Low and High Angular Momentum Accretion Flows in BHCs: Case Study of \xte
Roberto Soria
Kinwah Wu \address{Mullard Space Science Laboratory, UCL, Holmbury St.~Mary, Dorking, RH5~6NT, UK.}
Mike McCollough \address{SD50, NASA-MSFC, Huntsville, AL 35812, USA.}
The 1998 outburst of \xte started with a hard X-ray spike, rising in less than a day and declining after 3-4 days; at the same time, the soft X-ray flux was rising with a longer timescale ($\sim 10$ days). We suggest that the soft and the initial hard X-ray emission are produced by two different components of the accretion flow: a higher angular momentum flow, which forms the disk, and a lower angular momentum flow feeding the hot inner region. Thus, we argue that the onset of the outburst is determined by an increased mass transfer rate from the companion star, but the outburst morphology is also determined by the distribution of specific angular momentum in the accreting matter. In \xte, we attribute the initial, impulsive low angular momentum accretion to the breaking down of magnetic confinement by the magnetically active secondary star. We show that a hard X-ray spike is seen at the onset of an outburst in other BHCs.
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\newcommand{\etal}{{\it et al.} }
\newcommand{\asca}{{\it ASCA} }
\newcommand{\exosat}{{\it EXOSAT} }
\newcommand{\ginga}{{\it Ginga} }
\newcommand{\bbxrt}{{\it BBXRT} }
\newcommand{\rosat}{{\it ROSAT} }
\newcommand{\xmm}{{\it XMM-Newton} }
\newcommand{\chandra}{{\it Chandra} }
\newcommand{\conx}{{\it Constellation-X} }
\newcommand{\rxte}{{\it RXTE} }
\newcommand{\xte}{XTE J1550$-$564 }
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\markboth{R. Soria, K. Wu, D. C. Hannikainen, M. L. McCollough, 
R. W. Hunstead}{X-ray Emission from Accretion onto
Black Holes, JHU/LHEA Workshop, June 20--23, 2001}


\begin{document}

\title{Low and High Angular Momentum Accretion Flows in BHCs: Case Study of \xte }
 
\author{Roberto Soria, Kinwah Wu
\address{Mullard Space Science Laboratory,
UCL, Holmbury St.~Mary, Dorking, RH5~6NT, UK.},
Diana Hannikainen \address{Department of Physics and Astronomy,
  University of Southampton, Southampton, SO17~1BJ, UK.},
Mike McCollough \address{SD50, NASA-MSFC, Huntsville, AL 35812, USA.},
Richard Hunstead
\address{School of Physics, University of Sydney, 
  NSW 2006, Australia}}

\begin{abstract}
The 1998 outburst of \xte started with a hard X-ray spike, rising 
in less than a day and declining after 3-4 days; at the same time, 
the soft X-ray flux was rising with a longer timescale ($\sim 10$ days).
We suggest that the soft and the initial hard X-ray emission 
are produced by two different components of the accretion flow: 
a higher angular momentum flow, which forms the disk, and a lower 
angular momentum flow feeding the hot inner region. Thus, we argue that 
the onset of the outburst is determined 
by an increased mass transfer rate from the companion star, 
but the outburst morphology is also determined by the distribution 
of specific angular momentum in the accreting matter.
In \xte, we attribute the initial, impulsive low angular momentum accretion 
to the breaking down of magnetic confinement by the magnetically active 
secondary star. We show that a hard X-ray spike is seen at the onset 
of an outburst in other BHCs. 
\end{abstract}

\maketitle

%\cite{}

\section{Introduction}

%Many soft X-ray transients (SXTs) are binaries containing a low-mass
%  star transferring material to a black hole.
%SXT outbursts are usually attributed 
%  to accretion-disk or mass-transfer instabilities, and 
%  show a variety of patterns and properties. 
%The evolutionary status and magnetic properties of the companion star 
%may also be an important factor determining the outburst morphology.

The BH soft X-ray transient \xte was discovered on MJD~51063 
  (1998 Sep 7) 
  by the ASM on board 
  {\it RXTE} \cite{smi98} 
  and by BATSE  
  on board {\it CGRO} 
  \cite{mcc98}.   
The optical counterpart was identified shortly afterwards
  \cite{oro98}, 
  and a radio source was found at the optical position on MJD~51065 
\cite{cam98},
 with the Molonglo Observatory Synthesis Telescope (MOST), at 843MHz.
From the optical ellipsoidal modulations, 
an orbital period of $1.541\pm0.009$~d was inferred \cite{jai01}.

We used BATSE to monitor the hard X-ray emission (20--100 keV) 
from \xte, from the start of the 1998 Sep outburst until 
{\it CGRO}'s re-entry (2000 Jun). We plot the BATSE and ASM fluxes 
in Fig. 1, together with the ASM hardness ratios. Details of the technique 
used to obtain a daily lightcurve are given in \cite{kwu01}.
We also followed the radio source with MOST for the first 27 days after the 
initial detection. For a preliminary analysis 
of our radio data, see \cite{kwu01,han01,han02}.



\begin{figure*}[t] % fig.1
\vspace{-14pt}
\centerline{\psfig{file=f1a.ps,height=7.2in, width=4.0in, angle=270}}
\vspace{-30pt}
\caption{Top panel: BATSE (20--100 keV) and ASM (2--12 keV) fluxes from 1998 Sep 
to 2000 Jun. Bottom panel: ASM hardness ratios for the same period.
}\label{fig:smallfig}
\end{figure*} 

\begin{figure*}[t] % fig.2
\vspace{-30pt}
%\hspace{0.17in}
%{\epsfig{file=f1a_col.eps,width=2.7in,height=5.6in, angle=270}}\\
%\vspace{-20pt}
%\hspace{0.5in} 
{\epsfig{file=f1b_col.eps,width=6.8in, height=3.92in}}
\vspace{-37pt}
\caption{
%Top panel: Optical/IR/UV lightcurves for the 1998 outburst. 
%Open circles: data from \cite{san99}; 
%filled circles: data from \cite{jai99}. 
%The upper lightcurve 
%is in the standard $I$ band for the dataset \cite{jai99}, and 
%in Gunn $i$ for \cite{san99}. The other two lightcurves 
%are in standard $V$ and $U$ for both datasets. 
%A possible sharp flare at MJD 51076 
%was detected in the $V$ band by \cite{san99}, but was not 
%seen by \cite{jai99}. Bottom panel: 
Hard X-ray, soft X-ray and radio lightcurves of \xte; see \cite{kwu01} 
for details.
}\label{fig:largefig}
\end{figure*}


\section{Phenomenology of the outburst}  
Remarkable features of the 1998 outburst are: 
\begin{itemize}
\item {\it an initial impulsive rise in the hard X-rays} (Fig.\ 2).
The BATSE (20--100~keV) flux 
  reached about 0.3~ph~cm$^{-2}$s$^{-1}$ within one day, 
  peaking at about 0.4~ph~cm$^{-2}$s$^{-1}$ the next day. 
From {\it RXTE}/PCA$+$HEXTE data \cite{wil01}, 
the 20--200~keV flux peaked at 
$\approx 1.7 \times 10^{37}~{\rm erg}~{\rm s}^{-1}$, for a distance of 
2.5 kpc \cite{san99}.
Assuming an efficiency $\eta \sim 0.1$, 
  the total accreted mass required to account for the initial 
  hard X-ray burst is \ltae $10^{23}$~g;
\item {\it a subsequent giant flare} that occurred almost simultaneously 
      in both hard and soft X-rays 
      about twelve days after the initial rise of the hard X-rays. The radio 
flare occurred at MJD~51077.8$\pm0.1$, $\approx 1.8$~d 
  after the X-ray peak, 
  with the flux density reaching 380~mJy at 843~MHz.
The hard and soft X-ray flux then settled on a plateau, with spectral and timing properties  
similar to the ``very high state'' identified in other BHCs;
\item {\it an exponential decay in the
      soft X-ray flux} after about 40 days. 
      The same decay timescale for the soft flux is seen after 
      the 1998, 1999 and 2000 outbursts.
\end{itemize}

%Fig. 2 shows multiwavelength lightcurves for the 1998 outburst 
%(see \cite{kwu01}).





\section{Origin of the hard X-rays}

Hard X-ray emission  
is not thermal emission from an accretion disk (whose maximum 
temperature is $\sim 1$ keV), but is 
produced by inverse Compton scattering of softer photons 
off highly energetic ($E \sim 100$ keV) electrons near the BH.
Different models have been proposed to explain the source 
and location of the hot electrons.

\begin{itemize}
\item Accretion of matter 
with low angular momentum
can produce hard X-rays 
  \cite{igu99,bel00,cha95}.  
The physical justification is that free-falling 
electrons in a quasi-spherical inflow 
can acquire kinetic energies \gtae 100 keV as they approach 
the BH horizon, before reaching their 
circularisation radius. The kinetic energy of the infalling electrons 
powers the Comptonisation process, either directly (bulk-motion 
Comptonisation) or after it has been converted into thermal energy 
(thermal Comptonisation).

\item Alternatively, a hot, optically thin, quasi-spherical flow 
may already exist at small radii before the onset of the outburst 
(ADAF solutions;  \cite{nar95,esi97}). 
%The timescale for 
%accretion of the hot inner flow is shorter than its radiative 
%cooling timescale, hence energy is advected into the BH. 
The ADAF model may explain residual hard X-ray emission 
seen in some BHCs in quiescence.

\item Other physical mechanisms for the production and storing of  
the Comptonising electrons (eg, a magnetically heated corona 
on top of the disk) may be present 
at later stages of the outburst, and may explain 
the hard X-ray emission in the plateau phase (MJD 51080 -- 51100). 
But it is unlikely that they could operate at the onset 
of the outburst, before a disk was fully formed.

\end{itemize}





\section{Low angular momentum inflow}

We suggest that the initial hard X-ray spike 
is due to impulsive accretion of matter with low specific angular 
momentum, over the first few days. Meanwhile, high angular 
momentum matter would also be accreted through normal Roche lobe (RL) overflow.
{\it The outburst behaviour is determined by the mass transfer rate} 
  (which sets the energetics of the burst) 
{\it and the angular momentum distribution of the accretion flow}
  (which sets the initial spectral properties and the X-ray rise time).
%When the initial injection ends and the 
%  supply of low-angular momentum accreting matter is depleted,  
%  the hard X-ray flux declines.

In the ADAF scenario, the optically thin inner flow would be 
already present when the outburst starts, and 
would be responsible for the initial hard emission. The 
subsequent spectral softening  
  is attributed to the collapse of the ADAF 
  into an optically thick disk 
  when the mass accretion rate increases. 
%However, it is difficult to 
%explain why the initial hard spike lasts for a few days, 
%while the expected timescale for the collapse  
%is at least an order of magnitude shorter.
Only one parameter determines the morphology of the outburst 
in the ADAF model: the mass accretion rate. 

\section{Magnetic confinement}

Wind accretion is a typical example of low angular momentum 
accretion flows. In fact, the X-ray spectrum of \xte during the initial 
hard phase is very similar to the spectrum of Cyg X-1 in its 
low/hard state, when wind accretion is thought to dominate over 
RL overflow \cite{wil01}. However, 
the companion star in \xte is likely to be a K subgiant, 
therefore it cannot produce a wind. Is there an alternative source 
for the subkeplerian accretion flow?

We notice that fast-rotating K subgiant stars in binary systems 
tend to have a strong magnetic field. 
The main physical reason is   
that, while the tidally-deformed stellar envelope 
  can be locked in synchronous rotation with the orbit, 
  the degenerate stellar core 
  may not attain perfect rotational synchronism.
The differential rotation leads to dynamo action and magnetic activity.

For example, RS CVn systems 
are magnetic close binaries containing a G/K subgiant 
with a multi-kilogauss global dipolar field \cite{don92,vog99}. 
Given that the orbital period in \xte is even shorter, 
and tidal deformation of the companion is even stronger than in 
typical RS CVn systems, 
we suggest that the K subgiant companion in \xte is also strongly magnetic. 
%A moderate surface field strength of 1~kG, 
%  corresponding to a magnetic stress of
%  $4\times 10^4$~erg~cm$^{-3}$, 
%  would be strong enough to influence 
%  the mass- and energy-transport process in the stellar atmosphere.

Extended regions of cool, optically-thin, magnetically confined material 
  around a K star are a common feature of RS CVn systems  \cite{hal94}. 
We suggest that the same magnetic confinement occurs in \xte.
%The typical scale-height of these prominences above the photosphere 
%  is about twice the stellar radius, 
%i.e., a few times $10^{11}$ cm, 
%  implying a volume of $\sim 10^{34}$--$10^{35}$ cm$^3$ 
%  for the confining region.
In fact, the mass of the confined gas 
   can be even higher here, 
   because the magnetic companion star is close to 
   filling its RL. 
%The condition for confinement 
%is that the magnetic tension/stress must be greater than the gas pressure. 
We show in \cite{kwu01} that a mass $> 10^{23}$~g 
can be stored in this ``magnetic bag'', enough to account for 
the initial hard X-ray burst.
By analogy with RS CVn systems, 
  the confined gas in \xte 
  may be located well beyond the L$_1$ point, 
  deep into the RL of the primary, 
  but corotating with the binary (therefore with lower 
angular momentum than the material accreted via RL overflow).

%For magnetic confinement near L$_1$, 
%  the magnetic tension/stress must be greater than the gas pressure, 
%  implying a critical density for the confined matter:   
%\begin{equation}  
%  \rho_{\rm c} \approx  {\mu {m_{_{\rm H}}B^2}\over {8\pi kT}}    
%        \sim 2 \times 10^{-8} {\rm g~cm}^{-3} 
%        \biggl({B \over {1 ~{\rm kG}}} \biggr)^2  
%        \biggl({T \over {10^4~{\rm K}}}\biggr)^{-1} \ . 
%\end{equation}

%For the sizes of the confining regions in RS CVn systems, 
%  this corresponds to a mass upper limit of $\sim 10^{26}$~g.   
%By analogy with RS CVn systems, 
%  the confined gas in \xte1550 
%  may be located well beyond the L$_1$ point, 
%  deep into the Roche lobe of the primary, 
%  but corotating with the binary.
%In reality, the confined mass is likely to be limited 
%   by the mass loss rate of the secondary: 
%   a rate of $\sim 10^{17}$ g s$^{-1}$ 
%   would supply a mass of $\sim 10^{23}$~g 
%   over a timescale of a few days 
%   on which the prominence forms, 
%   enough to account for the observed hard X-ray luminosity.

The start of the outburst 
is determined by an increased mass loss from the secondary star. 
Part of this matter goes through L$_1$, is accreted with a higher 
specific angular momentum 
and is responsible for the formation of the disk. Another component is 
temporarily channelled 
into the magnetic bags above the companion star, until the confinement 
breaks down; it then 
falls towards the accreting primary with lower angular momentum, 
feeding the hot, hard X-ray emitting region.



\section{A subclass of BHCs with a hard spike}

\xte is not the only system showing a hard X-ray spike 
  at the onset of an outburst.   
XTE J1859$+$226 \cite{woo99,mcc99}
  and XTE J2012$+$381 \cite{vas00} are two other good examples.  
In these three systems, 
  the hard X-ray flux reached a maximum within the first day after detection, 
  and then declined after about 4 days.
The soft X-ray flux increased more steadily and at a slower pace, 
  for $\sim 10$ days. 
%  until a plateau was reached.     
The similarity in their X-ray spectral properties 
%  of \xte, XTE J1859$+$226 and XTE J2012$+$381 
  can be seen from their hardness-ratio plots (Fig. 3).
%  data from the public {\it RXTE}/ASM archive).     
There is probably a subclass of BH transients,  
  sharing a similar physical mechanism 
  that gives rise to a strong hard X-ray spike  
  at the onset of an outburst.  
A more detailed analysis of the hard X-ray properties 
  of this type of systems will be presented elsewhere.


\section{Conclusions}

We discuss evidence of a distribution of angular momenta in the accretion 
flow onto the BH transient \xte. We argue that the initial hard X-ray burst 
is due to a subkeplerian component, while matter accreted through the Roche  
lobe is responsible for the formation of a disk. 
We propose that the companion star 
is magnetically active, and its magnetic field creates a ``bag'' 
capable of confining $\sim 10^{23}$ g of gas inside 
the Roche lobe of the primary, corotating with the binary. 
When the confinement breaks down, this low angular 
momentum gas reaches the inner region near the BH horizon 
on a free-fall timescale.
We suggest that the onset of the outburst is determined 
by an increased mass transfer rate from the companion star, 
but the outburst morphology is determined by the distribution 
of specific angular momentum in the accreting matter. 




\begin{figure}[t] % fig.1
\vspace{10pt}
\centerline{\epsfig{file=f4_col.eps,width=2.1in,height=2.9in, angle=270}}
\vspace{-20pt}
\caption{
Evolution of 
the {\it RXTE}/ASM hardness ratios (5--12 / 3--5 keV band) 
during recent outbursts of the BHCs \xte (open circles), 
XTE J1859$+$226 (filled circles) and XTE J2012$+$381 (asterisks). 
See  \cite{kwu01} for details.
%All three systems show an initial hard spike.
%For \xte1550, time $=$ MJD $-$ 50063; 
%for XTE J1859$+$226, time $=$ MJD $-$ 51459, and for XTE J2012$+$381, time $=$ 
%MJD $-$ 50956. 
}\label{fig:smallfig}
\end{figure} 



%\section{References}


%do not change this
\small
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%do not change this
\normalsize

%\section*{ACKNOWLEDGEMENTS}
 
%Put all acknowledgements and grant numbers, etc. here, after the references.
%Please leave the section heading in capitals.                                            
\end{document}




\section{Figures}

Follow the example of Fig.~\ref{fig:smallfig} for figures that are
to appear in one column. Follow the example of Fig.~\ref{fig:largefig}
for figures that will span across two columns. You can adjust
the dimensions of the figure in either case within the \verb-\psfig{}-
command. Examine {\tt jhuwkshp.tex} to see how. Note that
the small figure in this article is a bad example because the
text in the figure is illegible. Please use your judgment.
The main difference between figures that span one column and 
those that span the
whole page is that the former use \verb-\begin{figure}- and
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In the text, figures should be referred to as 
Fig. or Figs. as opposed to Figure or Figures. The examples
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You may use color, but remember that if someone prints out the 
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\section{Equations}

Simple equations may be inserted in the text using mathmode, or
as a separate block using \verb-\begin{equation}-, \verb-\end{equation}- or
\verb-\begin{eqnarray}-, \verb-\end{eqnarray}- (the latter for multi-line
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and Equ.~\ref{eqn:twoleqn} uses \verb-\begin{eqnarray}-.

\begin{equation}
v = \sqrt{\left(\frac{3kT}{m}\right)}
\label{eqn:oneleqn}
\end{equation}


\begin{eqnarray}
\frac{L}{L_{\rm EDD}} & = & 1.236 \times 10^{38} \ M_{\odot} \nonumber \\
& = & 1.236 \times 10^{46} \ M_{8}
\label{eqn:twoleqn}
\end{eqnarray}

Equations are numbered automatically and sequentially, even if
you mix and match {\tt equation} and {\tt eqnarray}.
Note that when using {\tt eqnarray} you must have
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Refer to equations in the text by Equ. or Equs. and not Equation or Equations.
The examples
in {\tt jhuwkshp.tex} also show you how to label equations and refer to them
without explicitly referring to an equation number.  

%If you have a huge, complicated equation and can figure out how to
%make it span two columns, feel free to do so. 

%Note that the way errors on parameters are shown in this Table
%is only an example; you may use $xxx^{+xxx}_{-xxx}$ or $ xxx \pm xxx$ 
\begin{table*}[t]   
\caption{The Narrow Fe-K Line Component in Some AGN } 
\label{tab:largetab} 
\begin{tabular*}{\textwidth}{@{}l@{\extracolsep{\fill}}rrr}
\hline
& \multicolumn{3}{c}{Narrow-Line Parameters } \\
\cline{2-4}
Source & $E_{\rm Fe-K}$ (keV) & EW (eV) & FWHM (km/s) \\  
\hline
NGC 5548 & $6.402 \ (+0.027,-0.025)$ & $133\ (+62,-54)$ & $<8040$ \\
NGC 3783 & $6.347\ (+0.050,-0.021)$ & $86 \ (+35,-28)$ & $<6665$ \\
NGC 4151 & $6.386 \ (+0.014,-0.016)$ & $330 \ (+63,-67)$ & $<8169$ \\
Mkn 509 & $6.389 \ (+0.082,-0.030)$ & $61 \ (+35,-21)$ & $<6290$ \\
NGC 4051 & $6.444 \ (+0.051,-0.052)$ & $198 \ (+99,-82)$ & $<8851$ \\
NGC 3327 & 6.4 (fixed) & $<85$ & - \\
3C 273 & 6.4 (fixed) & $<25$ & - \\   
\hline 
\multicolumn{4}{@{}p{120mm}}{Errors and upper limits are for 90\% confidence, one interesting parameter.}

\end{tabular*}
\end{table*}

\section{Submission Process}

{\it Note: All filenames must use lower case}.

\begin{itemize}

\item{Name the {\tt .tex} file of your article
with your last name, followed
by an  underscore, followed your first two initials (one if
you don't have two), plus an extension {\tt .tex}. For example, 
T. Yaqoob would use {\tt yaqoob\_t.tex}.}

\item{Name your postscript figure files  
with your last name, followed by 
an  underscore, followed by your first two initials and
then another underscore, followed by {\tt fig1.ps}, {\tt fig2.ps}, etc.
For example, T. Yaqoob would use
{\tt yaqoob\_t\_fig1.ps}, {\tt yaqoob\_t\_fig2.ps}, etc.}

\item{Make a postscript file of the complete article, including the figures,
(with no errors remaining in the {\tt .tex} file),
and give it the same name as the 
{\tt .tex} file, except with an extension {\tt .ps}}

\item{Make a {\tt .tar} file of the {\tt .tex} file, the {\tt *fig*.ps} files,
and the {\tt .ps} file for the complete article. 
The name of the {\tt .tar} file should have a
root name which is the same as that of the {\tt .tex} file, plus an  underscore, 
plus {\tt wkshpfiles.tar}. For example {\tt yaqoob\_t\_wkshpfiles.tar}.
To make the {\tt .tar} file put all the files (the {\tt .tex} and all {\tt .ps} files)
in a separate directory then do 

{\tt tar -cvf  yaqoob\_t\_wkshpfiles.tar *}

Then {\tt gzip} the {\tt .tar} file.  }


\item{Mail the gzipped {\tt tar} file, as an attachment, to
{\tt yaqoob@starquake.pha.jhu.edu} by the deadline of September 14, 2001.} 

\end{itemize}

\section{References}

Follow ApJ style for references. Use the examples in
this article by examining {\tt jhuwkshp.tex}. References
should be prepared according to the sequential numeric system.
Citations in the text can be done
as in the following examples
(see {\tt jhuwkshp.tex}).
A single paper can be cited like this \cite{Tana1994} 
using only a command like \verb-\cite{Tana1994}-.
Several papers which form a complete sequence in the
reference list can be cited like this 
\cite{Tana1994,Merr1994,Nand1997} using a
command like \verb-\cite{Tana1994,Merr1994,Nand1997}-.
Several papers which don't form a sequence can be cited
like this \cite{Tana1994,Nand1997} using a command like
\verb-\cite{Tana1994,Nand1997}-. 
In the reference section,
all author names after the third
may be replaced by \etal as in \cite{Nand1997}.



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