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. Unless we state otherwise. (his rate (denoted as 74) is adopted in our calculations. | Unless we state otherwise, this rate (denoted as $r_{1}$ ) is adopted in our calculations. |
A higher estimate of 0.43Mpe‘Myr|! (denoted as r5) was obtained by assuming that all svstems survive early merger during the common envelope stage. | A higher estimate of $0.43\,\rm{Mpc}^{-3}\rm{Myr}^{-1}$ (denoted as $r_{2}$ ) was obtained by assuming that all systems survive early merger during the common envelope stage. |
We note that 7»leads to a detection rate for initial LIGO/Virgo of around 5 events per vear. | We note that $r_{2}$leads to a detection rate for initial LIGO/Virgo of around 5 events per year. |
A recent population | A recent population |
the dust is significantly settled (Le. 41=0.5). the bluest optically thick (uncleared) disc has a value of Ix - S] of around 1.2. | the dust is significantly settled (i.e. $H_D=0.5$ ), the bluest optically thick (uncleared) disc has a value of K - [8] of around 1.2. |
Only models with 445~0.1 can reach Ix - ὃς 1.2 and these represent a very extreme case of settling. | Only models with $H_D \sim 0.1$ can reach K - [8] $<$ 1.2 and these represent a very extreme case of settling. |
When these optically thick farec dises are subject to uniform draining they evolve (for XxRR+) alone the dashed lines shown. | When these optically thick flared discs are subject to uniform draining they evolve (for $\Sigma \propto R^{-1}$ ) along the dashed lines shown. |
As long as the disc is optically thick in the optical. the models are still redder (for à given surface density normalisation) than the fat. dise mocoels. since the geometry of the optical photosphere sets the temperature distribution in the disc. | As long as the disc is optically thick in the optical, the models are still redder (for a given surface density normalisation) than the flat disc models, since the geometry of the optical photosphere sets the temperature distribution in the disc. |
It is only at the point that the disc becomes optically thin also to the star's optical radiation that clise geometry no longer alfects the disc colours. | It is only at the point that the disc becomes optically thin also to the star's optical radiation that disc geometry no longer affects the disc colours. |
What is notable. however. is that (although the Dared cise models are redder than the flat dises in both colours for a given surface density profile) both flared and Hat draining models lic along the same trajectory in colour space. | What is notable, however, is that (although the flared disc models are redder than the flat discs in both colours for a given surface density profile) both flared and flat draining models lie along the same trajectory in colour space. |
7. We have therefore shown that there exists a well defined. draining sequence such that disces with a wide range of vertical structure and inclinations lie along this trajectory. | We have therefore shown that there exists a well defined draining sequence such that discs with a wide range of vertical structure and inclinations lie along this trajectory. |
The tracks for flared discs of various inclinations that are subject to clearing through the progressive erowth of inner holes are qualitatively similar to the tracks with inner holes for Hat discs except that. as expected. they are receler at Ix - 24]. | The tracks for flared discs of various inclinations that are subject to clearing through the progressive growth of inner holes are qualitatively similar to the tracks with inner holes for flat discs except that, as expected, they are redder at K - [24]. |
In Figure 3 we use the models computed above in order to Classify various regions of the two colour plane. | In Figure 3 we use the models computed above in order to classify various regions of the two colour plane. |
Regions A and B correspond to optically thick disces and. pure stellar photospheres. | Regions A and B correspond to optically thick discs and pure stellar photospheres. |
The blueward extent of region A depends on the disc geometry: however even. untlared disces (ie. with constant opening angle: b= 1) have Ix -. S] values L2 if the star is optically visible. unless the degree of dust settling is extreme (ic. the case of {1= 0.1). | The blueward extent of region A depends on the disc geometry: however even unflared discs (i.e. with constant opening angle: $b=1$ ) have K - [8] values $> 1.2$ if the star is optically visible, unless the degree of dust settling is extreme (i.e. the case of $H_D = 0.1$ ). |
The redward extent of region B depends on the spectral type: the solid. lines represent the enipivically derived. infrared. colours of an M5 star presented in the Appendix of Luhman et al (2010). ( | The redward extent of region B depends on the spectral type: the solid lines represent the empirically derived infrared colours of an M5 star presented in the Appendix of Luhman et al (2010). ( |
Ix - bi = (46. Ix - 24] = 0.69). | K - [8] = 0.46, K - [24] = 0.69). |
The clotted line encompasses our estimate of the uncertainties in this limit (see Section 3). | The dotted line encompasses our estimate of the uncertainties in this limit (see Section 3). |
Systems which lie outside these regions are "partially cleared: disces. apart from the small wedge of colour space (region €) which corresponds to "ultra-settled? clises viewed close to edge-on. C | Systems which lie outside these regions are `partially cleared discs', apart from the small wedge of colour space (region C) which corresponds to `ultra-settled' discs viewed close to edge-on. ( |
lUltra-settled denotes the case that the ratio of dust to gas scale height //p.« 0.1). | `Ultra-settled' denotes the case that the ratio of dust to gas scale height $H_D < 0.1$ ). |
legion. D contains objects that are compatible with being sources with cleared inner holes. | Region D contains objects that are compatible with being sources with cleared inner holes. |
The strip I5 is the ‘clraining locus’. which corresponds to the trajectory of discs in which the column density of dust is successively reduced by a (spatially constant) factor. | The strip E is the `draining locus', which corresponds to the trajectory of discs in which the column density of dust is successively reduced by a (spatially constant) factor. |
3clow we use this plot in order to classify the clise bearing properties of late type stars in IC 348. | Below we use this plot in order to classify the disc bearing properties of late type stars in IC 348. |
We now turn to the question of how we should interpret 10 various categories of partially cleared. dises classified by »evious authors. | We now turn to the question of how we should interpret the various categories of partially cleared discs classified by previous authors. |
Since many star forming regions (notably Taurus) contain relatively few disces that occupy the partially cleared. region. we here select. a cluster in which it has ovn claimed. that there is a relatively large population of xwtiallvy cleared disces (see Currie Ixenvon 2010. Luhman al 2010 for contrasting views on this issue). | Since many star forming regions (notably Taurus) contain relatively few discs that occupy the partially cleared region, we here select a cluster in which it has been claimed that there is a relatively large population of partially cleared discs (see Currie Kenyon 2010, Luhman et al 2010 for contrasting views on this issue). |
We presen statistical analysis of LC 348 in Section 3 below ane rere restrict ourselves to some exemplary. cases Cron. this cluster. | We present a statistical analysis of IC 348 in Section 3 below and here restrict ourselves to some exemplary cases from this cluster. |
Figure 3 combines the classification derived. above with Spitzer colours of the stars in LC 348 in the spectra vpe range M3 to ALS. dereddened: using the ο values of Lada et al (2006) ancl Müuench et al (2007) and the redadening aw of Rieke Lebofskv (1985). | Figure 3 combines the classification derived above with Spitzer colours of the stars in IC 348 in the spectral type range M3 to M5, dereddened using the $A_v$ values of Lada et al (2006) and Muench et al (2007) and the reddening law of Rieke Lebofsky (1985). |
We have. where available. revised the photometry obtained by Laca et al (2006) in accordance with the more sensitive measurements (ane ALPS photometry) obtained by Currie Ixenvon. (2009). | We have, where available, revised the photometry obtained by Lada et al (2006) in accordance with the more sensitive measurements (and MIPS photometry) obtained by Currie Kenyon (2009). |
The black svmbols represent cliscs classified as (optically) thick clises by Lada ct al (2006). the green svmbols are the ‘anaemic clises’ of Lada οἱ al (2006) and the blue clots are systems classified: as photospheres (on the basis of LRAC data alone). | The black symbols represent discs classified as (optically) thick discs by Lada et al (2006), the green symbols are the `anaemic discs' of Lada et al (2006) and the blue dots are systems classified as photospheres (on the basis of IRAC data alone). |
Ehe τος triangles are the sources observed. by Aluench et al (2007). | The red triangles are the sources observed by Muench et al (2007). |
The horizontal arrows represent the sources for which only upper limits are available at 244m which are unfortunately quite numerous in this cluster owing to the bright background emission. | The horizontal arrows represent the sources for which only upper limits are available at $24 \mu$ m which are unfortunately quite numerous in this cluster owing to the bright background emission. |
The red. vellow and blue asterisks denote sources that have variously been classified as weak excess’. warn excess’ and ‘classical transition’ (1.6. inner hole) sources by Muzerolle et al (2010). | The red, yellow and blue asterisks denote sources that have variously been classified as `weak excess', `warm excess' and `classical transition' (i.e. inner hole) sources by Muzerolle et al (2010). |
The sources with an outer red circle are those classified as “homologously depleted! by Currie Ixenvon (2009). | The sources with an outer red circle are those classified as `homologously depleted' by Currie Kenyon (2009). |
Lt is immediately obvious from this plot that the various classification schemes overlap and that we have sources that are associated with a variety of designations. | It is immediately obvious from this plot that the various classification schemes overlap and that we have sources that are associated with a variety of designations. |
Thus there are ped. asterisks with black borders (objects counted as weak excess by Muzerolle et al 2010. but as ‘thick discs! by Lada et al 2006) which have very similar colours to other sources that were counted as ‘anaemic’ by Lada et al (2006). | Thus there are red asterisks with black borders (objects counted as weak excess by Muzerolle et al 2010, but as `thick discs' by Lada et al 2006) which have very similar colours to other sources that were counted as `anaemic' by Lada et al (2006). |
Other anaemic sources (those with red border) are described: as ‘homologously depleted? by Currie Ixenson. (2009). | Other anaemic sources (those with red border) are described as `homologously depleted' by Currie Kenyon (2009). |
Llere we draw attention to three aspects of this plot: i) Almost all the “weak excess! disces of Muzerolle et al (2010) lie close to the upper end of the line FE in Figure 1. implving that they are describable as discs at a range of (relatively face on) inclinations (i « 60deg). | Here we draw attention to three aspects of this plot: i) Almost all the `weak excess' discs of Muzerolle et al (2010) lie close to the upper end of the line FE in Figure 1, implying that they are describable as at a range of (relatively face on) inclinations (i $< 60 \deg$ ). |
In order to check this. we have fitted these sources as such disces in which the only free parameter is the source inclination: the results of this exercise are shown in ligure 4 and demonstrate excellent agreement in 5/9 cases (namely sources LIRLLI35. LRLLITO. LRLL213. LRLL229 and LILL241). | In order to check this, we have fitted these sources as such discs in which the only free parameter is the source inclination; the results of this exercise are shown in Figure 4 and demonstrate excellent agreement in 5/9 cases (namely sources LRLL135, LRLL176, LRLL213, LRLL229 and LRLL241). |
Note that this is not à unique interpretation: these same sources could also be fit bv finite thickness | Note that this is not a unique interpretation: these same sources could also be fit by finite thickness |
In recent vears. the Sloan Digital Skv Survey (Yorketal.2000) has more than doubled ihe number of known dwarl galaxies that orbit the Galaxy (Willmanetal.2005a.b:Zucker2007:Walshetal.Belokuroy2010:Willman 2010).. | In recent years, the Sloan Digital Sky Survey \citep{York2000} has more than doubled the number of known dwarf galaxies that orbit the Galaxy \citep{Willman2005,Willman2005a,Zucker2006,Belokurov2006,Zucker2006a,Belokurov2007,Irwin2007,Walsh2007,Belokurov2010,Willman2010}. |
These newly discovered | These newly discovered |
9970508 (at órz100 d) from which the kinetic energy was inferred to be Ey~5«10?" erg. | 970508 (at $\delta t\gtrsim
100$ d) from which the kinetic energy was inferred to be $E_K\sim
5\times 10^{50}$ erg. |
Bergeretal.(2004) used the same approach to model the radio afterglow emission of 9980703 on timescales of z40 d. and to re-model 9970508. | \citet{bkf04} used the same approach to model the radio afterglow emission of 980703 on timescales of $\gtrsim
40$ d, and to re-model 970508. |
They found kinetic energies of Ex~3«10°! erg for both bursts. | They found kinetic energies of $E_K\sim 3\times 10^{51}$ erg for both bursts. |
Finally. Frailetal.(2005) modeled the radio emission from 0030329 at ὃνz50 d and found Ex~10°! ere. | Finally, \citet{fsk+05} modeled the radio emission from 030329 at $\delta
t\gtrsim 50$ d and found $E_K\sim 10^{51}$ erg. |
Only 3 bursts have been studied in this fashion so far because only those events have well-sampled radio light curves on the relevant timescales of 0t=100 d. However. the kinetic energy can still be estimated using the same methodology even from fragmentary late-time radio observations. | Only 3 bursts have been studied in this fashion so far because only those events have well-sampled radio light curves on the relevant timescales of $\delta
t\gtrsim 100$ d. However, the kinetic energy can still be estimated using the same methodology even from fragmentary late-time radio observations. |
Such an approach will naturally result in larger uncertainties for each burst. but it can be applied to à much larger sample of events. | Such an approach will naturally result in larger uncertainties for each burst, but it can be applied to a much larger sample of events. |
Here we present such an analysis for 24 long-duration GRBs with radio observations at 2100 d. but with only 1—3 data points (at 1.4 to 8.5 GHz) per burst. | Here we present such an analysis for $24$ long-duration GRBs with radio observations at $\gtrsim 100$ d, but with only $1-3$ data points (at 1.4 to 8.5 GHz) per burst. |
σιsing these observations we infer robust ranges for the kinetic energy of each burst and for the population as a whole. | Using these observations we infer robust ranges for the kinetic energy of each burst and for the population as a whole. |
The plan of the paper is as follows. | The plan of the paper is as follows. |
The radio observations are summarized in refsec:obs.. | The radio observations are summarized in \\ref{sec:obs}. |
The model for synchrotron emission from a Sedov-Taylor blastwave. and the various assumptions we employ are presented in refsec:model.. | The model for synchrotron emission from a Sedov-Taylor blastwave, and the various assumptions we employ are presented in \\ref{sec:model}. |
In refsec:res we detail the resulting kinetic energies and the range for the overall sample. and we compare these results to multi-wavelength analyses of early afterglows in refsec:comp.. | In \\ref{sec:res} we detail the resulting kinetic energies and the range for the overall sample, and we compare these results to multi-wavelength analyses of early afterglows in \\ref{sec:comp}. |
We conclude with a discussion of future prospects. | We conclude with a discussion of future prospects. |
We use radio observations of 24 long GRBs at ór=100 d since on those timescales the blastwave is expected to become non-relativistic and roughly isotropic (Livio&Wax-man 2000).. and the peak of the afterglow emission 1s at or below the centimeter band. | We use radio observations of $24$ long GRBs at $\delta t\gtrsim 100$ d since on those timescales the blastwave is expected to become non-relativistic and roughly isotropic \citep{lw00}, and the peak of the afterglow emission is at or below the centimeter band. |
This has been confirmed with detailed data in the case of GRBs 970508. 980703. and 030329 (Frailetal.2000:Berger2004:Frail 2005). | This has been confirmed with detailed data in the case of GRBs 970508, 980703, and 030329 \citep{fwk00,bkf04,fsk+05}. |
. We restrict the analysis to GRBs with a known redshift and with early-time detections. which for the case of a single detection or upper limit allow us to infer that the peak of the spectrum has transitioned below our observing frequency. | We restrict the analysis to GRBs with a known redshift and with early-time detections, which for the case of a single detection or upper limit allow us to infer that the peak of the spectrum has transitioned below our observing frequency. |
The observations are primarily from the Very Large Array Inc.)). with the exception of GRBs 980425 and 011121. which were observed with the Australia Telescope Compact Array (ATCA). | The observations are primarily from the Very Large Array ), with the exception of GRBs 980425 and 011121 which were observed with the Australia Telescope Compact Array (ATCA). |
The data were obtained between 1997 and 2009 as part of a long-term GRB radio program (e.g.. 20030). | The data were obtained between 1997 and 2009 as part of a long-term GRB radio program (e.g., \citealt{fkb+03}) ). |
For the purpose of our analysis. we separate the bursts into three categories based on the quality of the data. | For the purpose of our analysis, we separate the bursts into three categories based on the quality of the data. |
In Group A are 3 bursts with late-time detections at multiple frequencies that constrain the peak of the synchrotron spectrum (the same three bursts that have been studied in detail by Frailetal.2000:Bergeretal.2004:Frail 2005)). | In Group A are 3 bursts with late-time detections at multiple frequencies that constrain the peak of the synchrotron spectrum (the same three bursts that have been studied in detail by \citealt{fwk00,bkf04,fsk+05}) ). |
In Group B are 11 bursts with single-frequency detections. while Group C consists of 10 GRBs with late-time non-detections. | In Group B are 11 bursts with single-frequency detections, while Group C consists of 10 GRBs with late-time non-detections. |
The VLA measurements and relevant burst properties are listed in table Τ.. | The VLA measurements and relevant burst properties are listed in table \ref{tab:data}. |
Our modeling of the radio data follows the methodology of Frailetal.(2000) and Bergeretal.(2004) for the case ofa uniform densitycase. | Our modeling of the radio data follows the methodology of \citet{fwk00} and \citet{bkf04} for the case ofa uniform density. |
.. For the typical expected. parameters of long GRBs. the initially collimated blastwave approaches spherical symmetry and decelerates to non-relativistic velocity on similar timescales. 1©10sso/8 d and typ JOCERiss/nMre d. respectivelyΕκ (Livio&Waxman2000): here. n, 15 the circumburst density in units of em and r; is the "jet break" time at which the jet begins to expand sideways (1.e.. (rj)9i. where Τ is the bulk Lorentz factor). | For the typical expected parameters of long GRBs, the initially collimated blastwave approaches spherical symmetry and decelerates to non-relativistic velocity on similar timescales, $t_s\approx 150(E_{\rm K,iso,52}/n_e)^{1/4} t_{\rm j,d}^{1/4}$ d and $t_{NR}\approx 40(E_{\rm K,iso,52}/n_e)^{1/4}t_{\rm j,d}^{1/4}$ d, respectively \citep{lw00}; here, $n_e$ is the circumburst density in units of $^{-3}$ and $t_j$ is the “jet break” time at which the jet begins to expand sideways (i.e., $\Gamma(t_j)\sim\theta_j^{-1}$ , where $\Gamma$ is the bulk Lorentz factor). |
In this paper we assume that the blastwave has transitioned to the non-relativistic isotropic phase by the time of our observations and subsequently check forself-consistency. | In this paper we assume that the blastwave has transitioned to the non-relativistic isotropic phase by the time of our observations and subsequently check forself-consistency. |
The blastwave dynamies in the non-relativistic phase are described by the Sedov-Taylor self-similar solution with rt)x(Est? /m'!>. | The blastwave dynamics in the non-relativistic phase are described by the Sedov-Taylor self-similar solution with $r(t)\propto (E_{\rm ST}
t^2/n)^{1/5}$ . |
To calculate the synchrotron emission emerging from the shock-heated material. we make the usual assumptions: (1) the electrons are accelerated to a power-law energy distribution. N(5)x37? for 5>54. where 5 is the minimum Lorentz factor: (11) the value of p is 2.2 as inferred from several bursts (e.g.. Panaitescu&Kumar2001.2002:Yostetal. 2003): and (ii) the energy densities in the magnetic field and electrons are constant fractions(ερ and e,. respectively) of the shock energy density. | To calculate the synchrotron emission emerging from the shock-heated material, we make the usual assumptions: (i) the electrons are accelerated to a power-law energy distribution, $N(\gamma)\propto\gamma^{-p}$ for $\gamma>\gamma_m$, where $\gamma_m$ is the minimum Lorentz factor; (ii) the value of $p$ is 2.2 as inferred from several bursts (e.g., \citealt{pk01,pk02,yhs+03}) ); and (iii) the energy densities in the magnetic field and electrons are constant fractions$\epsilon_B$ and $\epsilon_e$, respectively) of the shock energy density. |
Accounting for synchrotron emissivity and self-absorption. and including the appropriate redshift transformations. the flux observed at frequency 7 and time f£ is given by (Frailetal. 2004): where the optical depth is given by: the synchrotron peak frequency. corresponding to electrons with 7255. is given by: and the function fj;Go) Is given by where F(v) is an integration over Bessel functions (RybickiLightman 1979). | Accounting for synchrotron emissivity and self-absorption, and including the appropriate redshift transformations, the flux observed at frequency $\nu$ and time $t$ is given by \citep{fwk00,bkf04}: where the optical depth is given by: the synchrotron peak frequency, corresponding to electrons with $\gamma=\gamma_m$, is given by: and the function $f_l(x)$ is given by where $F(y)$ is an integration over Bessel functions \citep{rl79}. |
. The temporal indices in the case of auniform density mediumare a,=11/10. 0-=1—3p/2. and à2—3. | The temporal indices in the case of auniform density mediumare $\alpha_F=11/10$, $\alpha_\tau=1-3p/2$, and $\alpha_m=-3$. |
The normalizations are such that Fo and τα are the flux density and optical depth at a frequency of 21 Hz att fj. and 14 is the synchrotron peak frequency in the rest frame of burst at 2 fo. | The normalizations are such that $F_0$ and $\tau_0$ are the flux density and optical depth at a frequency of $\nu=1$ Hz at $t=t_0$ , and $\nu_0$ is the synchrotron peak frequency in the rest frame of burst at $t=t_0$ . |
Furthermore. the synchrotron self-absorption frequency. 7. is defined by the condition τε)= 1. | Furthermore, the synchrotron self-absorption frequency, $\nu_a$ , is defined by the condition $\tau_\nu(\nu_a)=1$ . |
We fit this synchrotron model to our radio data using Fo. το. and το as free parameters. | We fit this synchrotron model to our radio data using $F_0$ , $\tau_0$ , and $\nu_0$ as free parameters. |
Since we have no knowledge about the expected values of the synchrotron spectrum parameters we assume that they follow a flat distribution in log-space. | Since we have no knowledge about the expected values of the synchrotron spectrum parameters we assume that they follow a flat distribution in log-space. |
Wenote that any further assumption | Wenote that any further assumption |
19 degrees of freedom. | 19 degrees of freedom. |
The amplitudes (125) iu the power law aud exponential compoucuts. imteerated from 10? to 100 Tz. ave. Lbd:0.354 and 5.7£0. respectively, | The amplitudes (rms) in the power law and exponential components, integrated from $10^{-3}$ to 100 Hz, are $4.4 \pm 0.3 \%$ and $5.7 \pm 0.4 \%$, respectively. |
We also searched. for kz quasi-periodic oscillations (QPO) but found no significant features in the 200 - 1200 IIz rango. | We also searched for kHz quasi-periodic oscillations (QPO) but found no significant features in the 200 - 1200 Hz range. |
The N-ray variability revealed by ls qualitatively very simular to tha seen in other N-raw binaries (see for example Wijuands van der Klis 1999). showing a broad excess UMnes clearly a OPO other times properly called a bunip) superposed on a broad band power law component. | The X-ray variability revealed by is qualitatively very similar to that seen in other X-ray binaries (see for example Wijnands van der Klis 1999), showing a broad excess (sometimes clearly a QPO other times properly called a 'bump') superposed on a broad band power law component. |
Unfortunately. this behavior is exlübited bv both black hole as well as neutron star SVSTCLUS. so it is not possible to distinguish the nature of the compact source based on the broad band N-ray variability alone. | Unfortunately, this behavior is exhibited by both black hole as well as neutron star systems, so it is not possible to distinguish the nature of the compact source based on the broad band X-ray variability alone. |
Louger observations will be required to make more sensitive measurements to search for neutron star signatures such as kIIz QPO. | Longer observations will be required to make more sensitive measurements to search for neutron star signatures such as kHz QPO. |
We note that no eclipse was seen with the PCA. either diving the TOO discussed in this section. or during the monitoring observations (see Fig. 2)). | We note that no eclipse was seen with the PCA, either during the TOO discussed in this section, or during the monitoring observations (see Fig. \ref{figpcalc}) ). |
Despite the lack of an optical ideutification. lis very likely a transicut LMXD1j)ecause it is located in a elobular cluster. | Despite the lack of an optical identification, is very likely a transient LMXB because it is located in a globular cluster. |
This is also supported by the 12.960 hr orbital period. | This is also supported by the 12.360 hr orbital period. |
The eclipse provides us with further constraints on the binary orbit. | The eclipse provides us with further constraints on the binary orbit. |
If one assumes that 1) the eclipse is caused by only the companion star: 2) the mass of the companion star is less than 0.5 (as is expected for a miain-sequence or (sub)eiaut star iu a elobular cluster): 3) the mass of the compact object is larger than 1.1 M. (Le. it must be heavier than or equal to that of a neutron star): and L) Roche geometry applies to the |iuuv. then the calculations by Torue (1985) iniply that the inclination angle is larger than | If one assumes that 1) the eclipse is caused by only the companion star; 2) the mass of the companion star is less than 0.8 $_\odot$ (as is expected for a main-sequence or (sub)giant star in a globular cluster); 3) the mass of the compact object is larger than 1.4 $_\odot$ (i.e., it must be heavier than or equal to that of a neutron star); and 4) Roche geometry applies to the binary, then the calculations by Horne (1985) imply that the inclination angle is larger than. |
The egress out of eclipse to ~90% of the out-ofeclipse level is 35 s. This duration is deteruunect by the size of the chussion region beige eclipsed and the sharpuess of the edge of the eclipsing olject. | The egress out of eclipse to $\sim$ of the out-of-eclipse level is 35 s. This duration is determined by the size of the emission region being eclipsed and the sharpness of the edge of the eclipsing object. |
Enhancements in Nyy curving ceress would poiut to absoption effects in the atmosphere of the companion star. | Enhancements in $N_{\rm H}$ during egress would point to absorption effects in the atmosphere of the companion star. |
Uufortunatelv. we are rot able to ueasure that. | Unfortunately, we are not able to measure that. |
We can oulv determine wpper Πές to the ransparent part of the companions atinosphere aud the size of the cussion region. | We can only determine upper limits to the transparent part of the companion's atmosphere and the size of the emission region. |
The upper luit ou the relative lückuess o ‘the ransparent part of the atinospliere is of the stellar radius. | The upper limit on the relative thickness of the transparent part of the atmosphere is of the stellar radius. |
The emissiou region shoul be sinaller han 2«4107 kin (as derive: if one applies I&eplers law uuder the assmuption that the combined mass of both ünarv comniponeuts is 2 M). | The emission region should be smaller than $2\times10^3$ km (as derived if one applies Kepler's law under the assumption that the combined mass of both binary components is 2 $_\odot$ ). |
By analoev to other eclipsing LMXDs nod(c.g. EXO 07Ls676. Darinar ct al. | By analogy to other eclipsing LMXBs (e.g. EXO 0748--676, Parmar et al. |
1986). if Is a hypothesis to attribute the residual cussion during the eclipse to photons scattered into the line o sieht by an accretion disk corona (ADC). plus perhaps a coutributiou of iuterstellar dust erains below 2 keV. The enismiou beine eclipsed would then result from à much nore confined region. very ikelv the immer acerction disk. | 1986), it is a good hypothesis to attribute the residual emission during the eclipse to photons scattered into the line of sight by an accretion disk corona (ADC), plus perhaps a contribution of interstellar dust grains below 2 keV. The emission being eclipsed would then result from a much more confined region, very likely the inner accretion disk. |
The flux coutributiion of he ADC outside the eclipse would1ο less than. | The flux contribution of the ADC outside the eclipse would be less than. |
Any nodulation introduced by the xwtial obscuration of the ADC outside the eclipse may. therefore. be masked by other variability of this trausieut source. | Any modulation introduced by the partial obscuration of the ADC outside the eclipse may, therefore, be masked by other variability of this transient source. |
Possibly. the slow wart of the egress as secu with the ΣΕΤ may be explained w this. | Possibly, the slow part of the egress as seen with the NFI may be explained by this. |
similar foreground. absorption. Ay~ Imag. and some clferential. redcdening. | similar foreground absorption, $\aV\sim1$ mag, and some differential reddening. |
‘Taken together. the presence of gas. dust. some VV anc VV stars. and the | PAIS CMD morphologics consistently constrain the age of the ECs (and 1. extended. stellar eroup) to less than zz10 MMwyr (the null of the stars are probably ~5 AlAIver old). with a time-oreacd of ~LO AAIvr for the star formation. | Taken together, the presence of gas, dust, some V and V stars, and the $+$ PMS CMD morphologies consistently constrain the age of the ECs (and the extended stellar group) to less than $\approx10$ Myr (the bulk of the stars are probably $\sim5$ Myr old), with a time-spread of $\sim10$ Myr for the star formation. |
Phe | PAIS gacollar masses are low. within zzGOAL. to zz200M... | The $+$ PMS stellar masses are low, within $\approx60\,\ms$ to $\approx200\,\ms$. |
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