ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

A secondary goal is to determine the distribution of impact velocities for those cases that end in planetary collisions.

The parameter space [or this study is large.

The parameter space for this study is large.

For (he sake of definiteness. the star has mass M,=LOAM. and the Jovian planet has starting semi-major axis α=0.05 AU (Poy724 dav).

For the sake of definiteness, the star has mass $M_\ast={1.0}M_\odot$ and the Jovian planet has starting semi-major axis $a=0.05$ AU $P_{orb}\approx{4}$ day).

The eccentricity of (he giant planet varies over (he range 0xe0.3 (these planets are expected to become tidally circularized. but only on much longer timescales).

The eccentricity of the giant planet varies over the range ${0}\le{e}\le{0.3}$ (these planets are expected to become tidally circularized, but only on much longer timescales).

The rocky planet starts just outside the 5:1 mean motion resonance (à£0.15 AU). with small eccenlricily e=0.001. and fixed mass mp=10M. :in this problem. the rocky planet acts like a test particle. so ils nass cannot greatly affect the dviamies.

The rocky planet starts just outside the 5:1 mean motion resonance $a\approx0.15$ AU), with small eccentricity $e=0.001$, and fixed mass $m_P=10\mearth$; in this problem, the rocky planet acts like a test particle, so its mass cannot greatly affect the dynamics.

The migration rate of the rocky planet varies with location. according to equation (7)): inside the disk edge (azz 0.05AU). migration ceases.

The migration rate of the rocky planet varies with location, according to equation \ref{awork}) ); inside the disk edge $a\lta{0.05}$ AU), migration ceases.

With (hese specifications. we consider the ellects of varying the mass anc

With these specifications, we consider the effects of varying the mass and

With (hese specifications. we consider the ellects of varying the mass ancl

With these specifications, we consider the effects of varying the mass and

channel).

).

Some of the "primordial binary systems may evolve to WD binary systems after the first mass transfer phase.

Some of the 'primordial' binary systems may evolve to WD binary systems after the first mass transfer phase.

In this case. the secondary continues to evolve and may fill its Roche lobe when it has a helium core.

In this case, the secondary continues to evolve and may fill its Roche lobe when it has a helium core.

If the mass transfer is stable. a wide EHB binary with a WD companion ts resulted given that the helium core is ignited later channel).

If the mass transfer is stable, a wide EHB binary with a WD companion is resulted given that the helium core is ignited later ).

The mass transfer is very likely to be dynamically unstable. and this results in the formation of a CE. and the CE ejection produces a close EHB binary with a WD companion if the helium core is 1gnited later channel).

The mass transfer is very likely to be dynamically unstable, and this results in the formation of a CE, and the CE ejection produces a close EHB binary with a WD companion if the helium core is ignited later ).

For a WD binary system. if the WD is a helium type and the secondary is on the first giant branch (1.e.. containing a helium core) when it fills its Roche lobe. the resultant CE contains helium WD pairs.

For a WD binary system, if the WD is a helium type and the secondary is on the first giant branch (i.e., containing a helium core) when it fills its Roche lobe, the resultant CE contains helium WD pairs.

The ejection of the CE leads to the formation of close helium WD pairs. and the close helium WD pairs may merge due to angular momentum loss of gravitational wave radiation to form single EHB stars channe!).

The ejection of the CE leads to the formation of close helium WD pairs, and the close helium WD pairs may merge due to angular momentum loss of gravitational wave radiation to form single EHB stars ).

Binaries with primary’s initial mass. Mj;~0.95—7M. and initial orbital period P;~1—1000d. may produce EHB stars via the channels described above.

Binaries with primary's initial mass, $M_{\rm 1i}\sim 0.95\,-\,7M_\odot$ and initial orbital period $P_{\rm i}\sim 1\, -\, 1000\, {\rm d}$, may produce EHB stars via the channels described above.

For the Ist stable RLOF channel. the primary fills its Roche lobe on the Hertzsprung gap or the first giant branch. and the mass transfer strips the primary of its envelope and leaves a naked helium core.

For the 1st stable RLOF channel, the primary fills its Roche lobe on the Hertzsprung gap or the first giant branch, and the mass transfer strips the primary of its envelope and leaves a naked helium core.

This leads to a wide EHB+MS binary. if helium is 1gnited.

This leads to a wide EHB+MS binary, if helium is ignited.

For a stellar population of age t<1Gyr (Mj;2 2M). a binary with P,~1-100d can evolve to a wide EHB+MS binary via stable RLOF on the Hertzsprung gap or the first giant branch. and a larger EHB mass corresponds to a small ¢ (a large Mj).

For a stellar population of age $t<1\,{\rm Gyr}$ $M_{\rm 1i}\ga 2M_\odot$ ), a binary with $P_{\rm i}\sim 1\,-\,100\,{\rm d}$ can evolve to a wide EHB+MS binary via stable RLOF on the Hertzsprung gap or the first giant branch, and a larger EHB mass corresponds to a small $t$ (a large $M_{\rm 1i}$ ).

For tf>IGyr (Mj;€ 2M.s). only RLOF near the tip of the first giant branch (1.e.. with a narrow range of Pi. typically Alog(P;/d)~ 0.5) can lead to EHB binaries. as the naked helium core is not ignited if the RLOF ts not close to the tip of the first giant branch (see Table 4 of Han et al. 2002)).

For $t>1\,{\rm Gyr}$ $M_{\rm 1i}\la 2M_\odot$ ), only RLOF near the tip of the first giant branch (i.e., with a narrow range of $P_{\rm i}$, typically $\Delta \log (P_{\rm i}/{\rm d})\sim 0.5$ ) can lead to EHB binaries, as the naked helium core is not ignited if the RLOF is not close to the tip of the first giant branch (see Table 4 of Han et al. \cite{han02}) ).

The corresponding orbital period P; is from ~10d to ~1000d for a stellar population with age ¢ of GGyr to GGyr.

The corresponding orbital period $P_{\rm i}$ is from $\sim 10\,{\rm d}$ to $\sim 1000\,{\rm d}$ for a stellar population with age $t$ of Gyr to Gyr.

The orbital period and the EHB mass of the resultant EHB binary would be larger for a large t.

The orbital period and the EHB mass of the resultant EHB binary would be larger for a large $t$.

For the Ist CE channel to form an EHB binary. the primary of a binary system needs to fill its Roche lobe while it has a helium core.

For the 1st CE channel to form an EHB binary, the primary of a binary system needs to fill its Roche lobe while it has a helium core.

However. the resultant CE cannot be ejected to form a close EHB binary due to a tight envelope ift<|Gyr (Mj;2 2M...

However, the resultant CE cannot be ejected to form a close EHB binary due to a tight envelope if $t<1\,{\rm Gyr}$ $M_{\rm 1i}\ga 2M_\odot$ ).

Fort>|Gyr. the primary of a binary system needs to fill its Roche lobe very close to the tip of the first giant branch (re.. with a typical range of initial orbital period Alog(P;/d)~ 0.1) (see Fig.

For $t>1\,{\rm Gyr}$, the primary of a binary system needs to fill its Roche lobe very close to the tip of the first giant branch (i.e., with a typical range of initial orbital period $\Delta \log (P_{\rm i}/{\rm d})\sim 0.1$ ) (see Fig.

| of Han et al. 2002).

1 of Han et al. \cite{han02}) ),

otherwise the helium core cannot be ignited after the CE ejection.

A large ¢ (a small Mj;) corresponds to a larger Pj (P;~100—1000d fort~1—I5 Gyr) and a larger orbital period P of the resultant EHB binary.

A large $t$ (a small $M_{\rm 1i}$ ) corresponds to a larger $P_{\rm i}$ $P_{\rm i} \sim 100\,-\, 1000\,{\rm d}$ for $t\sim 1\,-\,15\,{\rm Gyr}$ ) and a larger orbital period $P$ of the resultant EHB binary.

This channel produces a close EHB binary with a typical EHB mass of ~0.46M...

This channel produces a close EHB binary with a typical EHB mass of $\sim 0.46M_\odot$.

No EHB star can form from the 2nd stable RLOF channel in the HPMM model. as the mass ratio of the secondary (with an appropriate helium core) to its WD primary ts too large and the RLOF is not stable.

No EHB star can form from the 2nd stable RLOF channel in the HPMM model, as the mass ratio of the secondary (with an appropriate helium core) to its WD primary is too large and the RLOF is not stable.

For the second CE channel to form an EHB star. the primary of a binary system first needs to experience a stable RLOF (with P;~10— 10004) to form a WD binary. and the secondary of the WD binary needs to fill its Roche lobe when its helium core mass 1s in an appropriate range (see Fig.

For the second CE channel to form an EHB star, the primary of a binary system first needs to experience a stable RLOF (with $P_{\rm i}\sim 10\,-\,1000\,{\rm d}$ ) to form a WD binary, and the secondary of the WD binary needs to fill its Roche lobe when its helium core mass is in an appropriate range (see Fig.

| of Han et al. 2002)).

1 of Han et al. \cite{han02}) ).

The ejection of the resultant CE leads to a close WD+EHB binary.

As the WD can spiral in deeper in the envelope during the CE ejection. the WD+EHB can have a much shorter orbital period (as short as ~0.02d for a small 4) than the MS+EHB binary from the first CE ejection channel.

As the WD can spiral in deeper in the envelope during the CE ejection, the WD+EHB can have a much shorter orbital period (as short as $\sim 0.02\,{\rm d}$ for a small $t$ ) than the MS+EHB binary from the first CE ejection channel.

For a large f. the WD spirals in the envelope of a less massive secondary (the envelope is more loosely bound). and the resultant WD+EHB binary is wider.

For a large $t$, the WD spirals in the envelope of a less massive secondary (the envelope is more loosely bound), and the resultant WD+EHB binary is wider.

For the merger channel. a binary system first needs to experience a stable RLOF (with P;~4—250d and Mj;~0.95--2M.) to produce a helium WD binary. and the binary experiences a CE evolution to form a helium WD pair.

For the merger channel, a binary system first needs to experience a stable RLOF (with $P_{\rm i}\sim 4\,-\,250\,{\rm d}$ and $M_{\rm 1i}\sim 0.95\,-\,2\,M_\odot$ ) to produce a helium WD binary, and the binary experiences a CE evolution to form a helium WD pair.

The pair may coalesce to form a single EHB star due to angular momentum loss of gravitational wave radiation.

The EHB star produced from this channel has a wider mass range (0.4—0.54.) than from other channels.

The EHB star produced from this channel has a wider mass range $0.4\,-\,0.8M_\odot$ ) than from other channels.

The mass range is smaller for a small ¢ (0.56—0.64M.« for t=2Gyr) as only very close helium WD pairs have time enough to merge. but the range becomes larger with a large t(~ 04-O.8M.» for t=15Gyr).

The mass range is smaller for a small $t$ $0.56\,-\,0.64M_\odot$ for $t=2\,{\rm Gyr}$ ) as only very close helium WD pairs have time enough to merge, but the range becomes larger with a large $t$ $\sim 0.4\,-\,0.8M_\odot$ for $t=15\,{\rm Gyr}$ ).

In the process of CE ejection. the orbital energy released by the orbital decay of the embedded binary is used to overcome the binding energy of the CE.

In the process of CE ejection, the orbital energy released by the orbital decay of the embedded binary is used to overcome the binding energy of the CE.

As ts usual. I defined two parameters: the CE ejection efficiency ece. tthe fraction of the released orbital energy used to overcome the binding energy; and αμ. which defines the fraction of the thermal energy contributing to the binding energy of the CE.

As is usual, I defined two parameters: the CE ejection efficiency $\alpha_{\rm CE}$, the fraction of the released orbital energy used to overcome the binding energy; and $\alpha_{\rm th}$, which defines the fraction of the thermal energy contributing to the binding energy of the CE.

As in the HPMM model. I adopted gui=1.5 (for stable RLOF on the first giant branch). and ας.=ayy,0.75 as the best choices. and varied them to see their effects.

As in the HPMM model, I adopted $q_{\rm crit}=1.5$ (for stable RLOF on the first giant branch), and $\alpha_{\rm CE}=\alpha_{\rm th}=0.75$ as the best choices, and varied them to see their effects.

To obtain the distributions of properties of EHB stars at different ages. | have performed detailed Monte Carlo simulations with the binary population synthesis code developed for the HPMM model.

To obtain the distributions of properties of EHB stars at different ages, I have performed detailed Monte Carlo simulations with the binary population synthesis code developed for the HPMM model.

In the simulation. I followed the evolution of 10 million sample binaries according to grids of stellar models of solar metallicity and the evolution channels leading to EHB stars.

In the simulation, I followed the evolution of 10 million sample binaries according to grids of stellar models of solar metallicity and the evolution channels leading to EHB stars.

I adopted the following input for the simulations (see Han. Podsiadlowski Eggleton 1995)): Similar to the HPMM model. I have carried out 5 sets of Monte Carlo simulations altogether for Population | by varying the model parameters over a reasonable range (see Table 1)).

I adopted the following input for the simulations (see Han, Podsiadlowski Eggleton \cite{han95}) ): Similar to the HPMM model, I have carried out 5 sets of Monte Carlo simulations altogether for Population I by varying the model parameters over a reasonable range (see Table \ref{table1}) ).

central source in these obscured AGNs.

In section 4, we use mid-IR narrow-line emission and AGN-produced mid-IR continuum emission to determine the intrinsic luminosity of the obscured AGNs.

We combine each of these AGN luminosity indicators in order to reliably identify which sources are Compton-thick AGNs.

We use these results to further constrain the ubiquity of Compton-thick AGNs at zc0.1.

We use these results to further constrain the ubiquity of Compton-thick AGNs at $z \sim 0.1$.

Throughout, we adopt a standard ACDM cosmology of Ho=7lkms~'Mpc™', Qu=0.30, and Qa=0.70.

Throughout, we adopt a standard $\Lambda$ CDM cosmology of $H_0 = 71 \kmpspMpc$, $\Omega_M = 0.30$, and $\Omega_\Lambda = 0.70$.

We select our candidate Compton-thick AGN sample on the basis of their optical and X-ray properties.

Sources that are identified to be AGNs using traditional optical emission line diagnostics (e.g., ?)) but are undetected to faint limits in wide-fieldXMM-Newton observations (ie., fx/ffouy« 1) are strong candidates for containing heavily obscured AGNs (e.g., ??7)).! Here we provide the details behind the construction of our sample of X-ray undetected optically identified AGNs (i.e., candidate Compton-thick AGNs).

Sources that are identified to be AGNs using traditional optical emission line diagnostics (e.g., \citealt{bpt}) ) but are undetected to faint limits in wide-field observations (i.e., $f_X/f_{\rm [OIII]} < 1$ ) are strong candidates for containing heavily obscured AGNs (e.g., \citealt{bassani99, panessa02, akylas09}) Here we provide the details behind the construction of our sample of X-ray undetected optically identified AGNs (i.e., candidate Compton-thick AGNs).

We construct a parent sample of all optical spectroscopically identified galaxies in the zz100 deg? overlap region between the seventh data release of the SDSS (?;; hereafter SDSS-DR7) and the second source catalogue of theXMM-Newton Serendipitous survey (?;; hereafter 2XMMi).

We construct a parent sample of all optical spectroscopically identified galaxies in the $\approx 100$ $^2$ overlap region between the seventh data release of the SDSS \citealt{sdss_dr7}; hereafter SDSS-DR7) and the second source catalogue of the Serendipitous survey \citealt{watson09}; hereafter 2XMMi).

We define the redshift range for our shallow wide-area sample based on the combined available cosmological volume in the deep 2Ms “pencil-beam” CDF-North (zz448 arcmin?; ?)) and CDF-South (436 arcmin?; ??)) surveys.

We define the redshift range for our shallow wide-area sample based on the combined available cosmological volume in the deep 2Ms “pencil-beam” CDF-North $\approx 448$ $^2$; \citealt{dma03a}) ) and CDF-South $\approx 436$ $^2$; \citealt{giacconi02,Luo08}) ) surveys.

At z~0.5-2.5, where the CDFs are complete towards X-ray luminous AGNs, the encompassed comoving volume is V~4.57x10° Mpc?, which is equivalent to the comoving volume in the redshift range of z~0.03-0.2 in our SDSS-2XMMi selected sample.

At $z \sim 0.5$ –2.5, where the CDFs are complete towards X-ray luminous AGNs, the encompassed comoving volume is $V \sim 4.57 \times 10^6$ $^3$, which is equivalent to the comoving volume in the redshift range of $z \sim 0.03$ –0.2 in our SDSS-2XMMi selected sample.

The SDSS-DR'7 is currently the largest publicly available optical spectroscopic catalogue (9830 deg?) containing 929,555 spectroscopic source Previous studies have used past data releases of the survey to show that through careful spectral analyses, general galaxy and AGN properties can be derived from these large datasets (e.g., ??7)).

The SDSS-DR7 is currently the largest publicly available optical spectroscopic catalogue $\approx 9830$ $^{2}$ ) containing 929,555 spectroscopic source Previous studies have used past data releases of the survey to show that through careful spectral analyses, general galaxy and AGN properties can be derived from these large datasets (e.g., \citealt{kauff03b, heckman04, gre_ho07}) ).

We select all galaxies with well detected narrow A5007, Ha, [Nr1|A6585 emission-lines (S/N >5).? All galaxies with detected broad Balmer emission lines (here defined as a full-width half maximum >700kms!) are removed as these sources are unlikely to be intrinsically obscured by a gas/dust-rich geometrically thick torus.

We select all galaxies with well detected narrow $\lambda 5007$, $\alpha$, $\lambda 6585$ emission-lines (S/N $> 5$ All galaxies with detected broad Balmer emission lines (here defined as a full-width half maximum $> 700\kmps$ ) are removed as these sources are unlikely to be intrinsically obscured by a gas/dust-rich geometrically thick torus.

AGNs which are heavily obscured are often found to be hosted in dust-rich galaxies, and thus are likely to be strongly reddened (ie. Ha-H@ ratios >>3.1; eg., ?)).

AGNs which are heavily obscured are often found to be hosted in dust-rich galaxies, and thus are likely to be strongly reddened (i.e., $\alpha$ $\beta$ ratios $\gg 3.1$; e.g., \citealt{GA09}) ).

Hence, whilst useful in unambiguously discriminating between the properties of galaxies (e.g., Kauffmann et al.

2003; 7?)), we purposely do not limit our selection to only galaxies with well-detected Hj8 emission.

2003; \citealt{wild10}) ), we purposely do not limit our selection to only galaxies with well-detected $\beta$ emission.

Sources are separated by classification based on their optical emission-line ratios in a traditional diagnostic diagram (hereafter, BPT diagram; e.g., ?)).

Sources are separated by classification based on their optical emission-line ratios in a traditional diagnostic diagram (hereafter, BPT diagram; e.g., \citealt{bpt}) ).

We conservatively identify the narrow-line AGNs in the SDSS-DR7 as those which lie above the theoretical starburst limit of ?..

We conservatively identify the narrow-line AGNs in the SDSS-DR7 as those which lie above the theoretical starburst limit of \cite{kewley_bpt}.

See Fig. 1..

See Fig. \ref{fig:bpt}.

The 2XMMi catalogue identifies all X-ray sources detected in the 3491 observations made during the first c8 years ofXMM-Newton operations (Watson et al.

The 2XMMi catalogue identifies all X-ray sources detected in the 3491 observations made during the first $\approx 8$ years of operations (Watson et al.

2009).

Its unprecedented sky coverage (360 deg?) and sensitivity (median exposures of 20-50 ks) currently provides an exceptional resource for the unbiased identification of obscured AGN activity throughout the Universe.

Its unprecedented sky coverage (360 $^{2}$ ) and sensitivity (median exposures of 20–50 ks) currently provides an exceptional resource for the unbiased identification of obscured AGN activity throughout the Universe.

Using an automated reduction and analysis pipeline, the 2XMMi catalogue provides source positions,

observations of MCG—6-30-15 of eeV (appropriate for the above choice of abundances: see Reynolds. Fabian Inoue 1995).

observations of $-$ of eV (appropriate for the above choice of abundances; see Reynolds, Fabian Inoue 1995).

As suspected on the basis of the simple power-law fits in Table |. Fig.

As suspected on the basis of the simple power-law fits in Table 1, Fig.

7 (using a more complicated fit that includes the reflected spectrum) shows that the steepening in primary photon index during the brightest period of data is statistically significant.

There is also a suggestion of an increase in reflective fraction (with constant Fy) when going from the intermediate flux state o the low flux state.

There is also a suggestion of an increase in reflective fraction (with constant $\Gamma$ ) when going from the intermediate flux state to the low flux state.

Thus it appears that both changes in the oroperties of the primary X-ray source and changes in the amount of reflection are relevant for understanding spectral variability.

Thus it appears that both changes in the properties of the primary X-ray source and changes in the amount of reflection are relevant for understanding spectral variability.

As a urther check on the robustness of these results. we have recreated contours similar to those in Fig.

As a further check on the robustness of these results, we have recreated contours similar to those in Fig.

7 for several choices of abundances within the 68 per cent confidence statistical error range. and find hat none of the results changes materially based on these choices of abundance values.

7 for several choices of abundances within the 68 per cent confidence statistical error range, and find that none of the results changes materially based on these choices of abundance values.

The diserepaney in LP between 13 and the other two states does diminish if the abundance is allowed to be a free parameter.

The discrepancy in $\Gamma$ between i3 and the other two states does diminish if the abundance is allowed to be a free parameter.

It should be noted that the model as currently implemented does not include relativistic blurring of the reflection component.

À major caveat to RXTE spectral variability results lies in the background models.

A major caveat to RXTE spectral variability results lies in the background models.

The purpose of this paper was to show what unanswered questions can be addressed with the large area and wide-band coverage of even with the current uncertainties in spectral calibration.

The presence of a broad iron line is clearly evident as shown with a simple power law fit. and is one of the first detections where both features are seen simultaneously.

The presence of a broad iron line is clearly evident as shown with a simple power law fit, and is one of the first detections where both features are seen simultaneously.

We add a reflection component to our power law and gaussian fit to find that reflection is necessary to describe our data.

We note also that the steep intrinsic photon index coupled with a narrow //:7 FWHM implies that — 6-30-15 can be a possible narrow-line Seyfert | galaxy candidate.

We note also that the steep intrinsic photon index coupled with a narrow $H \beta$ FWHM implies that $-$ 6-30-15 can be a possible narrow-line Seyfert 1 galaxy candidate.

While spectral results may change in detail over the course of the next year with further improvements in calibration. we can already begin to place upper bound limits on the relationship between abundance values and reflective fraction.

While spectral results may change in detail over the course of the next year with further improvements in calibration, we can already begin to place upper bound limits on the relationship between abundance values and reflective fraction.

In Section 4. we study the effects of temporal variability on spectral components and find evidence to support the notion that variability may be due to changes in the amount of reflection seen (e.g. due to gravitational or Doppler beaming of the primary emission towards the disk).

In Section 4, we study the effects of temporal variability on spectral components and find evidence to support the notion that variability may be due to changes in the amount of reflection seen (e.g. due to gravitational or Doppler beaming of the primary emission towards the disk).

It is however not clear whether this effect may also be coupled with contributions from changes in the properties of the source itself (e.g. the temperature and optical depth of the coronal plasma).

We expect to be able to resolve these issues better with longer looks and simultaneous. ASCA observations.

We expect to be able to resolve these issues better with longer looks and simultaneous ASCA observations.

For the time being. the present results are important observational first steps in understanding some of the physics of AGN reprocessing mechanisms. and push the limits of our knowledge.

For the time being, the present results are important observational first steps in understanding some of the physics of AGN reprocessing mechanisms, and push the limits of our knowledge.

We thank all the members of the RXTE GOF for answering our inquiries in such a timely manner. with special thanks to Keith Jahoda for explanations of calibration issues.

We thank all the members of the RXTE GOF for answering our inquiries in such a timely manner, with special thanks to Keith Jahoda for explanations of calibration issues.

ICL thanks the Isaac Newton Trust. the Overseas Research Studentship programme (ORS) and the Cambridge Commonwealth Trust for support.

JCL thanks the Isaac Newton Trust, the Overseas Research Studentship programme (ORS) and the Cambridge Commonwealth Trust for support.

ACF thanks the Royal Society for support.

CSR thanks the National Science Foundation for support under grant AST9529175. and NASA for support under the Long Term Space Astrophysics grant NASA-NAG-6337.

CSR thanks the National Science Foundation for support under grant AST9529175, and NASA for support under the Long Term Space Astrophysics grant NASA-NAG-6337.

KI and WNB thank PPARC and NASA RXTE grant NAGS-685? for support respectively.

KI and WNB thank PPARC and NASA RXTE grant NAG5-6852 for support respectively.

expanding ejecta composed of nuu nüurshells with a wide-range of Lorentz factors are launched by the ceutral engine.

expanding ejecta composed of many mini-shells with a wide-range of Lorentz factors are launched by the central engine.

Tuternal shocks(Rees&Mészáros1991) are formed during the collisions of those shells aud produce the observed prompt GRB enmüssou (mostlv in (κατα band).

Internal shocks\citep{rees94} are formed during the collisions of those shells and produce the observed prompt GRB emission (mostly in Gamma-ray band).

Observationally this is the phase when GRBs trigger gamma-ray baud detectors. (

Observationally this is the phase when GRBs trigger gamma-ray band detectors. (

3) The ejecta are further decelerated by an ambient mediun (e.g. tuterstellar medium: ISAD) aud produce a lone-term broadband afterglow through an externaltforward shock (Mésszárros Rees 1997: Sari et al.

3) The ejecta are further decelerated by an ambient medium (e.g., interstellar medium; ISM) and produce a long-term broadband afterglow through an external-forward shock (Mésszárros Rees 1997; Sari et al.

1998) aucd/or externalreverse shock (Mésszárros Rees 1997. 1999: Sui Pian 19994.b). C

1998) and/or external-reverse shock (Mésszárros Rees 1997, 1999; Sari Piran 1999a,b). (

D) Th some cases. the ceutral cneine can be restarted during the afterglow phase and N-vav flares are produced through dissipation of a late wind launched from a long-lastiug central eugiue (Biurvows et al.

4) In some cases, the central engine can be restarted during the afterglow phase and X-ray flares are produced through dissipation of a late wind launched from a long-lasting central engine (Burrows et al.

2005a: Zhane et al.

2005a; Zhang et al.

2006: Fan Wei 2005:Re] Joka et al.

2006; Fan Wei 2005; Ioka et al.

2005: Wir et al.

2005; Wu et al.