁ D in IGR J17544-2619.
We show the observation
(solid blue line), the best-fit model (red dashed line), a model
√
with
√ higher Ṁ D (1.35 times the best-fit value, green dashed line), and a model with lower
Ṁ D (0.8 times the best-fit value, pink dashed line).
80
Chapter 10. The fitting procedure
Hα
1.15
1.10
Normalized flux
1.05
1.00
0.95
0.90
0.85
6525
6550
6575
6600
o
λ/A
Figure 10.7: Hα in IGR J17544-2619 at different orbital phases: φ ≃ 0.01, 0.61, 0.75, 0.97
(blue, green, pink and turquoise solid lines respectively).
√
product Ṁ D. Evidently, while the emission exhibited by the wings of H-α (shown in
Figure 10.8) is similar in all models, the absorption lines are strongly affected in a nontrivial manner. The reason for this unexpected behaviour is that many of the strong lines
in the spectrum (e.g. the Balmer series) are formed significantly beyond the photosphere
(τRoss ≈ 2/3), where the mass-loss rate already strongly affects the density stratification via
the continuity equation. Exploiting this effect, we find that D ≈ 4 provides the best results
for the overall spectrum. However, we warn that further observations are needed to better
constrain the clumping in this star. Nevertheless, we note that our final conclusions do not
strongly depend on the clumping factor and the implied mass-loss rate, as will be discussed
in Chapter 11.
The chemical composition was estimated from unblended spectral lines for He, C, N, O
and Si. The rest of the considered element abundances (see Table 10.2) were assumed solar
following Asplund et al. (2009). The fit yielded moderate overabundance of He and N, together with underabundance of C and O. In all, there are indications of chemical evolution
in the outer layers of the stellar atmosphere.
The photospheric microturbulent velocity (ξph ) was adjusted using He i and SiIV lines.
A higher ξph induce stronger absorption in several spectral lines, as shown in Figure 10.9.
The 3rot and 3mac were roughly estimated using the width of the He lines. The derived projected rotational velocity is around 0.3 times the critical rotation velocity (3crit =
10.1. IGR J17544-2619
81
Hα
1.15
Hδ
1.0
Normalized flux
Normalized flux
1.10
1.05
1.00
0.95
0.9
0.8
0.7
0.90
6550
6575
o
λ/A
6600
0.6
4080
4090
4100 4110
o
λ/A
4120
Figure 10.8: Hα and Hδ for the estimation of Ṁ in IGR J17544-2619. We show the
observation (solid blue line), the best-fit model ( Ṁ = 10−5.8 M⊙ /yr, D = 4, red dashed
line), a model with higher Ṁ ( Ṁ = 10−5.5 M⊙ /yr, D = 1, green dashed line), and a model
with lower Ṁ ( Ṁ = 10−5.9 M⊙ /yr, D = 8, pink√dashed line). Different Ṁ values do not yield
different Hα wings as long as the product Ṁ D remains constant. However, we observed
that other important lines like Hδ are sensitive these variations.
82
Chapter 10. The fitting procedure
1.00
HeI
Normalized flux
0.95
Normalized flux
HeI
1.00
0.90
0.85
0.95
0.90
0.80
0.85
0.75
0.70
5870
5875
o
λ/A
5880
0.80
7060
7065
o
λ/A
7070
Figure 10.9: Example of two He i lines in IGR J17544-2619, used in the ξph estimation. As
usual, the observation is plotted in solid blue line, and the best-fit model in red dashed line.
Models with ξph = 15, 35 km/s are also presented (green and pink dashed lines respectively).
10.2. Vela X-1
83
q
G M⋆ R−1
⋆ ). This high rotational velocity may favour the chemical mixing, in line with the
abundances derived in the fit.
To summarise, our NLTE analysis of optical and near IR spectra of IGR J17544-2619
showed that the optical O9I-type companion in this source is not-peculiar and has stellar
and wind parameters that are similar to other stars of the same spectral type, e.g. δ Ori
(Shenar et al., 2015).
10.2 Vela X-1
Vela X-1 is one of the most studied HMXBs, since it is a bright source discovered in the
early ages of the X-ray astronomy (Chodil et al., 1967). It is located at galactic coordinates
l = 263.06◦ , b = 3.93◦ . The distance was estimated to be 1.9 ± 0.2 kpc by (Sadakane et al.,
1985a). The system has a moderate eccentricity of e= 0.09 (Bildsten et al., 1997), and
orbital period Porb = 8.96 days (Kreykenbohm et al., 2008). The compact object is a neutron star that pulsates with Pspin = 283s (McClintock et al., 1976). The optical companion
HD 77581 (B0.5Iae) was identified by Vidal et al. (1973).
It is very likely that the wind of Vela X-1 is disturbed by the X-ray source. The photoionization produced close to the photosphere due to the intense X-ray luminosity might
hinder the acceleration of the wind and generate a structure known as photoionization wake
(Blondin et al., 1990; Krtička et al., 2015). This structure appears in the UV spectra as an
additional absorption component at phases larger than φ ∼ 0.5 (Kaper et al., 1994). In addition, the hard X-rays light curves of the source in near-to-eclipse phases show asymmetries
between ingress and egress, that have been interpreted as caused by the existence of this
type of structure trailing the neutron star (Feldmeier et al., 1996). Moreover, a density enhancement in the line of sight during the second half of the orbit is also observed in the
X-rays absorption, although the amount of absorbing material is highly variable from one
orbit to another.
We derived T ⋆ following the same procedure that we used for IGR J17544-2619. The
obtained T 2/3 is similar to previous estimations: Sadakane et al. (1985a) used the equivalent width (EW) of photospheric lines to estimate the effective temperature T 2/3 = 25000 K;
Fraser et al. (2010) used the TLUSTY code to estimate T 2/3 = 26500 K.
For the fit of the SED, we used photometry from the 2MASS catalogue (Cutri et al., 2003)
and the Stellar Photometry in Johnson’s 11-color system (Ducati, 2002), together with the
IUE observations. We made an estimation of the reddening, distance and RV ≡ A(V)/E B−V
by means of the SED fit. Then, we used the estimation of the stellar radius R2/3 = 31 R⊙
from Joss and Rappaport (1984), and T 2/3 from the successive fits, in order to derive the
luminosity (and the distance estimation) from the Stefan-Boltzmann law. Given that the
obtained T2/3 is very similar to previous estimations, the derived distance of 2.0 ± 0.2 is
almost equal to the value d=1.9 kpc given by Sadakane et al. (1985a). We show the results
of the SED analysis in Figure 10.10.
84
Chapter 10. The fitting procedure
Vela X-1
6.85 7.37
6.87
6.31
o
log f λ [erg s -1 cm -2 A -1 ]
6.05
5.833
-12
5.705
5.596
U
IUE
-14
3.0
3.2
B
V
R
I
J
Ducati 2002
3.4
3.6
o
3.8
H
K
2MASS
4.0
4.2
4.4
log λ / A
Figure 10.10: Fit of the SED of Vela X-1. In red we plot the best fit model with the
spectral lines in the domain where we have done the spectral analysis, and the continuum
where we have available photometry (marked in blue). We cite the references used for
the photometry. Note that the true continuum in the UV range do not correspond with the
apparent continuum from the observation, due to the number of spectral lines in this domain.
The employed values of extinction, distance and luminosity are shown in Table 10.1.
The estimation of geff was especially delicate in Vela X-1 because of its very low geff . A
higher value beyond the error given in Table 10.1 has a strong effect in the overall spectrum
and hinders a satisfying fit. The derived value enables a good fit, and it is in agreement with
previous estimations (Fraser et al., 2010).
We used UV resonance lines to find 3∞ . In Figure 10.11 we show the Si iv resonance
lines λ 1394, 1403 Å, where the effect of 3∞ is very clear. Models with higher terminal
velocities induce a shift towards the blue part of these spectral lines. The best description
of the observations is achieved for 3∞ = 700 km/s. This value is in agreement with the
estimation of van Loon et al. (2001): 3∞ = 600 km/s; and not too far from Watanabe et al.
(2006b), who estimated 3∞ = 1100 km/s using Chandra X-rays observations.
In contrast, it is in disagreement with the estimation of Dupree et al. (1980), namely
3∞ = 1700 km/s. These authors used a subset of the IUE observations used in this work,
and considered the UV resonance lines Si iv and C iv in the X-ray eclipse phases to make
their estimation. We have revisited our 3∞ estimation using only observations taken at
orbital phases φ = 0.9 − 0.1, in order to be able to directly compare to Dupree et al. (1980).
In Figure 10.12, we show the Si iv and C iv lines, as observed in the total averaged spectrum
and the spectrum averaging over φ = 0.9 − 0.1. C iv is almost the same in both cases. Then,
the disagreement in the estimates of 3∞ does not come from orbital phase variations but
from the overlook of the impact of the X-rays in the stellar wind by Dupree et al. (1980).
As we can see in Figure 10.12, when we introduce X-rays in the models we are able to
reproduce C iv without need for a high velocity, due to the significant enhancement of the
population of C iv in the wind. We note that the X-ray radiation we are introducing in the
models is an intrinsic radiation of the donor wind that is presumably produced in the shocks
SiIV
85
SiIV
10.2. Vela X-1
Normalized flux
1.0
0.5
0.0
1380
1390
1400
1410
o
λ/A
Figure 10.11: Si iv lines used for the estimation of 3∞ in Vela X-1. It is showed the observation (blue solid line), the best-fit model (red dashed line) and a model with 3∞ = 900 km/s
(green dashed line).
SiIV
Chapter 10. The fitting procedure
1.0
CIV
Normalized flux
Normalized flux
1.0
0.5
0.0
1390
1400
o
λ/A
1410
CIV
SiIV
86
0.5
0.0
1540
1545
o
λ/A
1550
Figure 10.12: Si iv and C iv resonance lines. We plot the total averaged spectrum (blue solid
line) and the averaged spectrum over orbital phases φ = 0.9 − 0.1 (pink solid line). We also
plot the best-fit model (3∞ = 700, red dashed line), and a model with 3∞ = 1400 km/s (green
dashed line).
within the stellar wind itself (Krtička et al., 2009, e.g.). This radiation is not coming from
the neutron star, since the effects are also noticeable at eclipsing phases. The impact of
the X-rays coming from the neutron star is a different and complex issue, and it has been
already studied by other authors (Watanabe et al., 2006b). Regarding the Si iv resonance
lines, in Figure 10.12 we show that high stellar wind velocities as derived by Dupree et al.
(1980) do not fit, neither using the total averaged spectrum, neither using the eclipsing
phases spectrum.
√
The value Ṁ D was estimated using Hα (see Figure 10.13). We did not find a good fit
of the blue part of the line, observed in absorption, but our model properly fits the emission
in the red part of the spectral line. Unfortunately, we do not have more optical observations
covering further orbital phases in order to check whether Hα is variable. Nevertheless,
previous studies of similar sources demonstrate that this might be the case: González-Galán
(2015) reported the variability of Hα in the very similar B0Iaep optical companion in the
SGXB system XTE J1855-026. Moreover, the shape of Hα in XTE J1855-026 at φ = 0 (see
Figure 5.12 in González-Galán 2015), when the neutron star is hidden behind the optical
counterpart, is strongly reminiscent of the shape that our model reproduces in Figure 10.13.
Hence, the relative disagreement between our best-fit model and our observation of Vela X1 (taken at φ = 0.68), might be produced by some kind of interaction of the neutron star
87
Hα
10.2. Vela X-1
Normalized flux
1.2
1.0
0.8
6550
6560
o
λ/A
6570
6580
√
Figure 10.13: Hα line for the estimation of Ṁ D in Vela X-1. We can see the observation
√
(blue solid line), the best-fit model (red dashed line), a model with 0.8 times
the Ṁ D
√
value of the best fit (green dashed line), and a model with 1.2 times the Ṁ D value of the
best fit (pink dashed line).
with the donor and/or the stellar wind, which is not possible to model using the assumption
of spherical symmetry that PoWR employs. This disagreement might be related to similar
features observed in other strong lines, as further discussed in Section 11.3.
We derived Ṁ and D from the Al iii resonance lines λ1855 and λ1863Å. As we can see
in Figure 10.14, the variation of Ṁ (and consequently D) directly affect these lines. Higher
(lower) Ṁ enhances (reduces) the density of the stellar wind, producing too strong (weak)
absorption.
Unfortunately, other resonance lines available in the spectrum (N v, C iv and Si iv) are
saturated in the models within a reasonable range of parameters around the best-fit, and consequently are not suitable for the Ṁ diagnosis. Interestingly, in contrast to the models, the
N v and Si iv resonance lines are slightly desaturated in the observations (see Figure 10.15).
The origin of this phenomenon might be related to the presence of optically thick clumps
(macroclumping), which directly affects the mass-loss rate estimations (Oskinova et al.,
88
Chapter 10. The fitting procedure
AlIII
AlIII
1.2
Normalized flux
1.0
0.8
0.6
0.4
0.2
0.0
1850
1855
o
λ/A
1860
1865
Figure 10.14: Al iii resonance lines λ1855 and λ1863Å, employed for the Ṁ estimation in
Vela X-1. We show the observation (solid blue line), the best-fit model ( Ṁ = 10−6.2 M⊙ /yr,
D = 11, red dashed line), a model with higher Ṁ ( Ṁ = 10−5.8 M⊙ /yr, D = 2, green dashed
line), and a model with lower Ṁ ( Ṁ = 10−6.3 M⊙ /yr, D = 20, pink dashed line).
2007; Šurlan et al., 2012). Undoubtedly, its study deserves further investigation, which is
beyond the scope of this work.
Based on the X-ray data analysis, Manousakis and Walter (2015) have suggested that the
velocity law with the parameter β = 0.5 fits better with the X-ray light curve of the system
in near-to-eclipse phases. However, a satisfying fit is not√possible when we assume β = 0.5.
We have tried models using β = 0.5 and adapting Ṁ D in order to fit Hα. However,
as shown in Figure 10.16, Hα in our observation is not compatible with β = 0.5. As we
mentioned above, Hα might suffer from important variability along the orbit. Moreover,
the X-ray irradiation from the neutron star might produce variations in the stellar wind. In
our opinion, this might be the cause of the apparent disagreement between the conclusions
extracted from the X-rays and the optical wavebands.
The chemical composition and was estimated following the same approach as it was done
for IGR J17544-2619. Interestingly, we found again indications of chemical evolution in
SiIV
89
SiIV
10.2. Vela X-1
1.5
0.5
NV
NV
Normalized flux
Normalized flux
1.0
1.0
0.5
0.0
1390
1400
λ/A
o
1410
0.0
1235
1240
o
λ/A
1245
Figure 10.15: Si iv and N v resonance lines. While the observations show slight desaturation, all the models within a reasonable parameter space around the best-fit model produce
saturated lines.
Chapter 10. The fitting procedure
Hα
90
Normalized flux
1.2
1.0
0.8
6550
6560
o
λ/A
6570
6580
Figure 10.16: Hα used the estimation of the parameter β in Vela X-1. We present the
observation (solid blue line), the best-fit model (β = 1.0, red dashed line), and a model with
β = 0.5 (green dashed line), as proposed by Manousakis and Walter (2015).
10.2. Vela X-1
91
1.0
OII
Normalized flux
OII
OII
OII
Normalized flux
OII
1.0
0.9
4065
4070
4075
o
λ/A
4080
0.9
4315
4320
o
λ/A
4325
Figure 10.17: Example of unblended lines in the spectrum of Vela X-1. In red we plot a
model with 3rot sin i = 56 km/s. In green we plot a model with 3rot sin i = 116 km/s.
the star, given the moderate overabundance of He and N, together with the underabundance
of C and O (see Table 10.2).
We adopted the value of 3rot sin i = 56 km/s derived by Fraser et al. (2010). Previous estimations pointed to much higher values around 115 km/s (Zuiderwijk, 1995; Howarth et al.,
1997), but such a high rotational velocity is not compatible with some of the lines that
we see unblended in the optical observation (see Figure 10.17). The rotational velocity
NS
directly affects the estimation of the neutron star mass (MVela
) from radial velocity
X-1
curves, as shown by Koenigsberger et al. (2012). If 3rot sin i = 56 km/s, it is feasible that
NS
MVela
∼ 1.5 M⊙ , close to the canonical value (1.4 M⊙ ), instead of a high mass neuX-1
NS
tron star MVela
& 1.8 M⊙ , as suggested by other authors (e.g. Quaintrell et al. (2003);
X-1
Barziv et al. (2001)).
To summarise, our new analysis of Vela X-1 is in a broad agreement with previous studies of this system. We find a rather low stellar wind velocity, while Ṁ is typical for the stars
of its spectral type. Like other studies, we note spectral line variability in dependence with
orbital phase, and attribute it to the influence of the X-rays and the compact object on the
stellar wind.
The final physical parameters of the the two sources obtained in this work are presented
in Table 10.1.
Chapter 11
Discussion
11.1 Wind-fed accretion
In SFXTs and SGXBs, the X-ray emission is powered by the accretion of matter from
the donor’s wind onto the compact object. The efficiency of the conversion of the potential
energy into X-ray luminosity depends on many factors including the properties of the stellar
wind, the properties of the compact object and the orbital separation.
The most efficient way of producing X-rays is the so called Bondi-Hoyle-Littleton accretion or direct accretion (Hoyle and Lyttleton, 1939; Bondi and Hoyle, 1944): the stellar
wind that is gravitationally captured by the neutron star free-falls onto the compact object.
The expected luminosity is close to the accretion luminosity Lacc . The following equations
Parameters
J17544-2619
Vela X-1
Porb (d)
Pspin (s)
a sin i (lt-s)
hLX i (erg/s)
B (G)
4.9d
71.49e , 11.58 f
1.45g × 1012
8.964357a
283.532a
113.89a
4.5b × 1036
2.6c × 1012
Table 11.1: Parameters used in Chapter 11.
References: (a) Kreykenbohm et al. (2008) (b) Sako et al. (1999) (c) Kreykenbohm et al.
(2002) (d) Clark et al. (2009) (e) Drave et al. (2012) (f) Romano et al. (2015)
(g) Bhalerao et al. (2015)
93
94
Chapter 11. Discussion
contain the most relevant parameters in this regime:
Ra =
2GMNS
32rel
R2a
4a2
GMNS Ṁ⋆
= fa
RNS
fa =
Lacc
(11.1a)
(11.1b)
(11.1c)
where Ra is the accretion radius (also called Bondi radius), that is to say, the maximum
distance to the neutron star where the stellar wind is able to avoid falling onto the compact
object; G is the gravitational constant; MNS is the mass of the neutron star, which in this
work is hereafter assumed to be the canonical value 1.4 M⊙ ; 3rel is the velocity of the wind
relative to the neutron star; fa is the fraction of stellar wind that is gravitationally captured
by the neutron star; a is the orbital distance; Lacc is the accretion luminosity, namely, the
luminosity that would arise if the whole potential energy of the accreted matter is eventually transformed in X-ray luminosity (note that Eq. 11.1c is the already introduced Eq. 2.2,
adapted for the flux of mass gravitationally captured by the neutron star from the stellar
wind); and RNS is the radius of the neutron star, which in this work is henceforward assumed to be 15 km.
The direct accretion regime is only feasible when centrifugal and magnetic forces do not
hinder it. When these forces are strong enough, the accretion is partially inhibited and it
entails a decrease in the efficiency of the transformation from mechanical energy to X-rays
radiation. The higher the magnetic field and the shorter the Pspin , the more likely it is that
some kind of inhibition mechanism is at work.
The onset of the differently inhibited accretion regimes depends on the relative size of the
spheres defined by Ra , RM and Rco ; where Ra is the already defined accretion radius, RM is
the magnetospheric radius (location where the pressure exerted by the gas equals the local
magnetic pressure), and Rco is the co-rotation radius (location where the angular velocity
of the neutron star equals the Keplerian velocity). These radii, in turn, depend on: Ṁ, 3rel ,
magnetic moment of the neutron star (µ), orbital separation (a) and Pspin .
First, we calculate the Rco :
G M
1/3
p
2πRco
NS 2
P
= G MNS /Rco ⇒ Rco =
spin
Pspin
4 π2
(11.2)
Next, according to Bozzo et al. (2008), we can distinguish among the following different
possibilities of regimes where the accretion is highly inhibited:
11.1.1 Outside the magnetospheric radius: RM > Ra
When RM > Ra , the matter cannot reach the point where it is channeled towards the
neutron star because of the action of the magnetosphere. In this case, RM is defined by the
11.1. Wind-fed accretion
95
location r where the magnetospheric pressure (µ2 /(8πr6 )) equals the ram pressure of the
gas (ρ32rel ). Assuming that the density of the local medium is ρ ≃ Ṁ/(4 π a2 3rel ), we obtain
(Bozzo et al., 2008):
−1/6 −1/6 1/3
RM = 3.3 × 1010 a1/3
µ33 cm
10d Ṁ−6 38
(11.3)
Where a10d is a in units of 4.2 × 1012 cm, Ṁ−6 is Ṁ in units of 10−6 M⊙ /yr, 38 is 3rel in units
of 1000 km/s, and µ33 is µ in units of 1033 G cm3 .
The two following situations are possible:
• RM > Rco : the Super-Keplerian Magnetic Inhibition Regime. The local medium is
forced to co-rotate close to RM at super-Keplerian velocities. As a natural consequence, matter cannot be accreted.
• RM < Rco : the Sub-Keplerian Magnetic Inhibition Regime. In this regime matter
can penetrate the magnetosphere through Kelvin-Helmholtz instability at the magnetopause boundary (Harding and Leventhal, 1992).
11.1.2 Inside the magnetospheric radius: RM < Ra
In this case the accreted matter is channeled towards the neutron star before being halted
at the magnetospheric radius. Apart from direct accretion, we can consider the following
cases:
• RM > Rco : the Supersonic Propeller Regime. In this regime, studied in detail by
Davies et al. (1979) and Davies and Pringle (1981), the accreted matter forms an hydrostatic shell in the region between Ra and RM when radiative losses inside Ra are
negligible. The previous definition of RM (Eq. 11.3) is no longer valid, and in this
regime it can be approximated as (Bozzo et al., 2008):
−2/9 4/9 4/9
RM ≃ 2.3 × 1010 a4/9
10d Ṁ−6 38 µ33 cm
(11.4)
• RM < Rco , Ṁ < Ṁlim : the Subsonic Propeller Regime. When the rotation of the
magnetosphere is subsonic with respect to the surrounding material, the structure
of the hydrostatic shell above the magnetosphere changes and the location of the
magnetospheric radius can be approximated as (Bozzo et al., 2008):
−2/7 8/7 4/7
RM ≃ 2.0 × 1010 a4/7
Ṁ−6
38 µ33 cm
10d
(11.5)
Again, in this regime the accretion is driven by Kelvin-Helmholtz instabilities, and the
shell is consistent unless a critical mass-loss rate Ṁlim is surpassed and the radiative
losses are no longer negligible (Bozzo et al., 2008):
5/2
lim
2
−3/2
Ṁ−6
= 2.8 × 10−14 P−3
spin a10d 38 RM [1 + 16 Ra /(5 RM )]
(11.6)
−6 M /yr. If Ṁ > Ṁ
lim is Ṁ
Where Ṁ−6
lim in units of 10
⊙
lim the source enters the direct
accretion regime.
96
Chapter 11. Discussion
11.1.3 Transitions across regimes
First, we equal RM in Eq. 11.3 and Ra in Eq. 11.1a, and obtain:
1/11 −2/11 −2/11
38 = 1.1 Ṁ−6
µ33 a10d
(11.7)
Second, for RM > Ra , we equal RM in Eq. 11.3 and Rco in Eq. 11.2, obtaining:
−1 2 2
38 = 45.7 Ṁ−6
µ33 a10d P−4
sp3
(11.8)
Third, for RM < Ra , we equal RM in Eq. 11.4 and Rco in Eq. 11.2, obtaining:
1/2 −1 −1
38 = 0.5 Ṁ−6
µ33 a10d P3/2
sp3
(11.9)
Finally, from Eq. 11.6 and Eq. 11.5 we obtain the condition for the onset of the subsonic
propeller regime:
1/4 −4/15 −1/2 21/60
38 = 0.6 Ṁ−6
µ33 a10d Psp3
(11.10)
11.1.4 IGR J17544-2619
In IGR J17544-2619, using the results of our spectral fitting and the data shown in Table 11.1, we obtain from Eq. 11.1c: Lacc = 7.2· 1035 erg/s. However, the source spends most
of the time showing LX < 5 · 1034 erg/s (Bozzo et al., 2015). A high eccentricity making
fa much lower during most of the orbit is ruled out by the obtained radius of the star. As
mentioned above, this value implies e < 0.25 given that periodical Roche overflow is not
observed in the source. Most likely, some inhibition mechanism is acting in IGR J175442619 (Bozzo et al., 2008; Drave et al., 2014).
Bozzo et al. (2008) discussed the application of their model to the light curve of an
outburst observed by Chandra in this source. Nowadays, the tentative estimations of the
spin period (Pspin = 71.49s by Drave et al. (2012) and alternatively Pspin = 11.58s by
Romano et al. 2015), along with the stellar wind parameters derived in this work, permit to
discuss the accretion regime transitions from a new perspective.
Using the values showed in Table 11.1 and the Eq. 11.7-11.10, we can elaborate diagrams
3rel - Ṁ and 3rel -Pspin where the different regions are occupied by the variety of possible accretion regimes.
In Fig. 11.1 we show the position of IGR J17544-2619 in the diagram 3rel - Ṁ for the two
currently available tentative estimations of the Pspin . The source lie in the direct accretion
regime for Pspin = 71.49s, and in the supersonic propeller regime for Pspin = 11.58s. We
can see in Fig. 11.2 the diagram 3rel -Pspin where we show that the inhibition of accretion
would occur at spin periods shorter than ∼ 30s. Hence, the shortest Pspin = 11.58s matches
better with the X-ray behavior of the source and its likelihood of staying in an inhibited
accretion regime.
11.1. Wind-fed accretion
97
As we can see in Fig. 11.1, modest variability in the stellar wind of IGR J17544-2619
may trigger the shift between different accretion regimes. In the case of Pspin = 11.58s, an
increase of density by one order of magnitude in the medium transited by the neutron star
would be required. Such a density jump is fully plausible, as demonstrated by hydrodynamical simulations of radiatively driven stellar winds (e.g. Feldmeier et al., 1997). These
clumps of higher density, intrinsic to stellar winds of hot stars, are sometimes invoked to
explain the X-ray variability of HMXBs (Oskinova et al., 2012). It seems that in objects
like IGR J17544-2619, the abrupt changes in the wind density may lead to the switching
from one accretion regime to the other.
Considering an alternative explanation for the X-ray variability of IGR J17544-2619,
Drave et al. (2014) invoked the quasi-spherical accretion model by Shakura et al. (2012b).
However, if the spin period is actually as short as 71.49s or 11.58s, the condition of a slowly
rotating pulsar, i.e. RM ≪ Rco , assumed by this approach, would be debatable. Even though
it raises doubts about the feasibility of applying this model, it cannot be ruled out until the
spin period and the magnetic field of the neutron star are firmly constrained.
In the case of Vela X-1, we can see in Fig. 11.3 that the source is well in the middle of
the zone where direct accretion is expected. Hence, more extreme density or velocity jumps
would be required to trigger any change of accretion regime. These extreme jumps are also
plausible, but much more unlikely. However, they might sporadically occur and lead to a
sudden decrease of the luminosity in Vela X-1.
Using the parameters shown in Table 11.1 and Eq. 11.1c, we obtain Lacc = 6.5×1036 erg/s
for Vela X-1. This luminosity is very close to the average X-rays luminosity (hLX i ≃ 0.7 ×
Lacc ). The very good agreement with Lacc implies that the direct accretion scenario can
describe the way that matter is accreted in Vela X-1.
The framework of different accretion regimes described by Bozzo et al. (2008) is able to
explain why IGR J17544-2619 is prone to show a high X-ray variability and inhibited accretion (assuming the shortest Pspin = 11.58s), and Vela X-1 is persistently very luminous in
the X-rays. As exposed in Fig. 11.1 and 11.3, the required variability in the stellar wind for
a transition in the accretion regime is far lower in IGR J17544-2619 than in Vela X-1. The
main ingredients that make the sources so different are the Pspin (shorter in IGR J175442619), and the 3rel (larger in IGR J17544-2619). Therefore, we may conjecture that the
radically different X-ray behavior of SFXTs and SGXBs is due to the combination of Pspin
and 3rel , that might tend to be shorter and higher, respectively, in SFXTs.
Pieces of evidence exist in this direction. In the following we expose our current knowledge about Pspin and 3∞ in SGXBs and SFXTs:
• In SGXBs, Pspin > 100s. The only known exception is OAO 1657-415 (Pspin =
38.2s). Even though such a short Pspin hinders the direct accretion, this difficulties
might be overcome in this system because of the extraordinarily low 3∞ ∼ 250 km/s
of the Ofpe/WN9 donor (Mason et al., 2012). As we can see in Fig. 11.4, such a low
98
Chapter 11. Discussion
SuperKeplerian magnetic inhibition
IGR J17544-2619
Supersonic propeller
Direct accretion
SubKeplerian magnetic inhibition
SuperKeplerian magnetic inhibition
nic
rso r
e
p
Su pelle
pro
eller
prop
c
i
n
o
ubs
S
IGR J17544-2619
Direct accretion
Figure 11.1: Position of IGR J17544-2619 in the 3rel - Ṁ diagram. We plot the following lines
defining the transition accross different regimes: Eq. 11.7 (solid line), Eq. 11.8 (dotted line),
Eq. 11.9 (triple-dot-dashed line) and Eq. 11.10 (dot-dashed line). Upper panel: diagram
calculated using Pspin = 11.58s. Lower panel: diagram calculated using Pspin = 71.49s.
11.1. Wind-fed accretion
99
SubKeplerian
magnetic
inhibition
SuperKeplerian magnetic inhibition
nic
Subso
IGR J17544-2619
ller
prope
(Pspin = 11.58s)
IGR J17544-2619
(Pspin = 71.49s)
er
l
c
pe
Su
i
on
o
pr
l
pe
rs
Direct accretion
Figure 11.2: Position of IGR J17544-2619 in the 3rel -Pspin diagram for the two tentative
estimations of Pspin .
Sub K
inhibition
Super K
Supersonic
pr
hibition
Subsonic
pr
Vela X-1
Dir
cretion
Figure 11.3: Position of Vela X-1 in the 3rel - Ṁ diagram.
100
Chapter 11. Discussion
SuperK
SubKe
e
me
ibition
inhibition
ller
ic
son
Sub
pe
pro
Supersonic propeller
Vela X-1
Dir
cretion
Figure 11.4: Position of Vela X-1 in the 3rel -Pspin diagram.
stellar wind velocity almost guarantees the permanence in the direct accretion regime
even for very short spin periods.
• Among SFXTs, there are two systems with Pspin < 100s: IGR J17544-2619 (tentative estimations, see Table 11.1) and IGR J18483-0311 (Pspin = 21.05s, Sguera et al.
2007b). There are SFXTs showing Pspin > 100s too. However, according to the recent review by Walter et al. (2015), their definition as SFXTs is debatable because
all of them are actually similar to other classic systems, or either they lack of good
observations. The Pspin of the rest of SFXTs remains unknown.
• The scarce number of available estimations of 3∞ in SGXBs points to low wind velocities, similar to the estimation for Vela X-1 in this work: 305 km/s in GX 301-2
(Kaper et al., 2006), 500 − 600 km/s in IGR J17252-3616 (Manousakis et al., 2012),
500 km/s in X1908+075 (Martínez-Núñez et al., 2015), and the already mentioned
3∞ ∼ 250 km/s in OAO 1657-415. Another SGXB, 4U1700-37, exhibits a larger
3∞ = 1700 km/s (van Loon et al., 2001), but interestingly, this source is candidate for
hosting a black hole instead of a neutron star, due to the lack of detected pulsations.
• Giménez-García et al. (2015) analysed a large sample of X-ray observations of HMXBs
and found a significant difference between SGXBs and SFXTs in the average column density (density from the X-ray source to the observer integrated over the line
of sight). The column density is fundamentally local in these systems. The group of
SGXBs showed systematically higher column densities than the SFXTs, in agreement
with our hypothesis of slower winds in the SGXBs’ donors.
• In SFXTs, Lorenzo et al. (2014) assumed 3∞ = 1230 km/s in their analysis of IGR J112155952. There are no further 3∞ estimations in SFXTs apart from this work.
11.2. Evolutionary tracks
101
In summary, we are far from a complete knowledge of Pspin and 3∞ in SGXBs and
SFXTs. However, in the light of this work and the available information from previous
studies, there are indeed grounds to think that SFXTs tend to show higher 3rel and shorter
Pspin than SGXBs. We note that a high 3rel or short Pspin are not neccessarily defining features of SFXTs. Rather, the specific combination of these parameters is what might make
them prone to shift between different accretion regimes. Moreover, the relative importance
of Pspin and 3rel for the onset of the different inhibited accretion regimes strongly depends
on the specific accretion regime involved. For instance, the onset of the subsonic propeller
regime is much more sensitive to 3rel than the onset of the supersonic propeller regime,
as we can see from the different slopes of their characteristic lines in Fig. 11.1-11.4. Provided that the intersection of the characteristic lines of the supersonic and subsonic propeller
regimes takes place at certain 3K , the stellar wind velocity gains special relevance whenever
3rel > 3K . In case 3rel < 3K , the transition between direct accretion and subsonic propeller
regimes is not longer possible, and the inhibition of accretion is produced only in case Pspin
is short enough.
Finally, we can compare the 3∞ and the 3esc that we obtain from the fits. In this regard,
Lamers et al. (1995) collected a large dataset from hot stars with radiatively driven winds,
and concluded that the ratio 3∞ /3esc steeply decreases from ∼ 2.6 to ∼ 1.3 when going from
high to low T eff at a point near T eff ≃ 21000 K, corresponding to spectral type around B1.
According to Vink et al. (1999), this drop is caused by a decrease in the line acceleration of
Fe iii in the subsonic part of the wind. In our case we have (see Table 10.1):
• IGR J17544-2619 (O9.5I): 3∞ /3esc = 2.4+0.7
−0.5
• Vela X-1 (B0.5I): 3∞ /3esc = 1.6+0.8
−0.4
These values follow the trend observed and described by Lamers et al. (1995). We suggest
that it might be the reason why IGR J17544-2619 shows higher 3∞ than Vela X-1. The action
of the X-rays can also make an important impact in the velocity of the stellar wind, as shown
by Karino (2014). However, this effect is probably local, since we do not observe important
differences in the terminal velocity between eclipsing and non-eclipsing orbital phases in
Vela X-1. Secondary features like asymmetries or additional absorption components in the
spectral lines, which might be related to the effect of the X-rays in the stellar wind, are
described and discussed below in Sect. 11.3.
11.2 Evolutionary tracks
In Fig. 11.5 we show the position of Vela X-1 and IGR J17544-2619 in the HertzprungRussell Diagram (HRD), and the evolutionary tracks from the Geneva Stellar Models (Ekström et al.,
2012). The two stars lie on the theoretical track of a star with initial mass ∼ 25 − 30 M⊙ .
In IGR J17544-2619 the spectroscopic mass obtained from the fits is compatible with the
evolutionary mass. Vela X-1 shows certain overluminosity, since its spectroscopic mass is
lower than the evolutionary mass. Nevertheless, the mass of the star obviously decreases
along its lifetime due to the stellar wind and possible mass transfer episodes. These phenomena might have been stronger or longer in Vela X-1 compared to IGR J17544-2619.
102
Chapter 11. Discussion
Figure 11.5: Evolutionary tracks from the Geneva Stellar Models with solar abundances
and rotation. The positions of IGR J17544-2619 (square) and Vela X-1 (diamond) are overplotted.
The overabundance of helium and nitrogen arising from the fits in the two stars might
trigger an increase in luminosity following the scaling relation L ∝ µα , where µ is the
average mean molecular weight and α > 1 (Langer, 1992). Then, we expect certain overluminosity in both sources. However, as already mentioned, the overluminosity is more
noticeable in Vela X-1. In all, the sources seem to be in a different evolutionary stage or to
have experienced a different evolutionary history.
The chemical evolution of the donors might have been driven by episodes of important
mass transfer in the past, given the close orbits of the systems, enhancing the helium and
nitrogen abundances due to the accretion of chemically enriched material (Langer, 2012).
Moreover, Roche-lobe overflow stages induce important spin-up in the mass gainer (Packet,
1981), inducing further chemical enrichment because of rotational mixing. This scenario is
supported by the observation of other HMXBs where indications of nitrogen enhancement
are also observed (González-Galán et al., 2014a).
11.3 Asymmetries in spectral lines of Vela X-1
Some of the lines in the spectrum of Vela X-1 show clear asymmetries that are not possible to reproduce with spherically symmetric models like PoWR (see Fig. 11.6). This
striking feature is specially noticeable for He i lines, but it is also observed in C, N, O or Si,
whenever the lines are strong enough.
11.3. Asymmetries in spectral lines of Vela X-1
Normalized flux
Normalized flux
HeI
1.0
HeI
1.0
103
0.8
3815
3820
o
λ/A
3825
0.8
4020
1.0
4025
o
λ/A
1.0
Normalized flux
SiIII
HeI
Normalized flux
4030
0.8
0.8
0.6
4465
4470
o
λ/A
4475
4550
4555
o
λ/A
Figure 11.6:
Example of four spectral lines showing notable asymmetries:
He i λ3820, 4026, 4471 Å, and Si iii λ4553 Å.
104
Chapter 11. Discussion
Asymmetries in spectral lines were also reported by Martínez-Núñez et al. (2015) in hydrogen lines of the infrared spectrum of X1908+75, a SGXB. A natural explanation for the
discrepancy between models and observations is the departure of the donor and/or the surrounding medium from the spherical symmetry. This departure may be triggered by tidally
induced effects and the persistent X-ray irradiation of the stellar wind and the stellar surface. In this regard, Koenigsberger et al. (2012) showed that tidal effects would produce
asymmetries in the line profiles.
The observed asymmetries might be related to the additional absorption that we observe in the blue part of other important lines, with special attention to Hα, Hβ, Hγ and
Si iv λ 1394, 1403 Å (see Fig. 11.7). Assuming that the absorption is produced by an independent component of matter moving at certain velocity, it is striking that the involved
velocities required for explaining such a blueshif are different depending on the lines:
∼ 200 − 300 km/s in Hα, Hβ and Hγ, ∼ 1000 km/s in the Si iv resonance lines.
In any case, we note that these asymmetries and additional absorption features have not
been observed in IGR J17544-2619. Hence, the physical cause at work is playing a significantly more important role in Vela X-1 than in IGR J17544-2619. This fact suggests that the
interaction of the X-ray source with the stellar wind might be fundamental for understanding
these asymmetries, given that the X-rays are on average more intense in Vela X-1. Indeed,
if we compare the wind mechanic luminosity Lmech = Ṁ 32∞ /2 to the X-ray luminosity LX
we obtain:
• IGR J17544-2619: Lmech ≃ 1036 erg/s. That is to say, at least two orders of magnitude
higher than the usual X-ray luminosity of the source.
• Vela X-1: Lmech ≃ 1035 erg/s. Namely, about one order of magnitude lower than the
X-ray luminosity of the source in quiescence.
Hence, there is a fundamental difference in the ratio Lmech /LX . The X-rays are much
more powerful with respect to the stellar wind in Vela X-1 rather than in IGR J17544-2619.
We suggest that this fact might be related to the asymmetries that we observe in the spectral
lines of Vela X-1, but not in IGR J17544-2619.
Hα
11.3. Asymmetries in spectral lines of Vela X-1
105
Hβ
1.0
1.0
6550
6560
6570
o
λ/A
6580
4850
4855
0.6
4335
4340
o
λ/A
4865 4870
1.0
Normalized flux
Normalized flux
Hγ
1.0
0.8
4860
o
λ/A
SiIV
0.8
0.8
SiIV
Normalized flux
Normalized flux
1.2
0.5
0.0
1390
1400
o
λ/A
1410
Figure 11.7: Hα, Hβ, Hγ and Si iv λ1394, 1403 Å. The observations (blue solid line) shows
an additional blueshifted component that we are not able to reproduce with the models (red
dashed line).
Chapter 12
Summary
We have performed a detailed analysis of the donors of the HMXBs IGR J17544-2619
and Vela X-1, using the code PoWR that computes models of hot stellar atmospheres. We
found the luminosity, extinction, stellar mass, stellar radius, effective temperature, effective
surface gravity, terminal velocity of the stellar wind, mass-loss rate, clumping factor, micro
and macro-turbulent velocity, rotational velocity, chemical abundances and extinction.
The estimation of the above mentioned parameters has implications on other physical
parameters of the system: the derived stellar radius of IGR J17544-2619 implies an upper
limit in the eccentricity of the source: e < 0.25. The rotational velocity derived for Vela X-1
NS
implies that the mass of the neutron star might be MVela
∼ 1.5 M⊙ , close to the canonical
X-1
value (1.4 M⊙ ).
The donors of IGR J17544-2619 and Vela X-1 are similar in many of the parameters that
physically characterise them and their spectrum. Moreover, they are also comparable in the
eccentricity and orbital separation. However, in the context of accretion regimes described
by Bozzo et al. (2008), their moderate differences in the stellar wind velocity and the Pspin
of the neutron star lead to a very different accretion regimes of the sources, which qualitatively explain their completely different X-ray behavior. Further explorations addressing the
estimation of the stellar wind properties of the donors in SGXBs and SFXTs, complemented
with Pspin measurements in SFXTs, will be necessary to confirm whether the conclusions
exposed here can be extrapolated to additional members of these groups of HMXBs.
In summary, this study shows that the wind terminal velocity and the neutron star pulse
period play a decisive role in determining the class of HMXB hosting a supergiant donor.
While the combination of these parameters allows direct steady accretion in SGXBs, the
high wind velocity and velocity jumps can easily shift the accretion mechanism from direct
accretion to propeller regimes in SFXTs. We conclude that this may be the mechanism responsible for these two major sub-classes of HMXBs with supergiant donors.
107
Part IV
Conclusions
Massive stars are a key element in the evolution of galaxies because of its role in the feedback and enrichment of material. Provided that the circumstellar medium of these stars
gives us insight about many of their properties, its study is highly relevant scientifically. In
the case of HMXBs, we have the additional possibility of using the X-ray source as an actual scanner of the region surrounding the massive star, which is a wide opened window for
further research. In this work we have taken profit of the range of possibilities offered by the
X-rays, infrared, optical and ultraviolet observations and we have used these possibilities to
describe and interpret the properties of the circumstellar medium in massive stars.
First, we have performed a comprehensive spectral analysis of the available set of observations with XMM-Newton until August 2013. The number of sources and FeKα detections
makes this study the most complete survey of FeKα in HMXBs to date. We have found that
the different groups of HMXBs show very different patterns in the Fe complex as a natural
consequence of their very distinct properties of their circumstellar medium. The Fe complex patterns of SGXBs, showing ubiquitous fluorescence, and γ Cass-analogs, showing
recombination and fluorescence, are specially clear.
We have observed correlation between the EW of FeKα and NH (the curve of growth),
indicating that there is a strong link between the reprocessing and the absorbing region.
This connection permits to derive physical properties of the absorber from the properties of
FeKα. Moreover, the observed correlations also provide information about the geometry of
the absorber, showing that its distribution is roughly isotropic and covers a large volume,
from the immediate surroundings of the compact object to a distance of the order of the
stellar radius.
The systems with a supergiant donor, SGXBs and SFXTs, tend to show more prominent FeKα due to the larger NH that they exhibit. This higher absorption is probably due
to the strong stellar wind of the donors, as shown from the orbital modulation of NH in
IGR J16320-4751 and 4U 1700-37. However, we have observed a significant difference in
the NH between SGXBs and SFXTs, being systematically higher in SGXBs. It suggests the
110
possibility that the stellar winds in SGXBs are in general more dense than in SFXTs, or the
interaction stellar wind - compact object is different.
In order to gain insight in the previous conjecture, we have performed a detailed analysis
of the donor in two representative sources of the SGXBs and SFXTs, namely, Vela X-1
and IGR J17544-2619, respectively. We have used the PoWR code in order to model the
atmosphere and stellar wind of these stars and fit their spectra in the infrared, optical and
ultraviolet. This analysis has permitted us to obtain an estimation of the following parameters of the donors: luminosity, extinction, stellar mass, stellar radius, effective temperature,
effective surface gravity, terminal velocity of the stellar wind, mass-loss rate, clumping factor, micro and macro-turbulent velocity, rotational velocity and chemical abundances. From
these parameters, we have been able to derive other ones from previous works of other authors: the stellar radius of IGR J17544-2619 constrains the eccentricity of the system to
e < 0.25. The rotational velocity of Vela X-1 implies that the mass of the neutron star might
be close to the canonical value.
Given the parameters that we have found in the analysis of both stars, we have seen that
the donors in Vela X-1 and IGR J17544-2619 are not particularly peculiar, but they agree
well with their respective spectral types, which, in turn, are not very different one to the
other. Moreover, the orbital parameters of the system are comparable too. However, we
have seen that in the framework described by Bozzo et al. (2008) their moderate differences
in the stellar wind combined with a different spin period of the neutron star may lead to
different accretion regimes and different likelihood of transition across regimes. This fact
would qualitatively explain their completely different behavior in the X-rays.
In all, we have seen along this work how the study of the X-rays, and FeKα in particular,
provide significant information about the surroundings of the compact object. This information can be used to understand better the circumstellar medium of massive stars, which in
turn has important consequences in many astrophysical fields. At the same time, a detailed
analysis of the donors’ spectra, using sophisticated models, gives a deep characterization
of the physical properties of those stars. We have shown why these detailed analyses are
crucial in order to interpret the current picture of the HMXBs zoo.
Appendix A
The fits of the FeKα survey
In this Appendix we show the results arising from the fits described in Part II. First, we
present the complete set of observations and fits. Second, we show the parameters obtained
from the fits, consisting of two tables: the first one shows the parameters of the continuum
and the second one the parameters of FeKα.
113
114
Appendix A. The fits of the FeKα survey
BeXB
2 -2
1A 0535+26 Obs 0674180101
Obs 0302970801
X Persei Obs 0151380101
0.1
0.01
5×10−3
2×10−3
1.2
1
0.1
1.2
1.2
ratio
1
ratio
ratio
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
1
1
0.8
0.8
0.8
2
5
2
5
5
Energy (keV)
Energy (keV)
Energy (keV)
X Per Obs 0600980101
RX J01469+6121 Obs 0201160101
AX J1820.5−1434 Obs 0511010101
1
0.1
0.01
5×10−3
0.1
0.01
2×10−3
10−3
1.2
1
1.2
ratio
ratio
1.2
ratio
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.02
1
1
0.8
0.8
0.8
2
2
5
2
Energy (keV)
5
Energy (keV)
5
Energy (keV)
RX J0440.9+4431 Obs 0653660101
RX J1037.5−5647 Obs 0550560101
SAX J2103.5+4545 Obs 0149550401
0.1
0.5
0.2
0.1
0.05
0.02
0.01
5×10−3
0.05
1.2
1
0.8
2
1
0.5
ratio
ratio
1
0.8
2
0.2
1.2
1.2
ratio
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
1
0.8
5
2
Energy (keV)
5
2
Energy (keV)
5
Energy (keV)
Swift J045106.8−69 Obs 0679381401
V0332+53 Obs 0506190101
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.1
0.05
0.02
0.01
10−3
10−4
5×10−3
1.2
ratio
ratio
1.2
1
0.8
1
0.8
2
5
Energy (keV)
2
5
Energy (keV)
Figure A.1: BeXBs data, model, model components and ratio data/model. The spectra are
typically soft, with no Fe emission lines.
115
SGXB
4U 1538−522 Obs 0152780201
4U 1538−522 Obs 0152780201
1700−37 Obs 0083280101
2
50
0.1
0.02
0.01
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
0.05
0.5
0.2
0.1
20
10
5
2
5×10−3
0.05
1.2
1.2
1.2
1
ratio
ratio
ratio
1.1
1
1
0.9
0.8
0.8
2
0.8
5
2
5
2
5
Energy (keV)
Energy (keV)
Energy (keV)
1700−37 Obs 0083280101
1700−37 Obs 0083280101
1700−37 Obs 0083280201
20
5
50
2
1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
10
5
2
1
0.5
20
10
5
0.5
1
1.2
ratio
1.2
ratio
ratio
1.2
1
0.8
0.8
2
5
2
5
2
Energy (keV)
Energy (keV)
1700−37 Obs 0083280301
1700−37 Obs 0083280401
2
20
10
5
2
normalized counts s−1 keV−1
normalized counts s−1 keV−1
10
normalized counts s−1 keV−1
5
Energy (keV)
1700−37 Obs 0083280201
20
5
2
1
0.1
0.05
1.2
1
0.8
2
0.2
1.2
ratio
ratio
1
0.8
1
0.5
0.5
1
1.2
ratio
1
0.8
1
0.8
5
2
5
Energy (keV)
2
5
Energy (keV)
1
1700−37 Obs 0083280401
7
Energy (keV)
Obs 0600950101
4U 1907+09 Obs 0555410101
10
10
2
1
0.5
0.5
0.2
0.1
0.05
0.2
1.2
5
2
1
0.5
1.2
ratio
1.2
1
ratio
ratio
normalized counts s−1 keV−1
5
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
1
0.8
1
0.8
2
0.8
5
2
Energy (keV)
2
5
Energy (keV)
5
Energy (keV)
4U 1907+09 Obs 0555410101
IGR J17255−3617 Obs 0206380401
IGR J17252−3616 Obs 0405640201
1
1
0.5
normalized counts s−1 keV−1
2
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.01
0.1
0.2
0.01
1.2
1.2
10−3
1
0.8
1
0.8
2
5
1
0.8
2
5
2
5
Energy (keV)
Energy (keV)
IGR J17252−3616 Obs 0405640401
IGR J17252−3616 Obs 0405640901
0.1
normalized counts s−1 keV−1
Energy (keV)
IGR J17252−3616 Obs 0405640301
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1.2
ratio
ratio
ratio
10−4
0.1
0.01
0.1
0.01
0.01
10−3
1
0.8
1.2
ratio
1.2
ratio
ratio
1.2
1
0.8
2
5
Energy (keV)
1
0.8
2
5
Energy (keV)
2
5
Energy (keV)
Figure A.2: SGXBs data, model, model components and ratio data/model. The spectra are
characteristically affected by high absorption, with Fe fluorescent lines.
116
Appendix A. The fits of the FeKα survey
IGR J17252−3616 Obs 0405640801
GX 301−2 Obs 0555200401
GX 301−2 Obs 0555200401
0.1
0.01
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
10
10
1
0.1
1.2
1.2
1
0.1
0.01
1.2
1
ratio
ratio
ratio
1.1
1
1
0.9
0.8
0.8
2
5
2
5
2
Energy (keV)
Energy (keV)
IGR J16318−4848 Obs 0154750401
IGRJ16138−4848 Obs 0201000201
0.05
0.02
0.01
0.1
0.1
0.01
0.01
10−3
10−3
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.1
10−4
10−5
10−6
10−7
10−8
10−9
10−10
10−11
10−12
0.8
1
0.8
2
10−6
10−7
10−8
10−9
10−10
10−11
1.2
ratio
ratio
ratio
1.2
1
10−5
10−13
10−14
1.2
10−4
10−12
10−13
5×10−3
1
0.8
5
2
Energy (keV)
5
2
5
Energy (keV)
IGRJ16138−4848 Obs 0201000301
Energy (keV)
IGRJ16138−4848 Obs 0201000401
1.2
IGR J16320−4751 Obs 0128531101
0.1
0.01
10−3
10−4
10−5
10−6
10−7
10−8
10−9
10−10
10−11
10−12
10−13
10−14
0.1
normalized counts s−1 keV−1
0.1
0.01
10−3
10−4
10−5
10−6
10−7
10−8
10−9
10−10
10−11
10−12
10−13
10−14
10−15
normalized counts s−1 keV−1
normalized counts s−1 keV−1
5
Energy (keV)
IGR J16207−5129 Obs 0402920201
0.2
normalized counts s−1 keV−1
0.8
0.05
0.02
0.01
1.2
1.2
ratio
ratio
ratio
1.1
1
1
1
0.9
0.8
0.8
2
5
0.8
Energy (keV)
Energy (keV)
5
Energy (keV)
IGR J16320−4751 Obs 0201700301
2
IGR J16320−4751 Obs 0201700301
5
IGR J16320−4751 Obs 0556140101
1
0.1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
0.1
0.01
0.1
0.01
1.2
1.2
1.1
1.1
1.1
1
ratio
1.2
ratio
ratio
0.01
1
1
0.9
0.9
0.9
0.8
0.8
0.8
2
5
2
5
2
5
Energy (keV)
Energy (keV)
Energy (keV)
IGR J16320−4751 Obs 0556140201
IGR J16320−4751 Obs 0556140301
IGR J16320−4751 Obs 0556140401
1
normalized counts s−1 keV−1
0.1
1.2
1.2
1.1
1.1
1.1
1
ratio
1.2
1
1
0.9
0.9
0.9
0.8
0.8
0.8
5
2
5
Energy (keV)
Energy (keV)
IGR J16320−4751 Obs 0556140501
IGR J16320−4751 Obs 0556140601
2
5
Energy (keV)
IGR J16320−4751 Obs 0556140701
1
0.1
1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
0.1
0.1
0.01
0.01
1.2
1.2
1.2
1.1
1.1
1.1
1
ratio
0.01
ratio
normalized counts s−1 keV−1
0.01
0.01
2
ratio
0.1
0.01
ratio
normalized counts s−1 keV−1
ratio
0.1
normalized counts s−1 keV−1
1
1
1
1
0.9
0.9
0.9
0.8
0.8
0.8
2
5
Energy (keV)
2
5
Energy (keV)
2
5
Energy (keV)
Figure A.2: SGXBs data, model, model components and ratio data/model. The spectra are
characteristically affected by high absorption, with Fe fluorescent lines (continued).
117
IGR J16320−4751 Obs 0556140801
IGR J16320−4751 Obs 0556141001
IGR J16465−4507 Obs 0164561001
1
0.05
0.1
0.01
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
0.1
0.02
0.01
5×10−3
0.01
1.1
1
1
0.9
0.9
0.8
0.8
2
1.2
ratio
1.2
1.1
ratio
ratio
2×10−3
1.2
1
0.8
5
2
Energy (keV)
SAX J1802.7-2017 Obs 0206380601
Energy (keV)
5
Energy (keV)
Vela X−1 Obs 0111030101
5
Vela X−1 Obs 0111030101
2
1
0.1
0.05
0.02
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.2
1
0.5
0.2
0.1
0.5
0.2
0.1
0.05
0.01
0.05
1.2
1.2
1.2
ratio
ratio
ratio
1.1
1
1
1
0.9
0.8
0.8
0.8
2
2
5
5
2
5
Energy (keV)
Energy (keV)
Energy (keV)
normalized counts s−1 keV−1
XTE J0421+560 Obs 0139760101
0.01
10−3
1.2
ratio
1.1
1
0.9
0.8
2
5
Energy (keV)
Figure A.2: SGXBs data, model, model components and ratio data/model. The spectra are
characteristically affected by high absorption, with Fe fluorescent lines (continued).
118
Appendix A. The fits of the FeKα survey
SFXT
AXJ1841.0−0536 Obs 0604820301
IGR J00370+6122 Obs 0501450101
IGR J11215−5952 Obs 0405181901
1
0.1
2
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
0.1
1
0.5
0.2
0.01
0.1
1.2
1.2
1.2
ratio
ratio
ratio
1.1
1
1
1
0.9
0.8
0.8
2
5
2
5
Energy (keV)
Energy (keV)
IGR J11215−5952 Obs 0405181901
IGR J16328−4726 Obs 0654190201
2
5
Energy (keV)
IGR J16418-4532 Obs 0206380301
0.1
0.5
0.1
0.2
0.1
0.05
0.05
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.8
0.02
0.01
5×10−3
0.01
0.02
2×10−3
1.2
1.2
1.2
ratio
1
1
ratio
ratio
1.1
1
0.9
0.8
0.8
2
5
2
0.8
5
Energy (keV)
Energy (keV)
2
5
Energy (keV)
IGR J16418−4532 Obs 0405180501
IGRJ16479−4514 Obs 0512180101
IGRJ16479−4514 Obs 0512180101
0.1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.05
0.02
0.01
normalized counts s−1 keV−1
0.02
0.01
5×10−3
2×10−3
0.1
0.01
5×10−3
10−3
1
1.2
1.1
1.1
1
0.8
2
ratio
1.2
ratio
ratio
1.2
0.9
0.8
0.8
5
2
5
Energy (keV)
T X
1
0.9
2
5
Energy (keV)
39-302 Obs 0554720101
X
Energy (keV)
39-302 Obs 0554720101
XTE1739−302 Obs 0561580101
0.2
0.02
0.01
5×10−3
2×10−3
1.2
0.01
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.05
5×10−3
2×10−3
0.1
0.05
0.02
0.01
10−3
1.2
ratio
ratio
1.2
ratio
1
1
1
0.8
0.8
2
2
5
5
0.8
Energy (keV)
Energy (keV)
2
5
Energy (keV)
IGR J17544-2619 Obs 0148090501
XTE1739−302 Obs 0561580101
IGR J18450−0435 Obs 0306170401
2
0.02
0.01
5×10−3
2×10−3
0.1
0.05
0.02
ratio
ratio
0.2
1.2
1.2
1
1
0.5
0.1
0.01
1.2
1
0.8
1
0.8
2
0.8
5
2
Energy (keV)
2
5
Energy (keV)
5
Energy (keV)
IGR J18450−0435 Obs 0306170401
IGR J18483−0311 Obs 0406140201
0.02
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.2
0.1
0.05
0.02
0.01
5×10−3
2×10−3
10−3
1.2
ratio
1.2
ratio
ratio
0.2
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.05
1
0.8
1
0.8
2
5
Energy (keV)
2
5
Energy (keV)
Figure A.3: SFXTs data, model, model components and ratio data/model.
119
γ Cassiopeae-like
GAMMA CAS Obs 0201220101
gamma Cassiopeiae Obs 0651670201
1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
gamma Cassiopeiae Obs 0651670301
10
10
10
1
1
0.1
0.1
0.1
1.2
1.2
1.2
1
ratio
ratio
ratio
1.1
1
1
0.9
0.8
0.8
2
0.8
5
2
5
Energy (keV)
Energy (keV)
gamma Cassiopeiae Obs 0651670401
gamma Cassiopeiae Obs 0651670501
2
5
Energy (keV)
HD 110432 Obs 0504730101
1
10
1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
10
1
0.1
0.01
0.1
0.1
1
1.2
ratio
1.2
ratio
ratio
1.2
1
0.8
1
0.8
2
0.8
5
2
5
Energy (keV)
2
5
Energy (keV)
H !
HD 119682 Obs 0551000201
Energy (keV)
Obs 0551020101
HD 161103 Obs 0201200101
0.1
0.1
0.01
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.1
0.01
0.01
−3
10
10−3
10−3
1.2
1.2
1.2
ratio
ratio
ratio
1.1
1
1
1
0.9
0.8
0.8
2
0.8
5
2
Energy (keV)
2
5
Energy (keV)
5
Energy (keV)
HD 45314 Obs 0670080301
SS "#$
SAO 49725 Obs 0201200201
Obs 0122700101
0.1
0.02
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.05
0.01
0.02
0.01
5×10−3
2×10−3
5×10−3
2×10−3
10−3
10−3
1.2
0.01
1.2
ratio
1
0.8
ratio
ratio
1.2
1
1
0.8
2
5
2
Energy (keV)
0.8
5
Energy (keV)
2
5
Energy (keV)
%% &'(
)) *+,
Obs 0122700201
.. /03
Obs 0122700301
Obs 0122700501
0.02
0.02
0.01
5×10
−3
2×10−3
10
−3
0.01
0.01
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
0.02
5×10−3
2×10−3
10
−3
5×10−3
2×10−3
10−3
5×10−4
5×10−4
1.2
1.2
ratio
1
ratio
ratio
1.2
1
1
0.8
0.8
2
5
Energy (keV)
0.8
2
2
5
5
Energy (keV)
Energy (keV)
Figure A.4: γ Cass-like systems data, model, model components and ratio data/model. The data
is usually fitted using thermal models, including Fe recombination lines. FeKα is also usually
visible.
120
Appendix A. The fits of the FeKα survey
HMGB
0.1
0.01
0.01
1.2
1.2
1
0.8
normalized counts s−1 keV−1
0.1
LS I +61 303 Obs 0505980901
1.2
1
5
1
0.8
0.8
2
0.1
0.01
ratio
normalized counts s−1 keV−1
LS I +61 303 Obs 0505980801
ratio
ratio
normalized counts s−1 keV−1
LS I +61 303 Obs 0207260101
2
5
2
5
Energy (keV)
Energy (keV)
Energy (keV)
LS I +61 303 Obs 0505981001
LS I +61 303 Obs 0505981101
LS I +61 303 Obs 0505981201
0.1
0.1
0.1
0.01
0.01
1
0.8
1.2
ratio
1.2
ratio
1.2
ratio
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
1
0.8
2
5
1
0.8
2
5
2
5
Energy (keV)
Energy (keV)
Energy (keV)
LS I +61 303 Obs 0505981301
LS I +61 303 Obs 0505981401
LS 5039 Obs 0151160201
0.1
0.1
0.1
0.01
0.01
0.01
1
0.8
1.2
ratio
1.2
ratio
1.2
ratio
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
1
1
0.8
2
5
1
0.8
2
Energy (keV)
5
Energy (keV)
2
5
Energy (keV)
normalized counts s−1 keV−1
LS 5039 Obs 0151160301
0.1
0.01
ratio
1.2
1
0.8
2
5
Energy (keV)
Figure A.5: HMGBs data, model, model components and ratio data/model. Soft spectra with no
sign of Fe lines.
121
Peculiars
4U 2206+54 Obs 0650640101
Cen X−3 Obs 0111010101
CYGNUS X−1 Obs 0202401101
1
0.5
0.2
normalized counts s−1 keV−1
2
normalized counts s−1 keV−1
normalized counts s−1 keV−1
2
1
0.5
0.2
1000
100
10
0.1
0.1
1
0.8
1.2
ratio
1.2
ratio
ratio
1.2
1
0.8
2
1
0.8
5
2
5
2
5
Energy (keV)
Energy (keV)
Energy (keV)
CYGNUS X−1 Obs 0202401201
CYGNUS X−1 Obs 0202760201
CYGNUS X−1 Obs 0202760301
50
100
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
50
1000
20
10
5
20
10
5
10
2
2
1
0.8
1.2
ratio
1.2
ratio
ratio
1.2
1
0.8
2
1
0.8
5
5
5
Energy (keV)
Energy (keV)
Energy (keV)
CYGNUS X−1 Obs 0202760401
CYGNUS X−1 Obs 0202760501
Cygnus X−1 Obs 0500880201
100
50
20
10
5
50
normalized counts s−1 keV−1
normalized counts s−1 keV−1
normalized counts s−1 keV−1
50
20
10
5
20
10
5
2
2
1
0.8
1.2
ratio
1.2
ratio
ratio
1.2
1
0.8
1
0.8
5
5
Energy (keV)
2
Energy (keV)
5
Energy (keV)
Cygnus X−1 Obs 0500880201
Cygnus X−1 Obs 0610000401
normalized counts s−1 keV−1
normalized counts s−1 keV−1
200
100
50
20
10
100
10
5
1.2
ratio
ratio
1.2
1
0.8
1
0.8
2
5
Energy (keV)
2
5
Energy (keV)
Figure A.6: Peculiar sources data, model, model components and ratio data/model. These
sources can hardly be categorized in any of the described HMXBs standard groups.
122
Appendix A. The fits of the FeKα survey
Unclassified
AX J1749.1−2733 Obs 0510010401
normalized counts s−1 keV−1
0.1
0.01
10−3
ratio
1.2
1
0.8
2
5
Energy (keV)
Figure A.7: AX J1749.1-2733 data, model, model components, and ratio data/model. We have
no references for the luminosity class of the optical start, although the high absorption we observed points to a supergiant companion.
Table A.1: Parameters of the continuum.
Source
Obs ID
State
Model
Gain
χ2Red
P0
L1−10 keV
(1033 erg/s)
NH
(1022 cm−2 )
Γ
kT 1
(keV)
kT 2
(keV)
kT 3
(keV)
1A 0535+26
2S 1845-024
X Persei
X Persei
X Persei
AX J1820.5-1434
AX J1820.5-1434
RX J0146.9+6121
RX J0440.9+4431
RX J0440.9+4431
RX J0440.9+4431
RX J1037.5-5647
RX J1037.5-5647
SAX J2103.5+4545
SAX J2103.5+4545
Swift J045106.8-694803
V0332+53
0674180101
0302970801
0151380101
0151380101
0600980101
0511010101
0511010101
0201160101
0653660101
0653660101
0653660101
0550560101
0550560101
0149550401
0149550401
0679381401
0506190101
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
T N1
T5
T 13
T N1
T N1
T5
N1
T N1
T6
T 10
T N1
T9
T N1
T 10
T N1
N1
T5
-
0.87
0.73
1.34
1.41
1.28
0.53
0.60
1.31
1.16
1.21
1.24
1.02
1.00
1.26
1.24
1.06
0.60
8.70E-01
7.67E-01
3.68E-03
7.44E-04
1.06E-02
9.92E-01
9.77E-01
8.58E-03
8.55E-02
4.14E-02
2.54E-02
4.16E-01
4.94E-01
1.77E-02
2.69E-02
3.28E-01
7.58E-01
0.72+0.06
−0.07
8.82+2.80
−4.11
0.50+0.04
−0.04
0.58+0.08
−0.06
0.39+0.09
−0.11
7.85+2.64
−1.78
14.21+3.60
−2.03
0.90
0.78+0.06
−0.07
1.17+0.05
−0.05
1.22+0.12
−0.13
1.04
1.04
0.79+0.06
−0.04
0.83+0.11
−0.11
0.29+0.21
−0.15
0.83
1.40+0.06
−0.08
1.63+0.14
−0.11
1.34+0.13
−0.14
1.34+0.35
−0.22
1.61+0.10
−0.08
0.95+0.12
−0.13
1.09+0.38
−0.12
0.96+0.24
−0.22
1.07+0.13
−0.10
-
-
0152780201
0152780201
0083280101
0083280101
0083280101
0083280201
0083280201
EC
I/E
D
F
Q
F
Q
N1
N2
N2
N1
N1
N2
T N2
-
1.33
1.44
1.38
0.97
1.16
1.28
0.90
7.12E-03
5.65E-04
9.69E-03
5.73E-01
1.04E-01
1.23E-02
7.98E-01
0.97+0.03
−0.03
80.48+13.16
−13.17
11.31+1.49
−0.99
10.67+0.46
−1.36
15.49
16.17+0.54
−2.39
6.02+1.69
−0.79
0.37+0.04
−0.05
0.86+0.05
−0.05
1.20+0.17
−0.11
1.10+0.07
−0.09
1.22
1.12+0.10
−0.03
1.35+0.54
−0.25
1.24+0.05
−0.06
+10.76
4.72−0.92
1.67+0.05
−0.07
1.59+0.04
−0.04
1.08+0.12
−0.15
2.05+0.29
−0.23
1.34+0.08
−0.09
0.91+0.08
−0.06
1.30+0.05
−0.05
1.38+0.09
−0.07
4.42+0.37
−0.45
1.55+0.32
−0.22
2.46+0.06
−0.06
2.15+0.12
−0.21
0.40+0.08
−0.07
10.87+5.49
−3.41
2.28+0.23
−0.12
5.82+0.35
−0.35
1.53+0.08
−0.15
-
4U 1538-522
4U 1538-522
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
11.03+9.20
−6.39
15.93+2.54
−0.98
72.42+27.82
−23.16
72.49+27.81
−23.17
86.94+33.07
−27.62
12.05+14.36
−8.39
12.37+14.36
−8.63
7.65+0.11
−0.12
65.96+22.21
−18.86
66.00+22.13
−18.91
66.06+22.26
−18.89
11.58+0.33
−0.43
11.53+0.38
−0.38
805.54+459.66
−355.73
806.39+462.07
−355.79
860.26+118.45
−112.25
0.19+0.14
−0.09
-
-
20.04+7.24
−6.07
659.55+227.25
−192.86
161.58+62.67
−52.25
1790.09+645.31
−545.05
624.17+228.25
−190.74
1952.17+693.31
−589.18
619.13+224.34
−189.62
P
2.59+0.14
−0.26
Table A.1: Parameters of the continuum (continued).
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1907+09
4U 1907+09
4U 1907+09
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
GX 301-2
GX 301-2
IGR J16207-5129
IGR J16318-4848
IGR J16318-4848
IGR J16318-4848
IGR J16318-4848
IGR J16318-4848
0083280201
0083280301
0083280401
0083280401
0600950101
0600950101
0555410101
0555410101
0555410101
0206380401
0206380401
0405640201
0405640201
0405640301
0405640301
0405640401
0405640401
0405640901
0405640801
0555200401
0555200401
0402920201
0154750401
0154750401
0201000201
0201000201
0201000301
Q
Q
EC
I/E
EC
EC
F
Q
Q
Q
Q
EC
EC
Q
Q
Q
Q
Q
Q
F
Q
Q
Q
Q
Q
Q
Q
N2
N2
N1
N2
N1
N2
T N1
T N2
N2
T5
N1
T5
N1
T5
N1
T5
N1
N1
T7
N3
N3
N2
T5
N1
T5
N1
T5
slope
slope
-
1.00
1.14
1.22
1.13
1.18
1.04
1.10
0.99
1.13
1.33
1.28
0.71
0.85
0.91
1.04
1.09
1.00
1.14
1.15
1.29
1.82
1.17
1.12
1.05
0.89
0.92
0.73
4.69E-01
1.20E-01
1.30E-01
2.01E-01
6.97E-02
3.58E-01
1.89E-01
5.16E-01
1.30E-01
9.84E-03
2.30E-02
7.47E-01
5.79E-01
7.42E-01
3.67E-01
2.69E-01
4.77E-01
1.50E-01
1.44E-01
1.21E-02
9.70E-09
8.16E-02
2.31E-01
3.63E-01
7.12E-01
6.36E-01
6.83E-01
623.31+226.45
−190.77
718.06+259.67
−217.98
+10.00
24.12−8.93
+116.85
313.50−99.84
16.64+5.93
−5.01
16.77+5.98
−5.05
1288.31+6.16
−4.92
721.87+2.75
−3.28
718.59+3.08
−3.04
+389.71
522.62−279.02
+393.88
528.75−282.05
+2.35
2.77−1.61
2.82+2.59
−1.72
474.87+357.48
−254.37
479.86+360.57
−257.24
+136.89
181.41−97.28
+137.96
182.90−98.01
203.03+150.31
−108.96
+215.12
287.66−153.67
3412.87+5.83
−5.37
1334.78+3.30
−2.47
94.98+2.93
−−0.63
+22.83
11.11−10.28
+22.66
11.15−10.31
+23.87
11.56−10.70
+23.87
11.62−10.75
+8.26
3.61−3.36
8.16+2.55
−0.72
23.76+0.65
−0.69
11.66
97.57
11.36+0.15
−0.16
15.96+0.22
−0.20
2.33+0.12
−0.12
6.39+0.52
−1.09
9.10+1.03
−1.21
17.65+0.86
−0.68
24.14+1.41
−0.95
17.78+6.87
−3.72
19.87+21.04
−3.85
13.10+0.77
−0.66
18.55+1.18
−0.88
36.68+2.28
−2.02
44.56+3.18
−2.00
30.84+1.59
−2.01
31.45+1.71
−3.11
200.37+0.53
−0.57
228.22+0.49
−0.57
16.27
227.17+19.61
−13.52
232.66+12.48
−8.47
206.14+15.62
−13.46
214.22+11.00
−9.48
309.57+49.52
−76.63
1.03+0.08
−0.04
1.00+0.04
−0.05
0.51
0.58
0.70
0.70
0.75+0.08
−0.09
0.45+0.10
−0.10
1.14+0.02
−0.02
0.78+0.09
−0.07
−0.19+0.75
−0.28
0.77+0.09
−0.08
0.63+0.14
−0.09
0.97+0.18
−0.07
0.81
1.04
1.27
0.91+0.19
−0.17
1.08+0.20
−0.16
-
0.87+0.03
−0.04
1.18+0.07
−0.03
2.78+0.10
−0.09
4.94
2.71+0.11
−0.09
3.14+0.21
−0.18
0.23+0.00
−0.07
2.89+0.43
−0.32
2.67+0.39
−0.27
2.15+2.90
−0.46
2.62+0.16
−0.09
-
-
Table A.1: Parameters of the continuum (continued).
IGR J16318-4848
IGR J16318-4848
IGR J16318-4848
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16465-4507
SAX J1802.7-2017
Vela X-1
Vela X-1
XTE J0421+560
0201000301
0201000401
0201000401
0128531101
0128531101
0128531101
0201700301
0201700301
0556140101
0556140201
0556140301
0556140401
0556140501
0556140601
0556140701
0556140801
0556141001
0164561001
0206380601
0111030101
0111030101
0139760101
Q
Q
Q
Q
Q
Q
F
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
F
Q
Q
N1
T5
N1
T5
T 11
N1
N2
N2
N2
N2
N2
N2
N2
N2
N2
N2
N2
N1
N1
N2
N1
N1
-
0.76
1.12
1.16
0.99
1.03
1.03
0.89
0.93
0.99
1.05
1.09
1.21
0.84
1.04
1.24
1.15
0.93
0.58
1.28
1.35
0.87
0.70
6.55E-01
2.57E-01
2.12E-01
4.94E-01
4.23E-01
4.16E-01
8.13E-01
6.96E-01
5.15E-01
3.47E-01
2.28E-01
6.25E-02
8.57E-01
3.73E-01
5.34E-02
1.31E-01
7.05E-01
9.23E-01
3.18E-02
4.24E-03
8.65E-01
9.51E-01
AXJ1841.0-0536
IGR J00370+6122
IGR J00370+6122
IGR J00370+6122
IGR J00370+6122
0604820301
0501450101
0501450101
0501450101
0501450101
F
Q
Q
Q
Q
N3
T 10
T6
T9
T N1
-
1.39
0.84
0.91
0.97
0.84
1.54E-03
9.23E-01
7.71E-01
5.92E-01
9.24E-01
3.64+7.93
−3.39
7.52+15.62
−6.96
7.56+15.54
−7.00
+69.98
78.37−45.61
+70.79
79.52−45.66
+69.23
79.73−46.82
+724.71
918.65−515.23
417.44+328.31
−233.15
+382.82
474.86−267.46
920.37+729.24
−515.52
+705.54
890.63−498.80
547.92+436.81
−307.19
+363.15
443.72−251.54
1145.80+902.57
−640.15
+333.69
418.19−234.89
533.78+425.72
−299.24
+530.68
666.20−373.27
74.04+6.11
−8.28
624.71+16.86
−20.14
+126.06
547.45−112.92
323.43+73.19
−65.56
3.99+9.48
−3.84
501.72+104.41
−94.20
47.94+0.64
−0.60
47.70+0.75
−0.62
47.95+0.59
−0.68
48.01+0.58
−0.74
317.51+56.91
−14.71
182.75+20.34
−18.20
191.92+21.91
−8.77
24.97+5.10
−4.87
34.91+1.50
−4.88
35.37+8.66
−4.55
20.53+1.00
−1.57
27.23+0.80
−0.70
45.79+3.23
−1.53
32.21+1.27
−1.01
32.11+1.58
−0.92
38.80+1.67
−1.26
37.64+3.36
−2.65
33.80+0.92
−0.69
71.96+3.22
−2.43
34.90+1.67
−1.24
30.68+2.11
−1.52
72.18+24.40
−15.83
15.72+1.99
−1.35
34.63+1.64
−0.77
39.68+0.82
−0.82
36.88+13.28
−3.76
41.90+8.22
−3.98
0.87+0.13
−0.09
0.65
1.17+0.05
−0.05
1.00+0.17
−0.13
1.66+0.66
−0.11
1.19+0.45
−0.20
1.41+0.45
−0.28
0.45+0.12
−0.13
0.65+0.07
−0.05
0.73+0.13
−0.08
0.39+0.07
−0.06
0.43+0.09
−0.05
0.72+0.08
−0.07
0.67+0.17
−0.12
0.32+0.05
−0.03
0.17+0.08
−0.07
0.53+0.09
−0.06
0.52+0.12
−0.08
0.42+0.55
−0.50
0.96+0.19
−0.13
1.66+0.13
−0.07
1.65+0.05
−0.06
0.57+0.44
−0.26
1.27+0.07
−0.05
1.16+0.19
−0.14
2.55+0.45
−0.31
2.19+0.33
−0.23
50.75
-
-
-
1.16+0.09
−0.08
0.76+0.07
−0.05
3.08+0.09
−0.09
1.28+0.11
−0.07
4.20+0.69
−0.55
1.89+0.18
−0.10
-
-
Table A.1: Parameters of the continuum (continued).
IGR J11215-5952
IGR J11215-5952
IGR J16328-4726
IGR J16328-4726
IGR J16328-4726
IGR J16418-4532
IGR J16418-4532
IGR J16418-4532
IGR J16418-4532
IGRJ16479-4514
IGRJ16479-4514
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
IGR J17544-2619
IGR J17544-2619
IGR J17544-2619
IGR J18450-0435
0405181901
0405181901
0654190201
0654190201
0654190201
0206380301
0206380301
0405180501
0405180501
0512180101
0512180101
0554720101
0554720101
0554720101
0554720101
0554720101
0554720101
0561580101
0561580101
0561580101
0561580101
0561580101
0561580101
0148090501
0148090501
0148090501
0306170401
F
Q
Q
Q
Q
Q
Q
Q
Q
EC
I/E
F
F
F
Q
Q
Q
F
F
F
Q
Q
Q
F
F
F
F
T N1
T N1
T5
T9
N1
T5
N1
T5
N1
N1
N3
T5
T 11
N1
T5
T 11
N1
T5
T 11
N1
T5
T 11
N1
T5
T 11
N1
T N1
-
1.04
0.98
1.04
1.03
1.07
0.90
0.91
1.11
1.10
0.95
0.74
1.25
1.08
1.20
1.20
1.39
1.52
0.96
0.84
0.88
1.24
1.02
1.09
0.90
0.81
0.87
1.04
3.64E-01
5.39E-01
3.77E-01
3.93E-01
2.98E-01
7.30E-01
7.01E-01
1.99E-01
2.26E-01
5.92E-01
9.67E-01
5.66E-02
2.88E-01
9.43E-02
2.51E-01
1.26E-01
7.23E-02
6.04E-01
8.79E-01
8.05E-01
6.93E-02
4.33E-01
2.77E-01
7.16E-01
8.91E-01
7.83E-01
3.51E-01
639.50+5.64
−5.72
76.78+0.87
−0.97
32.07+45.82
−25.42
32.35+46.00
−25.65
+46.43
32.54−25.80
122.35+4.58
−4.50
124.60+4.33
−5.00
78.92+2.30
−2.38
82.17+2.23
−2.32
5.30+4.63
−3.10
278.33+225.37
−159.69
1.61+1.05
−0.76
1.71+1.11
−0.80
1.73+1.11
−0.82
0.16+0.12
−0.08
0.17+0.13
−0.09
0.18+0.13
−0.09
5.68+3.63
−2.67
5.96+3.85
−2.79
6.00+3.81
−2.82
1.14+0.74
−0.54
1.20+0.79
−0.57
1.22+0.78
−0.58
16.58+13.39
−9.11
17.77+14.39
−9.78
+14.45
18.12−10.01
131.25+1.41
−1.55
0.89+0.09
−0.06
1.02+0.21
−0.12
16.67+1.32
−1.25
22.21+1.58
−0.95
25.13+2.26
−1.62
10.73+1.42
−1.18
17.64+2.35
−1.53
3.92+0.42
−0.36
8.25+0.77
−0.57
8.21+1.62
−1.14
69.51
1.51+0.21
−0.18
3.81+0.31
−0.32
4.29+0.46
−0.33
1.89+1.12
−0.81
4.15+1.50
−1.08
5.08+1.59
−0.92
2.24+0.26
−0.23
5.12+0.29
−0.42
5.31+0.50
−0.41
2.28+0.29
−0.25
4.94+0.38
−0.39
5.52+0.63
−0.42
2.10+0.34
−0.28
4.58+0.40
−0.43
5.17+0.75
−0.48
2.18+0.31
−0.24
1.04+0.15
−0.09
0.63+0.21
−0.15
1.41+0.18
−0.13
1.35+0.22
−0.15
1.49+0.12
−0.10
1.75+0.32
−0.26
1.12
1.81+0.13
−0.10
2.49+0.46
−0.30
1.43+0.11
−0.09
1.82+0.15
−0.10
1.92+0.21
−0.14
0.84+0.26
−0.22
1.87+0.26
−0.23
1.22+0.14
−0.06
2.14+0.11
−0.10
4.48+0.45
−0.59
2.11+0.15
−0.13
1.77+0.07
−0.07
1.37+0.05
−0.05
10.46+3.51
−1.84
1.00+0.13
−0.12
3.92+2.55
−1.37
1.67+0.06
−0.06
34.27+44.58
−8.96
1.42+0.06
−0.05
10.85+3.85
−1.85
1.32+0.07
−0.06
8.60+3.08
−1.63
1.58+0.26
−0.13
-
-
Table A.1: Parameters of the continuum (continued).
IGR J18450-0435
IGR J18483-0311
IGR J18483-0311
IGR J18483-0311
0306170401
0406140201
0406140201
0406140201
Q
Q
Q
Q
T N1
T5
T9
N1
-
1.05
1.17
1.28
1.39
3.39E-01
1.98E-01
9.12E-02
4.16E-02
γ Cassiopeiae
γ Cassiopeiae
γ Cassiopeiae
γ Cassiopeiae
γ Cassiopeiae
γ Cassiopeiae
HD 110432
HD 110432
HD 119682
HD 119682
HD 157832
HD 161103
HD 45314
SAO 49725
SAO 49725
SS397
SS397
SS397
SS397
SS397
SS397
SS397
SS397
0201220101
0651670201
0651670201
0651670301
0651670401
0651670501
0504730101
0504730101
0551000201
0551000201
0551020101
0201200101
0670080301
0201200201
0201200201
0122700101
0122700101
0122700201
0122700201
0122700301
0122700301
0122700501
0122700501
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
T3
T4
T3
T3
T3
T3
N1
T N5
T2
T N3
T1
T2
T2
T2
T N5
T1
N1
T1
N1
T1
N1
T1
N1
slope
offset
offset
-
1.39
1.22
0.89
1.43
1.15
1.18
1.24
1.03
1.11
1.18
1.07
1.20
0.96
1.31
1.19
0.84
1.00
0.89
0.93
1.03
1.08
1.01
0.95
1.42E-03
3.62E-02
8.39E-01
4.40E-04
1.09E-01
6.73E-02
2.81E-02
3.78E-01
2.06E-01
1.06E-01
3.12E-01
9.76E-02
5.63E-01
7.39E-02
1.67E-01
7.90E-01
4.67E-01
6.88E-01
6.04E-01
4.21E-01
3.35E-01
4.49E-01
5.45E-01
22.69+0.37
−0.46
1.25+0.48
−0.39
1.26+0.49
−0.39
1.29+0.48
−0.41
0.24+0.05
−0.05
0.27+0.06
−0.05
0.27+0.06
−0.05
0.24+0.05
−0.05
0.33+0.07
−0.07
0.23+0.05
−0.05
0.21+0.06
−0.06
0.21+0.06
−0.06
0.16+0.01
−0.01
0.16+0.01
−0.00
0.08+0.00
−0.00
0.41+0.34
−0.23
0.30+0.14
−0.10
0.52+0.39
−0.27
0.55+0.41
−0.29
0.17+0.07
−0.05
0.17+0.06
−0.05
0.20+0.08
−0.06
0.20+0.08
−0.06
0.13+0.06
−0.04
0.14+0.06
−0.05
0.20+0.09
−0.06
0.19+0.09
−0.07
2.47+0.58
−0.27
4.64+1.11
−0.75
7.42+1.28
−0.83
11.40+0.53
−0.33
0.56
0.23+0.01
−0.02
0.42+0.05
−0.04
0.56
0.56
0.56
0.17+0.02
−0.02
0.11+0.01
−0.02
1.81
1.81
0.34
0.94
0.73
0.77
0.77
1.37
1.45+0.43
−0.27
1.37
1.37
1.37
1.37
1.37
1.37
0.79+0.66
−0.19
2.44+0.15
−0.14
1.52+0.02
−0.02
0.94+0.18
−0.15
1.63+0.15
−0.16
0.82+0.40
−0.37
1.73+0.24
−0.16
1.58+0.12
−0.11
1.68+0.13
−0.11
1.42+0.20
−0.16
1.74+0.23
−0.08
1.36+0.09
−0.10
2.07+0.25
−0.22
-
-
-
0.96
28.62+1.14
−0.92
0.74
0.64
0.66
0.66
8.69+1.16
−1.14
0.22+0.00
−0.05
0.12+0.00
−0.01
6.50+0.90
−0.73
1.03
1.03
0.82
1.97+0.92
−0.53
7.13+2.49
−1.14
10.03+7.04
−2.09
6.92+3.32
−1.23
+51.89
28.01−11.74
-
5.27
3.42
3.95
3.58
3.40
14.53+9.04
−3.50
7.05
13.29
+61.04
18.86−2.80
-
33.04
31.13
40.83
38.62
38.38
-
Table A.1: Parameters of the continuum (continued).
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS 5039
LS 5039
LS 5039
LS 5039
LS 5039
LS 5039
LS 5039
LS 5039
LS 5039
LS 5039
0207260101
0505980801
0505980901
0505981001
0505981101
0505981201
0505981301
0505981401
0151160201
0151160201
0151160301
0151160301
0151160301
0202950201
0202950201
0202950201
0202950301
0202950301
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
T N1
T N1
T N1
T N1
T N1
T N1
T N1
T N1
T 11
N1
T 11
T 10
N1
T 11
T 10
N1
T 11
N1
-
1.09
1.03
1.19
1.09
0.92
0.99
1.06
1.16
1.18
0.97
1.21
1.20
1.11
1.15
1.10
0.93
1.19
0.92
2.25E-01
4.04E-01
8.86E-02
2.44E-01
7.24E-01
5.09E-01
3.03E-01
1.03E-01
9.13E-02
5.76E-01
6.01E-02
6.97E-02
1.99E-01
1.14E-01
2.09E-01
6.99E-01
8.41E-02
7.12E-01
4U 2206+54
Cen X-3
Cen X-3
Cygnus X-1
Cygnus X-1
Cygnus X-1
Cygnus X-1
Cygnus X-1
0650640101
0111010101
0111010101
0202401101
0202401201
0202760201
0202760301
0202760401
Q
I/E
I/E
Q
Q
Q
Q
Q
T N1
T N4
N2
T N4
T N3
N1
N1
N1
-
1.19
1.28
1.27
1.12
1.13
1.49
0.96
1.72
5.04E-02
1.24E-02
1.39E-02
1.54E-01
1.39E-01
2.94E-03
5.82E-01
6.18E-05
6.27+0.08
−0.07
5.89+0.13
−0.12
6.31+0.20
−0.17
6.00+0.18
−0.16
6.39+0.16
−0.11
11.94+0.24
−0.18
9.37+0.18
−0.15
6.22+0.17
−0.12
5.00+0.58
−0.52
5.10+0.56
−0.52
5.07+0.61
−0.53
5.08+0.64
−0.54
5.15+0.59
−0.55
5.93+0.62
−0.58
6.02+0.65
−0.62
6.05+0.63
−0.59
3.50+0.43
−0.38
3.60+0.41
−0.38
294.46+42.83
−39.89
2242.27+477.52
−430.86
2240.60+477.65
−430.78
10191.40+2586.02
−2289.25
8572.45+2176.95
−1929.60
5061.89+1284.05
−1137.75
3453.30+877.54
−777.01
5009.50+1269.64
−1126.79
0.90
0.90
0.90
0.90
0.90
0.90
0.90
0.90
1.03
1.03
1.03
1.03
1.03
1.03
1.03
1.03
1.03
1.03
0.61+0.03
−0.03
5.96+0.04
−0.05
6.07+0.08
−0.09
3.85+0.74
−0.51
0.77+0.02
−0.01
0.86
0.86
0.86
1.98+0.05
−0.08
2.34+0.12
−0.13
2.47+0.17
−0.24
2.28+0.16
−0.25
2.08+0.09
−0.16
1.89+0.07
−0.11
2.07+0.11
−0.12
2.06+0.11
−0.17
1.60+0.04
−0.04
1.48+0.04
−0.04
1.54+0.03
−0.03
1.69+0.04
−0.04
1.01+0.04
−0.04
0.68
0.71
2.67+0.06
−0.04
2.61+0.05
−0.03
2.29+0.03
−0.01
2.07+0.01
−0.01
2.10+0.02
−0.01
2.46+1.65
−0.26
1.78+0.33
−0.15
1.85+0.34
−0.15
1.79+0.57
−0.16
2.17+2.48
−0.24
2.58+1.74
−0.31
2.07+0.34
−0.15
2.12+1.23
−0.22
11.87+2.26
−1.18
18.50+4.55
−2.57
2.01+0.42
−0.16
14.66+1.68
−1.51
2.14+0.26
−0.16
9.24+1.38
−0.93
-
0.76+0.16
−0.07
0.82+0.08
−0.06
-
-
1.67+0.04
−0.03
0.94
0.46+0.01
−0.01
0.43+0.00
−0.01
-
-
-
Table A.1: Parameters of the continuum (continued).
Cygnus X-1
Cygnus X-1
Cygnus X-1
Cygnus X-1
Cygnus X-1
AX J1749.1-2733
AX J1749.1-2733
0202760501
0500880201
0500880201
0500880201
0610000401
0510010401
0510010401
Q
D
D
Q
Q
Q
Q
N1
T N4
N2
N1
T N4
T5
N1
-
1.17
1.19
1.19
1.46
1.13
0.92
0.95
1.38E-01
6.15E-02
5.54E-02
3.73E-04
1.43E-01
6.80E-01
6.06E-01
4842.41+1232.95
−1091.06
+494.23
1899.40−430.07
+490.38
1899.66−434.10
+676.00
2679.69−600.04
3840.30+977.45
−868.00
+156.80
290.03−119.53
293.29+156.74
−121.29
0.86
4.53+0.26
−0.34
5.04+0.35
−0.45
0.71
6.18+1.11
−0.07
22.81+2.70
−2.25
31.98+3.03
−2.53
2.18+0.02
−0.02
1.43+0.01
−0.03
1.42+0.04
−0.03
1.55
1.60+0.03
−0.01
1.32+0.18
−0.16
0.39+0.02
−0.02
4.63
0.57+0.02
−0.01
2.26+0.18
−0.16
-
-
Table A.1: Parameters of the continuum, including observation ID, state of the source, model (see Table 4.2), artificial gain (offset or slope),
reduced χ2 , null hypothesis probability (P0 ), luminosity between 1-10 keV (L1−10 keV ), total NH (adding every absorption component), photon
index (Γ), and temperature of thermal components (kT i ): bbody, diskbb, bremss, mekal or cemekl (depending on the model). Parameters
frozen or unbounded are included without an error estimation. Therefore they are not used in the plots and the subsequent discussion. The
possible states are: quiescence(Q), flare(F), eclipse ingress/egress(I/E), eclipse(E) and dip(D).
-
Table A.2: FeKα parameters.
Source
Obs ID
State
1A 0535+26
2S 1845-024
X Persei
X Persei
AX J1820.5-1434
RX J0146.9+6121
RX J0440.9+4431
RX J1037.5-5647
SAX J2103.5+4545
Swift J045106.8-694803
V0332+53
0674180101
0302970801
0151380101
0600980101
0511010101
0201160101
0653660101
0550560101
0149550401
0679381401
0506190101
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
4U 1538-522
4U 1538-522
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1700-37
4U 1907+09
P0
Energy
(keV)
8.70E-01
7.67E-01
3.68E-03
1.06E-02
9.92E-01
8.58E-03
8.55E-02
4.16E-01
1.77E-02 6.42+0.06
−0.04
3.28E-01
7.58E-01
-
0152780201 EC 7.12E-03 6.38+0.02
−0.01
0152780201 I/E 5.65E-04 6.40+0.01
−0.03
0083280101 D 9.69E-03
0083280101 F 5.73E-01 6.48+0.06
−0.16
0083280101 Q 1.04E-01
0083280201 F 1.23E-02 6.45+0.01
−0.01
0083280201 Q 7.98E-01 6.44+0.04
−0.04
0083280301 Q 1.20E-01 6.45+0.03
−0.03
0083280401 EC 1.30E-01 6.49+0.02
−0.02
0083280401 I/E 2.01E-01 6.48+0.04
−0.03
0600950101 EC 6.97E-02 6.40+0.00
−0.00
0555410101 F 1.89E-01 6.43+0.03
−0.02
Narrow FeKα
Width
LFeKα
(eV)
(1041 γ/s)
54.00
0.00
24.24
0.02
90.60+11.35
−33.19
125.45+48.71
−49.29
87.78+72.25
−46.84
112.97+23.82
−25.76
0.00
38.57+4.33
−4.67
34.05
EW
(keV)
< 6.92E − 02
< 0.06
< 1.22E + 00
< 0.60
< 2.11E − 01
< 0.03
< 1.76E − 01
< 0.02
< 4.03E − 01
< 0.21
< 4.25E − 02
< 0.06
< 4.73E − 01
< 0.06
< 1.72E − 01
< 0.13
+0.01
3.87E + 00+3.80E+00
0.04
−0.01
−2.59E+00
< 1.80E + 01
< 0.20
< 1.14E − 01
< 2223.78
9.58E − 01+5.48E−01
−3.93E−01
3.42E + 00+4.16E+00
−1.58E+00
< 6.03E + 00
6.86E + 00+4.54E+00
−5.30E+00
< 1.23E + 01
2.86E + 01+1.13E+01
−1.31E+01
9.07E + 00+5.72E+00
−4.56E+00
8.08E + 00+6.26E+00
−3.82E+00
3.54E + 00+1.80E+00
−1.45E+00
7.28E + 00+3.73E+00
−3.67E+00
2.47E + 00+9.42E−01
−7.77E−01
6.04E + 00+1.10E+00
−1.59E+00
Signif. Energy
(σ)
(keV)
3.27
-
−
-
0.18+0.04
8.70 6.54+0.02
−0.04
−0.02
+0.02
0.03−0.01 6.46
−
< 0.27
0.03+0.01
3.56
−
−0.02
< 0.13
+0.00
0.11−0.03 19.07
−
0.11+0.02
7.25
−
−0.03
+0.02
0.07−0.02 8.15
−
1.25+0.27
12.50
−
−0.28
0.13+0.01
6.82
−
−0.04
+0.05
0.92−0.04 > 25.5 6.54+0.04
−0.03
0.04+0.01
8.55
−
−0.01
Broad Fe feature
Width
LFeKα
(eV)
(1041 γ/s)
−
271.65+20.73
−18.92
−
−
−
−
−
−
−
470.98+47.86
−35.92
−
−
-
EW
(keV)
Signif.
(σ)
−
-
-
+0.14
3.15E + 00+1.41E+00
−1.10E+00 0.92−0.13 19.95
−
−
−
−
−
−
−
−
−
−
−
−
−
−
+4.84E−01
+0.02
1.07E + 00−3.69E−01 0.19−0.01 21.33
−
−
-
Table A.2: FeKα parameters (continued).
4U 1907+09
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
EXO1722-363
GX 301-2
GX 301-2
IGR J16207-5129
IGR J16318-4848
IGR J16318-4848
IGR J16318-4848
IGR J16318-4848
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16320-4751
IGR J16465-4507
0555410101 Q 5.16E-01 6.43+0.02
−0.02
0206380401 Q 9.84E-03 6.40+0.01
−0.01
0405640201 EC 7.47E-01 6.41+0.05
−0.03
0405640301 Q 7.42E-01 6.42+0.02
−0.02
0405640401 Q 2.69E-01 6.41+0.01
−0.01
0405640901 Q 1.50E-01 6.42+0.72
−0.48
0405640801 Q 1.44E-01 6.43+0.01
−0.01
0555200401 F 1.21E-02 6.37+0.00
−0.00
0555200401 Q 9.70E-09 6.38+0.01
−0.01
0402920201 Q 8.16E-02 6.42+0.06
−0.05
0154750401 Q 2.31E-01 6.40+0.00
−0.00
0201000201 Q 7.12E-01 6.41+0.01
−0.01
0201000301 Q 6.83E-01 6.42+0.02
−0.01
0201000401 Q 2.57E-01 6.41+0.01
−0.01
0128531101 Q 4.94E-01
0201700301 F 8.13E-01 6.41+0.02
−0.01
0201700301 Q 6.96E-01 6.42+0.00
−0.01
0556140101 Q 5.15E-01 6.44+0.03
−0.07
0556140201 Q 3.47E-01 6.42+0.01
−0.01
0556140301 Q 2.28E-01 6.42+0.01
−0.01
0556140401 Q 6.25E-02 6.43+0.01
−0.02
0556140501 Q 8.57E-01 6.41+0.02
−0.02
0556140601 Q 3.73E-01 6.42+0.01
−0.01
0556140701 Q 5.34E-02 6.42+0.01
−0.01
0556140801 Q 1.31E-01 6.42+0.01
−0.01
0556141001 Q 7.05E-01 6.42+0.01
−0.01
0164561001 Q 9.23E-01
-
56.22+27.09
−46.41
46.25
124.04+55.59
−39.53
54.05
45.19+20.89
−43.01
17.75
50.70+23.37
−37.21
25.43+1.57
−6.61
0.14
0.00
35.48+6.79
−7.68
26.59
32.52
27.11
31.55
0.00
0.00
26.65
0.00
68.31+19.44
−23.31
0.00
34.48+11.53
−31.19
45.43+8.55
−29.65
52.95+20.85
−29.37
21.72
-
5.29E + 00+7.56E−01
−8.56E−01
1.80E + 01+1.67E+01
−1.09E+01
1.02E + 00+1.23E+00
−6.50E−01
1.36E + 01+1.44E+01
−8.66E+00
1.46E + 01+1.33E+01
−8.60E+00
2.99E + 00+7.11E+01
−1.81E+00
1.57E + 01+1.49E+01
−9.36E+00
2.79E + 02+1.05E+00
−7.95E+00
1.41E + 02+4.70E+00
−3.54E+00
4.19E − 01+3.09E−01
−1.77E−01
2.55E + 02+8.46E+02
−2.39E+02
1.05E + 02+3.07E+02
−9.88E+01
3.96E + 02+2.58E+03
−3.88E+02
4.56E + 01+1.51E+02
−4.31E+01
< 3.78E + 00
1.92E + 01+1.84E+01
−1.17E+01
7.80E + 00+7.17E+00
−4.60E+00
1.83E + 01+1.62E+01
−1.67E+01
2.92E + 01+2.56E+01
−1.80E+01
2.23E + 01+2.17E+01
−1.39E+01
1.64E + 01+1.60E+01
−9.98E+00
1.13E + 01+1.27E+01
−7.30E+00
3.72E + 01+3.25E+01
−2.29E+01
2.83E + 01+2.46E+01
−1.77E+01
1.67E + 01+1.61E+01
−1.01E+01
1.72E + 01+1.82E+01
−1.05E+01
< 5.27E + 00
0.05+0.01
−0.01
0.15+0.02
−0.03
2.49+0.71
−0.59
0.14+0.03
−0.03
0.23+0.03
−0.04
0.05+0.66
−0.01
0.17+0.03
−0.03
0.46+0.00
−0.02
0.44+0.03
−0.02
0.03+0.02
−0.01
2.47+0.25
−0.21
1.12+0.13
−0.13
2.28+1.47
−0.68
1.15+0.15
−0.16
< 0.28
0.13+0.01
−0.02
0.11+0.01
−0.01
0.22+0.01
−0.19
0.19+0.01
−0.03
0.15+0.02
−0.03
0.17+0.02
−0.02
0.15+0.03
−0.03
0.20+0.01
−0.04
0.41+0.04
−0.09
0.19+0.02
−0.02
0.15+0.03
−0.02
< 0.50
11.60
13.19
7.92
8.99
14.26
4.02
14.20
> 25.5
> 25.5
2.71
> 25.5
> 25.5
10.97
22.68
5.38
9.58
15.57
6.60
5.52
17.17
7.71
15.70
10.46
19.00
8.66
-
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
-
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
-
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
-
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
-
-
Table A.2: FeKα parameters (continued).
SAX J1802.7-2017
Vela X-1
Vela X-1
XTE J0421+560
0206380601
0111030101
0111030101
0139760101
AXJ1841.0-0536
IGR J00370+6122
IGR J11215-5952
IGR J11215-5952
IGR J16328-4726
IGR J16418-4532
IGR J16418-4532
IGRJ16479-4514
IGRJ16479-4514
XTE J1739-302
XTE J1739-302
XTE J1739-302
XTE J1739-302
IGR J17544-2619
IGR J18450-0435
IGR J18450-0435
IGR J18483-0311
+0.01
0604820301 F 1.54E-03 6.46+0.07
0.00
1.67E + 00+1.18E+00
−1.19E+00 0.02−0.01
−0.06
0501450101 Q 9.23E-01
< 2.97E − 01
< 0.06
+0.04
+52.30
+1.38E+00
0405181901 F 3.64E-01 6.48−0.05 62.01−45.59 4.16E + 00−1.45E+00 0.05+0.02
−0.02
< 1.16E + 00
< 0.13
0405181901 Q 5.39E-01
0654190201 Q 3.77E-01
< 4.94E − 01
< 0.09
0206380301 Q 7.30E-01
< 1.34E + 00
< 0.07
0405180501 Q 1.99E-01
< 6.39E − 01
< 0.05
+0.02
+5.75E−01
0.00
4.89E − 01−3.34E−01 0.45+0.11
0512180101 EC 5.92E-01 6.42−0.00
−0.14
0512180101 I/E 9.67E-01
6.40
0.00
1.72E + 00
0.03
0554720101 F 5.66E-02
< 3.81E − 02
< 0.21
< 1.83E − 02
< 2.27
0554720101 Q 2.51E-01
0561580101 F 6.04E-01
< 9.58E − 02
< 0.12
0561580101 Q 6.93E-02
< 2.35E − 02
< 0.17
0148090501 F 7.16E-01
< 2.39E − 01
< 0.12
0306170401 F 3.51E-01
< 2.24E + 00
< 0.13
< 4.27E − 01
< 0.14
0306170401 Q 3.39E-01
0406140201 Q 1.98E-01
< 4.23E − 02
< 0.27
γ Cassiopeiae
γ Cassiopeiae
γ Cassiopeiae
γ Cassiopeiae
γ Cassiopeiae
0201220101
0651670201
0651670301
0651670401
0651670501
Q
F
Q
Q
Q
Q
Q
Q
Q
3.18E-02
< 1.36E + 01
< 0.14
+3.78E+00 0.07+0.02
55.91
7.27E
+
00
4.24E-03 6.45+0.02
−0.02
−0.01
−2.56E+00
+18.62 4.00E + 00+1.50E+00 0.06+0.01
8.65E-01 6.45+0.01
45.04
−0.01
−25.41
−1.15E+00
−0.01
9.51E-01 6.43+0.03
0.00
2.78E − 01+1.61E+00
0.15+0.21
−0.03
−0.04
−2.69E−01
1.42E-03
3.62E-02
4.40E-04
1.09E-01
6.73E-02
6.43+0.01
−0.01
6.39+0.03
−0.02
6.44+0.26
−0.03
6.41+0.02
−0.02
6.39+0.02
−0.03
0.00
23.94
6.39
1.50
0.00
6.86E − 04+2.62E−04
−1.95E−04
1.08E − 03+5.03E−04
−4.24E−04
1.00E − 03+2.86E−03
−3.77E−04
1.53E − 03+6.94E−04
−4.17E−04
8.55E − 04+3.70E−04
−3.18E−04
0.03+0.00
−0.00
0.04+0.01
−0.01
0.04+0.11
−0.01
0.05+0.01
−0.00
0.04+0.01
−0.01
9.75
−
−
−
−
+115.79
+1.91E+01
+0.08
+0.15
19.86 6.62−0.11 1220.66−76.84 4.29E + 01−1.24E+01 0.75−0.09 > 25.5
2.65 6.56+0.21
402.38+4955.35
7.08E − 01+2.67E+01
0.43+4.28
3.51
−0.13
−58.40
−6.86E−01
−2.68
2.78
3.49
6.36
1.48
-
−
−
−
−
-
−
−
−
−
-
−
−
−
−
-
−
−
−
−
-
-
13.80
7.82
7.41
10.26
7.91
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
−
-
Table A.2: FeKα parameters (continued).
HD 110432
HD 119682
HD 157832
HD 161103
HD 45314
SAO 49725
SS397
SS397
SS397
SS397
0504730101
0551000201
0551020101
0201200101
0670080301
0201200201
0122700101
0122700201
0122700301
0122700501
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
2.81E-02 6.44+0.03
−0.04
2.06E-01
3.12E-01
6.40
6.40
9.76E-02
5.63E-01 6.39+2.77
−0.65
7.39E-02
7.90E-01
6.88E-01
4.21E-01
4.49E-01
-
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS I +61 303
LS 5039
LS 5039
LS 5039
LS 5039
0207260101
0505980801
0505980901
0505981001
0505981101
0505981201
0505981301
0505981401
0151160201
0151160301
0202950201
0202950301
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
Q
2.25E-01
4.04E-01
8.86E-02
2.44E-01
7.24E-01
5.09E-01
3.03E-01
1.03E-01
9.13E-02
6.01E-02
1.14E-01
8.41E-02
4U 2206+54
Cen X-3
Cygnus X-1
Cygnus X-1
0650640101
0111010101
0202401101
0202401201
Q
I/E
Q
Q
5.04E-02
< 1.21E + 00
< 0.03
+0.01
+17.30
+2.06E+01
1.24E-02 6.41−0.01 42.06−13.71 5.77E + 01−1.45E+01 0.14+0.02
−0.01
1.54E-01 6.59+0.05
0.00
6.22E + 00+4.51E+00
0.01+0.01
−0.04
−0.01
−3.50E+00
1.39E-01
−
−
−
−
-
0.14
0.00
0.00
0.00
-
1.21E − 03+7.92E−04
0.05+0.01
−0.02
−5.97E−04
< 7.53E − 03
< 0.63
+0.06
1.32E − 03+8.17E−04
−9.18E−04 0.10−0.07
+9.96E−03
4.68E − 03−4.19E−03 0.09+0.07
−0.07
5.45E − 03
0.17
< 3.26E − 02
< 0.71
< 1.00E − 02
< 0.42
< 1.00E − 02
< 0.43
< 1.10E − 02
< 0.57
< 2.69E − 02
< 1.31
< 2.16E − 02
< 5.80E − 02
< 8.59E − 02
< 9.20E − 02
< 4.72E − 02
< 1.16E − 01
< 5.54E − 02
< 6.98E − 02
< 4.72E − 02
< 1.06E − 01
< 4.72E − 02
< 8.39E − 02
< 0.04
< 0.12
< 0.14
< 0.17
< 0.08
< 0.10
< 0.06
< 0.12
< 0.10
< 0.22
< 0.08
< 0.28
5.36
2.47
2.18
1.98
-
−
−
−
−
-
−
−
−
−
-
−
−
−
−
-
−
−
−
−
-
-
-
-
-
-
-
-
16.08
−
−
−
−
+100.30 1.66E + 02+8.33E+01 0.44+0.09 18.93
3.46 6.42+0.06
873.82
−0.09
−110.40
−4.86E+01
−0.06
+156.75 1.31E + 02+7.54E+01 0.46+0.13 19.49
6.50+0.07
−0.09 944.58−84.66
−4.24E+01
−0.06
Table A.2: FeKα parameters (continued).
Cygnus X-1
Cygnus X-1
Cygnus X-1
Cygnus X-1
Cygnus X-1
Cygnus X-1
Cygnus X-1
AX J1749.1-2733
0202760201
0202760301
0202760401
0202760501
0500880201
0500880201
0610000401
0510010401
Q
Q
Q
Q
D
Q
Q
Q
2.94E-03
5.82E-01
6.18E-05
1.38E-01
6.15E-02
3.73E-04
1.43E-01
6.80E-01
6.55+0.01
0.00
−0.02
6.54+0.02
0.00
−0.02
+0.03
6.48−0.01
0.00
6.53+0.02
0.00
−0.02
−
−
6.41
0.00
+142.20
6.46+0.03
−0.09 82.68−34.87
-
1.07E + 01+4.57E+00
−3.49E+00
5.54E + 00+2.72E+00
−2.11E+00
6.01E + 00+3.04E+00
−2.43E+00
1.12E + 01+5.53E+00
−3.99E+00
−
4.94E + 00+1.25E+00
−1.11E+00
+6.63E+00
7.05E + 00−2.98E+00
< 6.11E + 00
0.01+0.00
−0.00
0.01+0.00
−0.00
0.01+0.00
−0.00
0.01+0.00
−0.00
−
0.02
0.02+0.01
−0.01
< 0.12
5.82
7.72
7.47
8.97
6.95
5.09
-
6.55+0.02
−0.03
6.41+0.04
−0.04
6.47+0.03
−0.03
6.46+0.03
−0.06
6.36+0.21
−0.22
−
−
-
1101.52+54.39
−35.58
1243.17+66.00
−52.76
1199.67+57.79
−37.68
1297.32+91.10
−52.59
595.01+365.26
−181.09
−
−
-
4.82E + 02+1.80E+02
−1.26E+02
3.29E + 02+1.19E+02
−9.09E+01
4.64E + 02+1.65E+02
−1.18E+02
6.03E + 02+2.45E+02
−1.67E+02
2.13E + 01+1.59E+01
−1.10E+01
−
−
-
Table A.2: Fit parameters of FeKα and broad Fe feature detected. We present the centroid energy (keV), the width (eV), the luminosity
(photons/s), EW (keV) and the statistical significance of the Gaussian lines. In the cases where we don’t detect FeKα (σ sign < 2), we give the
upper limits of EW and the luminosity. Parameters frozen or unbounded are included without an error estimation. Therefore they are not used
in the plots and the subsequent discussion. The possible states are: quiescence(Q), flare(F), eclipse ingress/egress(I/E), eclipse(E) and dip(D).
0.65+0.07
5.82
−0.04
0.61+0.05
>
25.5
−0.04
+0.06
0.61−0.02 4.02
0.83+0.11
> 25.5
−0.06
+0.03
0.09−0.03 4.99
−
−
-
Appendix B
The fits of the supergiant donors
In this Appendix we show the entire spectra used for the fits described in Part III. First,
we show the fit of IGR J17544-2619 and second, the fit of Vela X-1.
135
136
Appendix B. The fits of the supergiant donors
B.1 IGR J17544-2619
Normalized flux
HeI
O II
O II
HeI
Hγ
O II
O II
He II
HeI
Si IV
HeI
Hδ
OII
OII
NII
SiIV
HeI
Hε
HeI
1.2
1.0
0.8
0.6
4000
4100
4200
4300
4400
4500
o
λ/A
Normalized flux
HeI
Hβ
HeI
He II
NIII
NIII
OII
OII
SiIII
SiIII
He II
NII
OII
CIII
1.2
1.0
0.8
0.6
4500
4600
4700
4800
4900
5000
5300
5400
5500
5900
6000
o
λ/A
Normalized flux
He II
S III
HeI
HeI
1.2
1.0
0.8
0.6
5000
5100
5200
o
λ/A
Normalized flux
HeI
SiIII
CIII
OIII
1.2
1.0
0.8
0.6
5500
5600
5700
5800
o
λ/A
Figure B.1: Optical spectrum of IGR J17544-2619 (blue), and the best fit model (red).
B.1. IGR J17544-2619
137
Normalized flux
He II
He II
He II
He II
He II
1.2
1.0
0.8
0.6
6000
6100
6200
6300
6400
6500
o
λ/A
Normalized flux
HeI
HeI
He II
Hα
He II
1.2
1.0
0.8
0.6
6500
6600
6700
6800
6900
7000
7300
7400
7500
o
λ/A
Normalized flux
HeI
HeI
He II
1.2
1.0
0.8
0.6
7000
7100
7200
o
λ/A
Figure B.1: continued.
Normalized flux
Normalized flux
1.0
0.8
1.0
0.8
1.0
0.8
8500
15000
1.0
0.8
19500
He II
H I.....22 - H I......4
H I.....17 - H I......3
H I.....16 - H I......3
H I.....15 - H I......3
He II
H I.....21 - H I......4
H I.....14 - H I......3
H I.....20 - H I......4
20000
H I.....13 - H I......3
H I.....19 - H I......4
H I.....12 - H I......3
H I.....18 - H I......4
H I.....11 - H I......3
H I.....17 - H I......4
9000
15500
20500
H I.....16 - H I......4
HeI
H I.....10 - H I......3
HeI
HeI
12000
H I.....15 - H I......4
21000
H I......9 - H I......3
He II
HeI
H I.....14 - H I......4
9500
16000
21500
HeI
HeI
H I......8 - H I......3
HeI
HeI
H I.....13 - H I......4
HeI
H I......7 - H I......4
He II
o
12500
22500
16500
10000
H I.....12 - H I......4
HeI
H I......7 - H I......3
He II
HeI
HeI
HeI
23000
H I.....11 - H I......4
He II
HeI
H I......5 - H I......3
HeI
I Band
10500
J Band
H Band
17000
23500
K Band
HeI
HeI
HeI
HeI
HeI
11000
HeI
13000
17500
24000
HeI
H I.....10 - H I......4
HeI
H I......6 - H I......3
NIII
HeI
HeI
Appendix B. The fits of the supergiant donors
λ/A
o
λ/A
o
λ/A
22000
o
λ/A
Figure B.2: Infrared spectrum of IGR J17544-2619 (blue), and the best fit model (red).
HeI
HeI
Normalized flux
138
Normalized flux
NV
CIII
CIII
O VI
Al V
O III
C III
Si IV
O IV
O III
O III
O III
NIV
O III
NIV
O III
Si IV
O III
N IV
NIV
O VI
O VI
Si III
CIII
N IV
Si III
Si III
Si III
Si III
CIII
N IV
CIV
O III
O III
O III
C II
Al V
C II
C II
O IV
N III
O III
O III
O III
O III
O III
NIII
NIII
FeII
FeIII
C IV
C IV
1350
N III
1550
1750
1950
139
1900
Si III
SV
OV
Al V
O III
B.2. Vela X-1
Al V
O III
C III
C III
Al V
P IV
O III
NV
NV
1.4
1.2
1.0
NV
1300
o
λ/A
1500
o
λ/A
1700
o
λ/A
o
λ/A
N IV
C II
Si IV
Si IV
Al III
0.8
N IV
NV
O III
Lyα
NV
Al V
O III
O II
O II
Si IV
N II
O II
Si IV
C III
N VI
OV
Mg III
0.6
OV
N III
NIV
N III
NV
OV
0.4
1850
OV
OV
O III
Si III
Al IV
OV
OV
OV
C III
Al IV
Al V
C III
C III
C III
Al IV
Si III
NIV
C IV
Al IV
N IV
NIV
Al IV
1250
1450
1650
C IV
He II
O VI
OV
C III
OV
0.2
0.0
1200
1.4
1.2
1.0
0.8
0.6
0.4
0.2
0.0
1400
1.4
1.2
1.0
0.8
0.6
N IV
O III
O III
N III
O II
O II
O II
O IV
Al V
O III
NIII
NIII
O II
C III
OV
O III
O III
O III
N IV
C III
O II
C III
C III
Al V
Al IV
C IV
O III
C III
OV
1400
1600
1800
O II
Si IV
Si IV
0.4
NV
NV
C III
O III
O II
C III
Al III
NV
Mg III
NV
Al III
N VI
C II
O III
O IV
O III
C III
O IV
C III
Mg III
O III
Al III
O VI
Mg III
Al IV
Si IV
Si III
Al III
O II
O II
N II
Mg III
C III
Al IV
N III
P IV
Normalized flux
Al III
Al IV
Al IV
O II
O VI
O III
Si III
O III
0.2
0.0
1600
1.4
1.2
1.0
0.8
0.6
0.4
0.2
0.0
1800
Figure B.3: Ultraviolet spectrum of Vela X-1 (blue), and the best fit model (red).
N III
N III
NIV
NV
NIV
Normalized flux
X-1
Normalized flux
B.2 Vela
Normalized flux
Si III
Si III
Si III
O III
O III
O III
Al V
O III
C IV
O III
Al V
C IV
O III
Si III
Al V
Si III
Si III
Al V
OV
Al V
Al V
Al IV
Al III
C III
NV
Si IV
Normalized flux
Normalized flux
HeI
O II
4500
4600
N II
5200
5700
CIII
AlIII
o
4100
O II
O II
HeI
O II
OII
HeI
NII
NIII
NIII
HeI
HeI
NII
NII
5300
5800
OII
OII
NII
OII
SiIV
OII
Hδ
Si IV
HeI
OII
HeI
OII
4200
4800
HeI
O II
O II
He II
Si IV
NII
NII
NII
HeI
Hβ
5400
O II
4300
4900
5900
CII
He II
HeI
OII
OII
CIII
Hγ
O II
O II
O II
O II
4400
5000
5500
6000
HeI
Appendix B. The fits of the supergiant donors
O II
He II
OII
HeI
HeI
OII
H......8 - H I......2
OII
OII
HeI
Hε
OII
O II
O II
N II
N II
N II
N II
N II
NIII
NIII
OII
O II
λ/A
4700
o
λ/A
o
λ/A
o
λ/A
SiIII
NII
OII
CIII
NIII
AlIII
NIII
He II
SiIII
SiIII
4000
N II
N II
OIII
SiIV
H I.....11 - H I......2
OII
OIII
SiIII
H I.....10 - H I......2
SiIII
HeI
H I......9 - H I......2
3900
5100
5600
OIII
1.2
3800
N II
HeI
N II
O II
O II
N II
HeI
1.0
O II
0.8
0.6
1.2
1.0
0.8
0.6
4400
1.2
1.0
0.8
0.6
5000
1.2
1.0
0.8
0.6
0.4
0.2
5500
Figure B.4: Optical spectrum of Vela X-1 (blue), and the best fit model (red).
N II
N II
N II
HeI
Normalized flux
140
Normalized flux
B.2. Vela X-1
141
Normalized flux
N II
He II
He II
He II
He II
He II
1.2
1.0
0.8
0.6
6000
6100
6200
6300
6400
6500
o
λ/A
He I
1.0
0.8
Hα
Normalized flux
HeI
He I
1.2
0.6
6500
6600
6700
6800
o
λ/A
Figure B.5: continued.
6900
7000
7100
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Agradecimientos
Quiero dar mi sincero agradecimiento a todas las personas que han hecho posible
la realización de este trabajo.
A mis padres y hermanos, que me apoyasteis para que me dedicara a lo que
siempre había querido ser.
A mis directores de tesis, Jose Miguel y Guillermo, por toda la confianza que me
habéis demostrado y vuestro ejemplo como investigadores.
Al resto del grupo, Silvia y Jose Joaquín, por darme vuestra ayuda y paciencia.
Al grupo de Potsdam, por vuestra acogida y la oportunidad de conocer otra manera
de vivir y trabajar.
A mis compañeros de piso estos años, de despacho, de terrazas de bar, de barrio,
de bici, de escalada... a los parrillitos canarios, a los lolos y lolas del IAC, a los que
almuerzan pulguita a las 11, a los potsdamitas... Sois muchos pero espero poder
conservar vuestra amistad muchos años.
Y a Irantzu, que ha endulzado todos los días de esta tesis.
163
Even if I knew that tomorrow
the world would go to pieces,
I would still plant my apple tree
Martin Luther King