uddipan kasi talk

8/2/2019 Uddipan KASI Talk

http://slidepdf.com/reader/full/uddipan-kasi-talk 1/100

Properties of

High Mass X-ray Binary Pulsars :

emphasis on the Stellar Windof the companion

Uddipan Mukherjee (T.I.F.R., India)

Advisor : Prof. Biswajit Paul (T.I.F.R. & R.R.I., India)



Outline of the Talk

X-ray Binaries : Pulsars & High Mass X-ray Binaries

Stellar Winds

Detectors & Analysis

Orbital Phase Spectroscopy of HMXBs

4U 1538-52

GX 301-2 OAO 1657-415

Vela X-1

Be/X-ray binary 3A 0535+262 in quiescence

Spectral studies in the high and low states

Cen X-3

2S 0114+650 Future Work



What are X-ray binaries?

Most bright X-ray

sources are in binaries

Compact object issucking material fromthe other star

Systems composed of 2 stars

The other star iscalled the companion

Binaries



According to the nature of the companion

If Mcomp

≤ Msun

Low Mass XRBs

If Mcomp

> Msun

High Mass X-ray Binaries

Classification of XRBs



Compact object is a NS with strongmagnetic field (1012 G)

Matter can only move following themagnetic field lines

Accumulation of plasma on the NSis called Accretion

X-rays make way out the magneticfield in a narrow beam

X-ray pulsars

The radiation is in pulses and theNS is an X-ray pulsar



High Mass X-ray Binaries

HMXRB contain supergiants (O &B) as companion & a NS(or a BH : Cyg X-1)

Companion is very bright (104 –105 × Lum of Sun)

Companion can be viewed in IR, optical or UV

NS can only be seen in X- rays

Most HMXRB are X-ray pulsars



As the supergiant expands, it increases in

brightness

The outer layers of its atmosphere push outa stream of particles

NS can pick up some of these particles andstart accretion in X rays



What is a supergiant ?

Brightest & largest kind of stars

Radii : 20 to several hundred times of Sun

Red (Betelguese) Blue (Rigel)

Both types can explode as Supernovae



Red Giant

Star of at least 15 solar mass & exhausting H

He - C burning expands even larger

Red Supergiant

Emanates a vigorous Stellar Wind

Lose extended atmospheres

smaller, hotter Blue Supergiants



What are Stellar Winds ?

Stars emit not only radiation but also particles

The emission of particles is called the Stellar Wind

Continuous phenomena, not episodic outbursts

In a star like the sun, the wind arises from the corona

In hotter stars, the high radiative flux, drives the wind

Primarily by means of line scattering



2 important parameters of a Stellar Wind derived fromobservations :

Mass Loss Rate : M & Terminal Velocity : v∞

A star like Sun loses about 10-14 solar mass yr-1

from winds blowing with 700 km s-1

Hot luminous stars exhibit stronger winds blowing

at speeds up to 2000 km s-1

And loses up to 10-5 solar mass yr-1 via winds

They are important to know the mechanism of wind generation

•



Line Driven Winds

Winds of luminous hot stars are driven by absorption in spectral lines

Hot stars emit bulk of their radiation in UV

Radiative acceleration in the winds of hot stars

mainly by the absorption & re-emission of UV photons

resonance lines of ions of abundant elements

C, N, O & Fe-group elements

Winds from quasars may also be radiation driven



Resonance line of N IV 765 Å

N+++ ion absorbing it increases its velocity by (h mc) cm s-1

To accelerate a single N+++ ion to the terminal velocity of 2000 km s-1 ,

we need 5

absorptions

In a plasma, momentum shared with other constituents

Hence, on an average : 1011

absorptions per ion to boost the velocityto the terminal value

The terminal velocity of a wind is reached within a few stellar radii

For example



To solve the Eqn. Of Motion for a stationary,

time independent wind

Assumptions :

Photosphere as a point source : radiation in the radialdirection

A photon emitted from the photosphere is absorbed by aline transition in the wind in only a narrow interactionregion

The wind is isothermal and behaves like a perfect gas

Then, the solution is :

v(r) = v ∞(1 – R/r)0.5

where R is the stellar radius

S o b o

l e v A

p p

r o x i m a

t i o n

Velocity Profile of Line Driven Winds

CAK Wind Profile



Observations & Analysis

● To go above the Earth's atmosphere ●

Possible via balloons, rockets & satellites● Data from ASCA, BeppoSAX, Chandra,

RXTE & XMM-Newton

X-rays from Stars

Satellite

Image Spectrum Light Curve

X-ray emissionmechanisms !!

Any Periodicities ??

Position



RXTE – PCA : Large Area (1000 sq. cm. @ 6 keV)

Broad Energy Band (2 – 60 keV)Time Resolution : 1 sEnergy Res : 18% @ 6.0 keV

A very brief description

ASCA : 0.4 – 10 keV

Energy Res : 2% @ 6 keV Area : 100 sq. cm. @ 6 keV

BeppoSAX : 0.1 – 10 keV

Energy Resolution : 8% @ 6.0 keV Area : 150 sq. cm. @ 6 keV



........Continued

Chandra :

ACIS : 0.1 – 10 keV

HETG : Energy Res : 0.5% @ 6 keV Area : 30 sq. cm. @ 6 keV

XMM-Newton :

PIC : 0.1 – 15 keV Area : 800 sq. cm. @ 6 keV (PN)

Energy Res : 2% @ 6 keV

Both have high earth orbits which helps toprovide long, continuous observations



A Glimpse of the Spectral Analysis

Assume a model for the incident spectrum

Generally for pulsars : power law (PL), black-body (BB) &

high energy cut-off describe their continuum

Any emission lines are modeled as gaussian features

Exponential line of sight absorption

Incident spectrum + detector response = observed spectrum

try to match the observed spectrum by varying the parameters of

the assumed model



Lineof sight

NS Eccentricorbit

Stellar Wind

4r2(r)v(r) = M

NH = ∫(x)dxX : path

Schematic diagram for HMXB pulsars



What we have done ............

Evaluated the Column Densities by a numerical integration along the line of sight from the NS to the observer at infinity as the NS traversed the

elliptical orbit

CAK velocity profile was assumed

The observer was on a diferent plane than the orbit

An assumed inclination angle was taken

Binary parameters , Mass-loss rate & Terminal velocity were known



Our Work

Part I

U. Mukherjee & B. Paul 2004, A&A,427, 567

U. Mukherjee et al., 2006a

U. Mukherjee, B. Paul & S. Naik,2006b

i i h h fl i i d



4U 1538-52 : a nice system with smooth outflowing wind

Properties

Orbital Period 3.73 d

Eclipse duration 0.6 d

Eccentricity 0.18

Distance 5.5 kpc

X-ray Luminosity 1036 erg s-1

Companion Star : B-type supergiant (Parkes et al. 1978)

Mass-loss rate 10-6 solar mass yr-1

v ∞

: 1000 km s-1

Pulse Period 529 s

Thus, it is a suitable candidate to study the windstructure of the companion star



Observations : We had proposed toobserve the source with RXTE-PCA

P.I. : Prof. B. Paul

from 2003-07-31 to 2003-08-07

out-of-eclipse phases for 2 binary orbits

25 data segments : exposure of (1.5 – 6.0) ks



Set III

1998-07-29 to 1998-08-01 with

BeppoSAX covering one binary orbit

40 data segments : exposure of

1.7 – 4.0 ks

Set II

1997-01-01 to 1997-01-05 with

PCA

14 datasets : (1.5 – 10.0) ks

on-source time

.....supplemented with archival data



Spectral Analysis

Backgrnd sub source spectra in 3 – 20 keV

Model : P L + exp Cut-Off + line-of-sightexp absp + gaussian @ 6.4 keV

All spectral parameters kept free for 2003dataset

Line centre and width frozen for 1997dataset

MECS and LECS data fitted in (1.8 – 10.0)keV & (0.3 – 4.5 ) keV

No high-energy cut-off for SAX

The soft excess at 0.1 keV was too faint

(0.8 – 1.5) : RXTE & (0.6 – 1.3) : SAX

LECS

MECS

RXTE-PCA

Spectral Parameters



Spectral Parameters

(1.0 – 1.5), lack of any

binary phase modulation

Similarly forE

c: RXTE-PCA 14 keV

& Ef 7 keV

Slightly different from Ec 16

keV &E

f 10 keV : BeppoSAX (Robba et

al. 2001)

Fluorescent Fe--line fluxmeasured with the averagespectrum taken over 2 -- 3 ks

Does not show any considerable variation along the orbit



Variation of Column Density with orbital phase

Column density : smooth variation over the orbitalphase

Model is consistent withobservations

2 HMXB pulsars showedsimilar increase in columndensity pattern near eclipse :

X1908+075 (Levine et al.2004) & SMC X-1 (Woo et al.

1995)

They may also have isotropic wind pattern from thecompanion stars



4U 1538-52 has a moderate X-ray luminosity : 1036 erg s-1

No significant perturbation in the stellar wind acceleration

4U 1700-37 (Haberl et al. 1989), 4U 1907+09 (Robertset al. 2001) & GX 301-2 (Leahy 1991, Pravdo & Ghosh2001) :

a simple spherical wind is not sufficient to explain thecolumn density profile



Conclusions for 4U 1538-52

Continuum X-ray spectrum is hardly affected by the NSrevolution

NH

shows a smooth variation over the orbital phase

A spherically symmetric stellar wind from the companion star mayproduce the observed orbital dependence of N

Hfor certain range of

the orbital inclination

Orbital phase resolved NH

measurements can be an

independent way to estimate the orbital inclination,especially for non-eclipsing binaries

GX 301 2 t ith l i d t t



GX 301-2 : a system with clumpy wind structure

• Orbital period ~ 41.5 d

• Eccentricity ~ 0.46

• Distance ~ 5.3 kpc

• X-ray Luminosity ~ 1035--37 erg s-1

• Companion Star : B-type Supergiant (Parkes et al. 1980)

• Mass-loss Rate ~ (3-10) × 10-6 solar mass yr -1

• v

∞~ 400 km s-1

• Pulse Period ~ 680 s



So, a large mass-loss rate & unusually low wind

velocity result in formation of clumped blobs of matter

Also, a large absorption column density is created

Thus, a variable luminosity & an eccentric orbit provide

a good site to probe the wind structure of the companion

A hi l RXTE Ob ti



Archival RXTE Observations

Set I : from 1996-10-05 to 1996-15-06 :

17 data segments

did not cover 0.85 - 0.98 phases

useful obs duration : 34 ks

Set II : from 2000-12--10 to 2000-19-11 :

39 data segments

almost a full phase coverage

useful obs duration : 262 ks

Choice of the spectral model



Choice of the spectral model

Her X-1 : Endo et al. (2000) : absorption had two components

One component absorbs the entire spectrum while the otherabsorbs it partially

Spectrum fitted : Partial Covering Absorption Model (PCAM)

PCAM : 2 different power laws with same photon index butdifferent normalisations & absorbed by different column densities

Endo et al. (2002) & Saraswat et al. (1996) : ASCA spectra of GX 301-2 with the PCAM + Fe K-alpha, K-beta emission lines and anFe absorption edge

The Chandra-HEG spectra of GX 301-2 was also fitted with the PCAM



A broad compton back-scattered peak @ 6.3 keV (660 eV FWHM) in the Chandra spectrum (Mukherjee & Paul, 2003& Watanabe et al. 2003)

For RXTE, we used the model used by Endo et al. (2002) + ahigh energy exponential cut-off

The back-scattered peak was not included in the RXTEspectrum due to the moderate energy resolution

Chandra

Wavelength (Ang)

C o u n

t s

/ s / A

n g

6.3 keV



RXTE Spectral Results

Almost all datasets : goodReduced between 0.6 and 1.6for 44 degrees of freedom (d.o.f.)

The fits with poor Reduced

showed wavy residuals with dips

around 10, 20 and 30 keV

Not Clear if it is a systematicphenomenon for this source

Systematic deviations at 20 &40 keV observed with Beppo-SAX also (Orlandini et al. 2000)

Poor

Good

The Continuum Parameters



10 spectral parameters varied : NH1

NH2, , the 2 normalizations, the 2iron line intensities, edge depth, E

c

& Ef

5 kept frozen : the iron-lineenergies & their FWHM along with

the edge energy (at the values of Mukherjee & Paul 2003)

Photon Index

Cut-off Energy

e-folding energy

No remarkable change in thecontinuum parameters

Ec, E

f & are more or less

consistent with the previouslymeasured values (White et al.1983; Orlandini et al. 2000).

The Continuum Parameters



Geometry of the GX 301-2/Wray 977 system (from Pravdo & Ghosh 2001)

Variation of the Column Densities



Considerable increase in column density near periastron; but also a substantialscatter in the values at intermediatephases

NH1

NH2

Large variation throughout the binary orbit(from 1022 to 1024 H atoms cm-2) seems to beone of the characteristics of the X-rayspectrum of GX 301--2

The large variations of NH1 & NH2 at allorbital phases indicate clumpiness of thestellar wind at different size scales.

Peak between 0.1 and 0.2 was expected

since the line of sight passes through thedensest parts of the wind

Thus it is clear that the observed variationin column density cannot be explained by aspherically symmetric CAK wind only

Comparison with the Models



2 models put forward to explain the NH

variation

Leahy (1991), from TENMA data, proposes a

spherically symmetric stellar wind + linear gasstream

The peak near periastron was fitted when thegas stream was introduced (see right, Leahy1991)

A gas stream can be due to the dynamicaleffect of the neutron star on the companionwind and the effect would be strongest atperiastron

Pravdo & Ghosh (2001), proposed the existence

of an equatorially enhanced circumstellar disk about Wray 977

The model describes two interactions of the NSwith the disk which gives rise to the two peaks inthe column density



Model calculations are very different from the observed variation with the RXTE-PCA data

There are probably strong inhomogeneties in the wind thatare causing large fluctuations in the column densities

Almost throughout the binary orbit the Covering Fractionremains high

Indicates the presence of clumpy inhomogeneous material

A spherical wind plus + stream fitted the variation in columndensity for 4U 1907+09 (Roberts et al. 2001)



Covering Fraction (CF)

defined as :

Norm2 / (Norm1+Norm2) ;

Norm1 and Norm2 arerespectively the normalizationsof the two power-laws

CF remains substantially high almost throughout theorbit which means thatthere is dense and clumpy material present throughout

Iron-line Flux & Equivalent Width



2 iron lines show large increases in flux near periastron and a possible smallincrease near 0.1 (at least for 1996)

The peak near periastron (phase 0.9) is not very evident in the 1996 data due tothe lack of enough observations

Fluxes for both lines in the intermediate phases show a more or less steady value

EQW has a clear correlation with NH2

, with considerable scatter



Makino et al. (1985) & Endo et al. (2002) have showedthat a correlation between N

H

& EQW exists in GX 301--2

Endo et al. (2002) discusses that the EQW increasesmonotonoically with the column density as long asthe fluorescing plasma is optically thin and fully surrounds the pulsar

The scatter seen in our correlation may be due to theclumpiness of the reprocessing material

For small values of the NH, the EQW is constant

This can happen if the fluorescing material does notsurround the pulsar completely

Conclusions for GX 301-2



Conclusions for GX 301-2

PCAM describes the X-ray spectrum well throughout thebinary orbit

Column Densities measured are very high with a large variation indicating a clumpy nature of the stellar wind

Correlation of the Fe-line EQW with NH2

suggests that

most of the Fe-line is produced by the local clumpymatter

Both the models do not clearly explain our results

OAO 1657-415 & Vela X-1 : orbital phase dependentt



spectroscopy

OAO 1657-415 Vela X-1 Orbital Period (d) : 10.44 8.96 Eclipse (d) : 1.7 2.0

eccentricity : 0.1 0.1

Distance (kpc) : 6.4 2.0

X-ray Luminosity (erg s-1) : 1037

Companion Star : B Supergiant B Supergiant

Pulse Period (s) : 38 283

Vela X-1 has fluctuations of luminosities to the order of 10%

These 2 comparable pulsars provide a good opportunity tostudy their stellar winds through spectral variability

Observations



OAO 1657-415 :

RXTE archival data from1997-10-31 to 1997-11-11 :

26 obs On-source time : 1--18 ks

Fairly good coverage of theorbit

Most purposes, pulsar was inthe eclipsed state

Also, we have analysed theMECS obs on 2001-08-14 for104 ks

Vela X-1 :

RXTE archival data from 2005-01-01 to 2005-01-09 : 38 obs On-source time : 2--18 ks

PCU0 had lost the propane layer,hence we did not use its data inany further analysis

For a very high luminosity pulsarlike Vela~X-1, this is not a majorissue

The dataset does not provide acomprehensive orbital coverage,i.e. it lacks data for the 0.4--0.7

orbital phase duration

Spectral Models



Quite similar to that used for GX 301-2

MECS spectrum of OAO 1657-415 best fitted withPCAM + 2 gaussian lines, at 6.5 keV (narrow) &@ 7.0 keV

Initial RXTE values of the parameters for theiteration : From MECS fit

Similarly for Vela X-1 : Same model + only one

gaussian line worked best

RXTE is a non--imaging instrument

& OAO 1657-415 under the Galactic ridgeemission

So, explicitly incorporate the ridge background asa separate spectral component : Raymond-Smithplasma + power-law with appropriatenormalizations (Valinia & Marshall 1998)

For the RXTE spectra : systematics to the tune of

1%

OAO 1657-415 Vela X-1



Continuum

OAO 1657-415 Vela X-1



C o v

e r i n

g

F r a c

t i o

n

C o v e r i n g

F r a c t

i o n

ColumnDensity

&

CoveringFraction

An Concl sions ??



Any Conclusions ??

The PCAM model appears to be somewhat generic for pulsars which have variable column density, at least

when observed with RXTE/SAX

For highly luminous pulsars, the CAK Model of stellar wind does not suffice to describe the column density variations

Moderately luminous pulsars probably validate the

spherical wind model

The continuum parameters do not show any significant variation over the orbit



Part II

U. Mukherjee & B. Paul 2005, A&A, 431, 667

Be/X-ray Binaries



Courtesy : Negueruela

Circumstellat Disc

Formation not clear

Probably due to rapid rotation

Be stars : major subclass of the B stars

Be means a non-supergiant star of spectral class Bwhich shows emission lines

B stars : HeI absorption lines in optical spectrum

High Temp : He loses 1 e

He II Absorption lines : O stars : hottest of all stars

Accretion Barrier ?



Be phenomenon is episodic transient

No outburst NS in quiescence

Low accretion rate : rm

> rc

Centrifugal Barrier exists

X-ray transients like 4U 0115+63 (Campana et al. 2001withBeppoSAX) &

GRO J1744-28 (Wijnands et al. 2002 with Chandra) in quiescencewith luminosities of 1033 erg s-1

Also been thought to be in the centrifugally inhibited regime

3A 0535+262 : Accretion and pulsations



during quiescence

The transient pulsar :

Magnetic Field : 1013 G

Pulse Period : 104 s

Orbital Period : 111 days

Optical Companion : O9.7IIIe Be star

Orbital Geometry

3 obs with the Beppo-SAX on

2000 Sep 4 (5:14 UT) : A

Oct 5 (00:42 UT) : B &

2001 Mar 5 (22:52 UT) : C



Distance of 2 kpc for this system (Steele et al. 1998)

The 2--10 keV X-ray luminosities measured from the threeobservations are in the range of (1.5-4.0) 1033 erg s-1,

We detected pulsations in C down to 2 1033 erg s-1

At this accretion rate the system is expected to be in the

centrifugally inhibited regime

Useful on-source times for A, B & C : 32 ks, 39 ks & 51 kseach for MECS and 19 ks, 30 ks & 43 ks for LECS

Data from only two of the MECS detectors were available

Timing Analysis



MECS light curves with res : 1.0 s

Average Count Rate for A & B 0.04 count s-1

For C 0.09 count s-1

Only C showed coherent pulsations103.41 s

Another period : 101.57 s whenperiod search over a wider rangebetween (101--105) s

Long term pulse period history rulesout 101.57 s

Probably arose due to the windowfunction of the light curve

Extracted 5 separate light-curves each



with data only with avg count rate lessthan 0.15 , 0.125 , 0.1 count s-1, 0.08 &0.06 count s-1 respectively

Detection of pulsations down to anaverage count rate of less than0.08 count s-1

But pulse profile for 0.06 count s-1 showspulsations (last panel)

Pulse profiles are background subtracted

Pulse Fraction = [(Maximum - Minimum)/Maximum] 50%

Thus it is fair to conclude that pulsationsexist for C down to at least a count rate of 0.06 count s-1

Spectral Analysis



Spectra fitted with PL or BB + exp. line of sight absp

BB does not fit A, B or C : red > 2.0

PL + exp. line of sight absp fits A, B well but not C

PL + BB + exp. line of sight absp fits C well

BB temp : 1.33 keV

Radius of BB 0.1 km



The BB component !!



Quite different from the soft excess usually observed in some binary X-raypulsars (Yokogawa et al. 2000, Paul et al. 2002)

In general modeled as a BB with temp 0.1 keV and the area of theemitting region 1015 -16 sq. cm.

Hickox, Narayan and Kallman (2004) have exploredthe physical origin of this so called “soft excess”

For sources with luminosity < 1036 erg s-1, it may be due to emission byphotoionized or collisionally heated diffuse gas, or thermal emission fromthe surface of the NS

This soft BB comp resembles that of 4U 0352+309 (Coburn et al. 2001)

& RX J0146.9+6121by Palombara & Mereghetti (2006)

This BB comp. has been attributed to polar cap emission

Conclusions



For all the 3 obs, A, B & C : accretion is expected to becentrifugally inhibited

Detection of pulsations in C at these low flux levels indicatesthat some matter may have leaked through onto the NS

surface

When the emission is non-pulsating, the X-ray spectrum isPL type

Origin of an additional BB component in the pulsedobservation is not very clear though it seems to be essentialin fitting the spectra



Part III

U. Mukherjee & B. Paul, 2006c

U. Mukherjee & B. Paul, 2006, J. Astrophys. Astr., 27, 37

B. Paul, H. Raichur & U. Mukherjee,2005, A&A, 442, L15

Centaurus X-3 : Any reason behind the high-low states ?



Detailed spectral analysis of the out of eclipse obs

with ASCA, BeppoSAX, Chandra, XMM-Newton & RXTE

in its different intensity states

High & low states have separate domains of

From RXTE : NH

has a ceiling in the high states

Low states have high EQW

Compare our results with that of LMC X-4 & Her X-1

Properties of Cen X-3



Ps 4.8 s

Porb

2.1 d

D 8 kpc (Krzeminski 1974).

O-type supergiant V779 Cen Companion (Krzeminski 1974).

V779 Cen radius of 12 R S

& M 17-19 MS(Hutchings et al.

1979).

QPOs at 40 mHz (Takeshima et al. 1991; Audley et al. 1996)strengthen the case for the presence of an accretion disk

Broadband (0.12-100) keV out-of-eclipse pulse-phase-



Broadband (0.12 100) keV out of eclipse pulse phaseavg spec generally described by :

an abs PL + a broad Fe emission line @ 6.7 keV +cut-off @ 14 keV (Santangelo et al. 1998, Burderi etal. 2000)

Soft excess < 1 keV : BB with kT 0.1 keV (Burderi et al. 2000)

Cyclotron resonant feature @ 28 keV

B (2.4-3.0) 1012 G (Santangelo et al. 1998)

1.5 – 3.0 keV



RXTE-ASM in different energy bands shows that Cen X-3 has aflux 40 times more in the bursting state as compared to the lowstate (Paul, Raichur & Mukherjee 2005)

Also, the low and high states last between a few to upto 110 dwithout having any periodicity

3.0 – 5.0 keV

5.0 – 12.0 keV

Not only Cen X 3 but several other X ray



Not only Cen X-3, but several other X-rayPulsars show long term X-ray intensity

modulations

Her X-1 & LMC X-4 show periodic X-ray

intensity variations, modulated at theirsuperorbital periods

Cen X-3 does not possess any such periodicbehaviour at all

There are complex X-ray spectral changesassociated with the intensity modulations at the



associated with the intensity modulations at thesuperorbital periods of Her X-1 & LMC X-4

(Naik & Paul 2003, 2004)

Broad band X-ray spectrum of Cen X-3 is more or

less similar to these pulsars

It is pertinent to explore the spectral changes of

this pulsar corresponding to its different intensity states

Observations



80 obs with RXTE from 1996 to 2000 covering the various

intensity states

2 obs. With BeppoSAX on 1999-06-24 (High) & 2000-06-06(Low)

XMM-Newton observed Cen X-3 on 2002-08-01 in the Low state

The obs with ASCA was on 1993-06-24 (Low)

With Chandra on 2000-12-30 (High)

All the observations were chosen in the out-of-eclipse state

Observation Log



RXTE Analysis



Continuum fitted with line of sight absp. + PL + a cut-off

In addition, a cyclotron absorption feature was also added

Fe-emission lines @ 6.4 keV, 6.67 keV and 6.95 keV were also included in the form of gaussian lines

The centre energies were fixed to the laboratory values and thewidths of the gaussian lines were fixed to that obtained with theXMM-Newton PN spectra respectively

All the other parameters were allowed to vary

A systematic uncertainty of 1.0% was added to the spectralchannels



For flux < 6.0 10-10 erg cm-2 s-1, thereseems to be a correlation between theCont. Flux &

The spectrum becomes harder at lowerflux levels

At a flux level > 2 10-9 erg cm-2 s-1,0.8 < < 1.3

Chandra (H)

SAX (H)

SAX (L)

ASCA (L)

XMM (L)

Naik & Paul (2003) have reported fromRXTE th t LMC X 4 i d ib d i th



RXTE that LMC X-4 is described in the

High state with 0.7 < <1.0 & Low state with < <

Thus for both these pulsars, a separate

range of values of may be designatedfor the low and high flux statesrespectively

Variation of NH

shows a few instances of high

to moderate values at low flux levels



For higher flux levels (> 2 10-9 erg cm-2 s-1),

there seems to be a ceiling of 4 1022 cm-2

For both there appears to be a lot of

scatter in their values as obtained with thePCA for intermediate flux levels

ASCA (L)SAX (L)

XMM (L) SAX (H)Chandra (H)

Variation of NH

does not quite pesent a very

conspicuous picture either



co sp cuous p ctu e e t e

Her X-1 shows an increase in NH

by almost 2

orders of magnitude in its low state comparedto its high state over its 35-d superorbitalperiod (Naik & Paul 2003)

On the other hand, LMC X-4 does not show anychange in N

Halong its superorbital period

(Lang et al. 1998)

Hence we find that Cen X-3 shows resultswhich are in between these two extremes

Fe-line fluxes against (7 - 25)keV PCA Flux does not showany systematic trend (see figureright)

SAX (H)



Large & considerable scatter atall flux levels for all the 3 lineintensities

EQW of the Fe-lines show a

reasonable trend (bottom panel,right)

For low fluxes, the values of EQW are quite high (a fewhundred eVs, especially for

Fe-K )

Then decreasing to acquire amore or less steady value of 100 eV for higher fluxes(> 6 10-9 erg cm-2 s-1)

ASCA

( )

SAX (L)

Chandra

XMM



LMC X-4 Her X-1

Naik & Paul, 2003

A similar variation of the EQW with continuum flux was seen in case of LMC X-4 & Her X-1

All the 3 sources show very high values of the EQW in the lowstates whereas during the high states, no significant change is seen

Unlike LMC X-4 & Her X-1, Cen X-3 does not show anyglobal dependence of the iron-line intensities with thecontin m fl



continuum flux

Low states : it may seem Fe-line intensity has somepositive correlation with the continuum intensity, but

High states : no such dependence

LMC X-4 & Her X-1 : a clear correlation between thefluxes at high states, indicating the production of theiron line near the continuum X-ray source

Picture for Cen X-3 is far more intricate & itsuggests that not all the Fe-line is produced closeto the pulsar

Inference



Our work attempts to compare the X-ray

observational results for Cen X-3 in both its highand low states with that of Her X-1 & LMC X-4

It is more or less established that Her X-1 &LMC X-4 have an warped accretion disk which isbelieved to be accounting for their flux behaviour

However spectral results obtained by us are notfully in corroboration with that measured for thethe other 2 pulsars

Hence, based on these results it seems difficult toput forward the same reasoning for the flux



behaviour in case of Cen X-3 as for the other two

Obscuration in the low states is present for

Cen X-3 (cases of high NH)

But that is not the only criterion to unequivocallyattribute the presence of high/low states to thewarped accretion disk

Ogilvie and Dubus (2001) : LMC X-4 & Her X-1 are 2possible candidates to show a super periodicity while



possible candidates to show a super-periodicity whileCen X-3 is unlikely to exhibit such

Numerical simulations of the evolution of theprecessing accretion disks in XBPs :

Iping & Petterson (1990) propose for Cen X-3

a precessing accretion disk could describe a partof its aperiodic X-ray flux variability

Day & Stevens (1993), based on EXOSAT data :

th t l t tt f di k



the central source accretes matter from a disk

which is fed by an X-ray excited wind emanatingfrom the companion

Scope for further theoretical and observationalwork in this direction remains open

Especially extensive orbital coverage at lower

energies with detectors of the genre of Chandra& XMM-Newton as well as a detailed broad bandcoverage

BeppoSAX

Model : PL + BB + a gaussian line

High



Model : PL + BB + a gaussian lineas additive comp. & an abs along

the line of sight as the multipl.comp.

Both the high and low state spectrawere well fitted

A broad Fe-emission line @6.65 keV

Presumably the blend of theHe-like triplet of Fe XXV

EQW (High) : 877 eV

EQW (Low) : 116 eV

Low

Model as used by

ASCA



Model as used byBurderi et al. (2000)

Fixed both the 6.44keV and 6.96 keV lines along with theirwidths as given in

Burderi et al. (2000) Red. : 1.47 for267 d.o.f.

2.0 keV Si XIV emission line

Low

Iaria et al. (2005) have

Chandra

Hi h



analyzed the Chandra-HEG

spectrum in the (6.0 - 7.6)keV & resolved the Fe XXV triplet

We used only the first ordergrating spectra

The zero-order source position determined by examining the HEG/MEGdispersion lines & the data read out streak of the zeroth-order image

We used the same model as of Iaria et al. (2005) but allowed the cont.

params to vary

We froze the centre energies of the triplet and their respective widthsthough we kept the params of the Fe K line thawed

High

XMM-Newton



The 100 s binned PN-light curve

showed a broad coverage of eclipse,egress and out of eclipse for about 68 ks.

Avg PN count rate was 15 counts s-1

For our spectral analysis, weselected out-of-eclipse dataonly (34 ks onwards)

Out of eclipse



Si XIV

S XVI Ar XVII+

Ar XVIII

Fe XXVI

Ca XX

Continuum Parameters

(*ASCA (L) & **SAX (H) : from Burderi et al. (2000))



4U 0114+65 : a luminosity dependentstudy



Spectral characteristics of the high and low statesof the pulsar 4U 0114+65 with ASCA

Examine the change in the parameters of thespectral model

Compare these spectral characteristics to theresults with other satellites

And with other X-ray pulsars ..........

About the Pulsar ..........



B1Ia supergiant optical companion

At a distance 7 kpc (Reig et al. 1996)

Crampton, Hutchings & Cowley (1985) reported anOrb Period 11.59 d (optical radial velocity measurements)

Their measurements were not able to distinguish between acircular or elliptical orbit

Corbet et al. (1999) obtained X-ray intensity modulations at aperiod 11.63 d from the long term light curve of RXTE-ASM

X-ray light curve : considerable variability,flaring activity for a few hours (Apparao et al. 1991,Finley et al. 1992)



Also short-term flickering for minutes (Koenigsbergeret al. 1983)

Presence of a 2.8 hr periodicity observed in the X-ray

light curves from EXOSAT and ROSAT data (Finley et al.1992)

Farrell et al. (2005) reported the detection of a

superorbital period of 30.7 d in this pulsar fromRXTE-ASM

Observations



Carried out by ASCA on 1997-02-10

Orbital phase0.19

Useful exposure time 25 ks

Total time span between start & end was55 ks (about 5.5% of the orbital period)

GIS : 0.7 – 10.0 keV



The intensity at the peak of the flare is 12-15 times than the

persistent low level emission

At times, the intensity almost drops down to zero (@ 22 & 32 ks)

High state : 0-18 ks & rest : low state

SIS : 0.5 – 9.0 keV

Bin : 100 s

High



Low

Most significant difference between thehigh & low state spectra is a change inthe Fe-line flux by a factor similar to thechange in the continuum



NH

corroborates well with that of RXTE & Ginga

But SAX measures a high NH



Masetti et al. (2005) reported for (1.5-100.0) keV whereas,we are able to go down to 0.5 keV

Fe emission line @ 6.4 keV was reported with SAX, RXTE,EXOSAT & Ginga

But the line was better detected during the low states

We report the detection of the Fe-line in both the stateswith appreciable EQWs (Table 2)

ASCA has the best spectral resolution in comparison,so the detectability of the Fe line can be supposed tobe most reliable

EQW of the Fe-line is found to increase appreciably in the low

Summarising...........



EQW of the Fe line is found to increase appreciably in the low

states of several other pulsars like Her X-1 (Naik & Paul 2003),LMC X-4 (Naik & Paul 2004) & SMC X-1 (Vrtilek et al. 2005)

High & low states of these pulsars are ascribed to the super-orbitalperiod precessing warped inner accretion disk

For 4U 0114+65, an abrupt decrease in X-ray flux over a fewthousand secs cannot be assigned to a super-orbital period

Rather, change in lum during our obs may be ascribed to a localchange in the density of the stellar wind

Rise & fall times of 1-2 ks as can be seen in the light curve

An assumed orbital velocity of a few hundred km/s indicatesclumpiness of the stellar wind at a length scale of about 1010-11 cm

Future Work ?!???!!



Binary Millisecond X-ray Pulsars

Techniques used for HMXBPs can be directly appliede.g. compare the characteristics of 3A 0535+262, to theobservations of BMXPs in quiescence

Furthermore, observations in the optical and IR will be

useful

Pulse Phase Resolved Spectroscopy of



p py f

LMXBs like 4U 1626-67 with Chandra

Detection of Centrifugally Inhibited Accretion for other Be/X-ray binaries with

BeppoSAX/XMM-Newton

uddipan kasi talk

Documents