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The combination of the high velocity relative to the galaxy, the low metallicity, and the large distance from the galaxy leads us to the conclusion that the halo cloud at v=+365—445kms! is an accreting, metal-poor gas cloud, providing fuel for future star formation in galaxy227.
The combination of the high velocity relative to the galaxy, the low metallicity, and the large distance from the galaxy leads us to the conclusion that the halo cloud at $v = +365-445\kms$ is an accreting, metal-poor gas cloud, providing fuel for future star formation in galaxy.
19.. The velocity separation between the ccloud and galaxy is Av=85kms!, which is typically close enough to consider the gas to exist in the galaxy halo.
The velocity separation between the cloud and galaxy is $\dv = 85\kms$, which is typically close enough to consider the gas to exist in the galaxy halo.
It is interesting to note that a cloud with such a low velocity in the MW halo may not even be classified as a halo cloud, given its low velocity separation from the galaxy.
It is interesting to note that a cloud with such a low velocity in the MW halo may not even be classified as a halo cloud, given its low velocity separation from the galaxy.
Rather, the ccloud would likelybe considered part of the ~3kpc high, ionized gas surrounding the thick disk (c.f.?),, as opposed to a distant halo cloud.
Rather, the cloud would likelybe considered part of the $\sim 3\kpc$ high, ionized gas surrounding the thick disk \citep[c.f.][]{savage-etal-03-OVI-disk-halo}, , as opposed to a distant halo cloud.
The detection of aat a distance d>95kpc from the disk may have implications for the interpretation of the MW hhalo clouds.
The detection of at a distance $d \geq 95\kpc$ from the disk may have implications for the interpretation of the MW halo clouds.
To interpret the ccloud, we considered the feasibility of several scenarios to explain physical conditions of the cloud.
To interpret the cloud, we considered the feasibility of several scenarios to explain physical conditions of the cloud.
In doing so, we adopt the single component metal-line column densities from the Voigt profile fitting in the previous section (Table 3)).
In doing so, we adopt the single component metal-line column densities from the Voigt profile fitting in the previous section (Table \ref{tab: profile_fitting}) ).
The principle source of error, which propagates to our derived physical quantities, is the large permitted range of ccolumn density.
The principle source of error, which propagates to our derived physical quantities, is the large permitted range of column density.
A pure collisional ionization equilibrium (CIE) model can be quickly excluded by examination of the absorption line widths and column density.
A pure collisional ionization equilibrium (CIE) model can be quickly excluded by examination of the absorption line widths and column density.
ttraces collisionally ionized gas in the range 5.2— 6.5, peaking at logT=5.5 (??).. a
traces collisionally ionized gas in the range $\logT = 5.2 - 6.5$ , peaking at $\logT = 5.5$ \citep{sutherland-dopita-93-gas-cooling, gnat-sternberg-07-non-equilibrium-cooling}.
and mmayC coexist at the low end of this range, but the tto rratio is a very steep function of temperature; the measured range in the column density ratio Novi/Noii=0.6—1.7 allows temperatures only in the narrow range logT'—5.255.3.
and may coexist at the low end of this range, but the to ratio is a very steep function of temperature; the measured range in the column density ratio $\NOVI/\NCIII = 0.6 - 1.7$ allows temperatures only in the narrow range $\logT = 5.25 - 5.3$.
At this temperature, the implied thermal doppler line-width for iis Dtherm=\/2kT/m5Tkms!, or a full-width at halfmaximum width of!.
At this temperature, the implied thermal doppler line-width for is $b_{therm} = \sqrt{2kT/m} = 57\kms$, or a full-width at half-maximum width of.
. This line width is clearly excluded by a visual inspection of the data; a Voigt profile fit (accounting for the COS line-spread function) yields by1~32kms+, confirming that a pure collisional ionization equilibrium model is ruled out.
This line width is clearly excluded by a visual inspection of the data; a Voigt profile fit (accounting for the COS line-spread function) yields $b_{\scHI} \simeq 32\kms$, confirming that a pure collisional ionization equilibrium model is ruled out.
The detection of both aandVI,, and the limits on aand ccolumn densities give us some handle on a simple photoionization scenario.
The detection of both and, and the limits on and column densities give us some handle on a simple photoionization scenario.
Simple plane-parallelCloudy photoionization models with the canonical HM05 spectrum can reproduce the observed metal line column densities and line ratios.
Simple plane-parallel photoionization models with the canonical HM05 spectrum can reproduce the observed metal line column densities and line ratios.
The ratio of tto ffixes the range of allowed ionization parameter (and hence density), since the column density ratio is independent of metallicity for a fixed abundance pattern, and the UV background is also fixed.
The ratio of to fixes the range of allowed ionization parameter (and hence density), since the column density ratio is independent of metallicity for a fixed abundance pattern, and the UV background is also fixed.
The allowed metallicity range can then be obtained using the permitted density ranges, and varying the metallicity until the observed column densities are obtained.
The allowed metallicity range can then be obtained using the permitted density ranges, and varying the metallicity until the observed column densities are obtained.
The observed conditions are reproduced for a model with logU=—1.1——0.4 and metallicity in the range [M/H]=—0.3——1.1 for a solar abundance pattern.
The observed conditions are reproduced for a model with $\logU = -1.1 - -0.4$ and metallicity in the range $\MH =-0.3 - -1.1$ for a solar abundance pattern.
Given the input HM05 spectrum is fixed at all points hence the number of ionizing photonsis also fixed), this (andlogU range corresponds to a density range logng— —4.9.
Given the input HM05 spectrum is fixed at all points (and hence the number of ionizing photonsis also fixed), this $\logU$ range corresponds to a density range $\loghden = -4.4 - -4.9$ .
However, this assumes both the shape and
However, this assumes both the shape and
and Strohmaver&Brown (2002).. Cumming&Bildsten(2001) and Brown(2004) do not explicitly state that Chev (rack the evolution of the carbon mass fraction Ζονο when they solve lor the equilibrium configuration of the laver.
and \citet{SB02}, , \citet{CB01} and \citet{B04} do not explicitly state that they track the evolution of the carbon mass fraction $Z_{\mathrm{CNO}}$ when they solve for the equilibrium configuration of the layer.
To our knowledge. these authors keep Zoxo constant throughout the laver. but we cannot state this with certainty.
To our knowledge, these authors keep $Z_{\mathrm{CNO}}$ constant throughout the layer, but we cannot state this with certainty.
To conduct an accurate comparison between our global analvsis and (he various approximations. we solve for the ecuilibrium configuration of the accreted laver in an identical fashion for both methods.
To conduct an accurate comparison between our global analysis and the various one-zone approximations, we solve for the equilibrium configuration of the accreted layer in an identical fashion for both methods.
This ensures that any differences in the results are cue onlv to the stabili caleulation.
This ensures that any differences in the results are due only to the stability calculation.
We find that cloe,../dla7zz2 in all of our caleulations.
We find that $\mathrm{d} \ln \epsilon_{\mathrm{cool}} / \mathrm{d} \ln T \approx 2$ in all of our calculations.
Therefore. setüng dlne,/dinf=2 is appropriate for all scenarios in our opinion.
Therefore, setting $\mathrm{d} \ln \epsilon_{\mathrm{cool}} / \mathrm{d} \ln T = 2$ is appropriate for all scenarios in our opinion.
In contrast. we find that setüng dnec/dl1nT=26 is appropriate only if the core temperature is high (2 105IX) or the aceretion rate is near the Eddington limit.
In contrast, we find that setting $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T = 26$ is appropriate only if the core temperature is high $\gtrsim 10^{8} $ K) or the accretion rate is near the Eddington limit.
At sufficiently high densities aud low temperatures. dInec/dln7«26. often bv many orders of magnitude. so (his approximation is inappropriate in (hese situations.
At sufficiently high densities and low temperatures, $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T \ll 26$, often by many orders of magnitude, so this approximation is inappropriate in these situations.
For a given equilibrium configuration. we calculate the one-zone approximation in four different wavs.
For a given equilibrium configuration, we calculate the one-zone approximation in four different ways.
When calculating dec/dZ at the base of the accreted laver we (1) use the value of Zexo and dInec/dIn7 derived from our equilibrium configuration. (i1) artificially set Zoxg constant aud use the value of dInec/dIn1 derived from our equilibrium configuration. ii) use (he value of Zexo derived [rom our equilibrium configuration and set 26. or (iv) artificially set Zexo constant and set dinec/dIn7=26.
When calculating $\mathrm{d} \epsilon_{\mathrm{C}} / \mathrm{d} T$ at the base of the accreted layer we (i) use the value of $Z_{\mathrm{CNO}}$ and $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T$ derived from our equilibrium configuration, (ii) artificially set $Z_{\mathrm{CNO}}$ constant and use the value of $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T$ derived from our equilibrium configuration, (iii) use the value of $Z_{\mathrm{CNO}}$ derived from our equilibrium configuration and set $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T = 26$ , or (iv) artificially set $Z_{\mathrm{CNO}}$ constant and set $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T = 26$.
Artificially setting Zexo constant results in superbursts at all accretion rates for each caleulation we performed.
Artificially setting $Z_{\mathrm{CNO}}$ constant results in superbursts at all accretion rates for each calculation we performed.
This is clearly incorrect since all our global stability analysis models indicate a minimum M below which superbursts are absent.
This is clearly incorrect since all our global stability analysis models indicate a minimum $\dot{M}$ below which superbursts are absent.
Therefore. the criteria Gi) and (iv) above are very inaccurate and should be avoided.
Therefore, the criteria (ii) and (iv) above are very inaccurate and should be avoided.
Regarding the other two criteria. we find Chat for neutron stars with hieh core temperatures. such as those with cores (hat radiate neutrinos via mocified URCA reactions. the results from the ealeulations in which πες111 is set to 26 and in which dlnec/din is caleulatecl sell-consistentlv are almost identical.
Regarding the other two criteria, we find that for neutron stars with high core temperatures, such as those with cores that radiate neutrinos via modified URCA reactions, the results from the calculations in which $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T$ is set to $26$ and in which $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T$ is calculated self-consistently are almost identical.
Therefore. the approximation dine/dln—26 is valid in (his situation.
Therefore, the approximation $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T = 26$ is valid in this situation.
HLowever. lor neutron stars wilh low core temperatures. such as those with cores that racliale neutrinos via pionic reactions. the results from the (wo calculations differ significantly, so in this case dine/dlaZ must be caleulatecl sell-consistentlv.
However, for neutron stars with low core temperatures, such as those with cores that radiate neutrinos via pionic reactions, the results from the two calculations differ significantly, so in this case $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T$ must be calculated self-consistently.
see Figures 17 and 18 [or a comparison between (he results from our global linear stability analvsis and the one-zone approximations (1) and (iii).
See Figures 17 and 18 for a comparison between the results from our global linear stability analysis and the one-zone approximations (i) and (iii).
For the calculations in which both Zexo and dinec/dinZ' are calculated sell-consistenilv. the results of the one-zone approximation and the global linear stability analvsis generally agree «quite well.
For the calculations in which both $Z_{\mathrm{CNO}}$ and $\mathrm{d} \ln \epsilon_{\mathrm{C}} / \mathrm{d} \ln T$ are calculated self-consistently, the results of the one-zone approximation and the global linear stability analysis generally agree quite well.
The errors produced by the one-zone analvsis are probably smaller (han those due to uncertainties in other parameters. such as (he accretionrate. impurity concentration. elemental composition
The errors produced by the one-zone analysis are probably smaller than those due to uncertainties in other parameters, such as the accretionrate, impurity concentration, elemental composition
legacy project. MIPSGAL (?)..
legacy project, MIPSGAL \citep{care09}.
The 6" beam of Spitzer at this wavelength and its much better sensitivity 15 a significant improvement to that of the Midcourse Source Experiment satellite (MSX: ?)).
The $''$ beam of Spitzer at this wavelength and its much better sensitivity is a significant improvement to that of the Midcourse Source Experiment satellite (MSX; \citealt{pric99}) ).
In this paper. we present multi-wavelength observations from centimeter to submillimeter wavelengths towards a sample of 20 6.7 GHz methanol masers selected from a blind survey.
In this paper, we present multi-wavelength observations from centimeter to submillimeter wavelengths towards a sample of 20 6.7 GHz methanol masers selected from a blind survey.
We combine our observations with existing data to determine SEDs from centimeter to near-infrared wavelengths.
We combine our observations with existing data to determine SEDs from centimeter to near-infrared wavelengths.
The combination of the radio continuum data and the results from SED modeling give very strong evidence that the methanol masers are associated with rapidly acereting massive stars.
The combination of the radio continuum data and the results from SED modeling give very strong evidence that the methanol masers are associated with rapidly accreting massive stars.
To determine the spectral energy distributions of 6.7 GHz methanol masers. we observed a sample of 20 masers with the Very Large (VLA) at 3.6 em. 1.3 em and 6.9 mm. the IRAM 30-meter telescope at 1.2 mm and the Atacama Path Finder Experiment telescope. citepgues06.. at 870 gm. The source sample was selected from the AMGPS and included all sources detected between Galactic longitudes of 38.67 and 43.17. and Galactic latitudes |b]€0.42%. the latitude constraint being imposed for the sake of completeness.
To determine the spectral energy distributions of 6.7 GHz methanol masers, we observed a sample of 20 masers with the Very Large (VLA) at 3.6 cm, 1.3 cm and 6.9 mm, the IRAM 30-meter telescope at 1.2 mm and the Atacama Path Finder Experiment telescope, \\citep{gues06}, at 870 $\mu$ m. The source sample was selected from the AMGPS and included all sources detected between Galactic longitudes of $38.6^\circ$ and $43.1^\circ$, and Galactic latitudes $|b| \leq 0.42^\circ$, the latitude constraint being imposed for the sake of completeness.
The peak flux density of the masers range from 0.2 Jy to 26 Jy. and their luminosities (takenfrom?) range from 9.9x107 to 2.6x107? Les.
The peak flux density of the masers range from 0.2 Jy to 26 Jy, and their luminosities \citep[taken from][]{pand09} range from $9.9 \times 10^{-8}$ to $2.6 \times 10^{-5}~L_\odot$ .
14: out of 20 sources have accurate positions (astrometry better than 0.1) determined from phase-referenced observations with the MERLIN interferometer (these observations. were part of a larger project to determine accurate positions. to the AMGPS sources. and will be described in detail in Pandian et al.
14 out of 20 sources have accurate positions (astrometry better than $''$ ) determined from phase-referenced observations with the MERLIN interferometer (these observations were part of a larger project to determine accurate positions to the AMGPS sources and will be described in detail in Pandian et al.
2010. in. preparation).
2010, in preparation).
The remaining sources have positions determined from 24 um point sources in MIPSGAL.
The remaining sources have positions determined from 24 $\mu$ m point sources in MIPSGAL.
It is generally found that 6.7 GHz methanol masers coincide with 24 m point sources with good accuracy (?.. Pandian et al..
It is generally found that 6.7 GHz methanol masers coincide with 24 $\mu$ m point sources with good accuracy \citealt{xu09}, Pandian et al.,
in preparation).
in preparation).
For instance. the mean deviation between accurate maser positions and their corresponding 24 jm MIPSGAL point sources for the 14 sources above is 1.1.
For instance, the mean deviation between accurate maser positions and their corresponding 24 $\mu$ m MIPSGAL point sources for the 14 sources above is $''$.
Consequently. for sources without good interferometer positions. we can use MIPSGAL point sources to determine positions accurate to ~1”.
Consequently, for sources without good interferometer positions, we can use MIPSGAL point sources to determine positions accurate to $\sim 1''$.
As indicated in ?.. the AMGPS positions have a root mean square positional accuracy to 7”. which translates to an uncertainty of 18” at the confidence level.
As indicated in \citet{pand07b}, the AMGPS positions have a root mean square positional accuracy to $7''$, which translates to an uncertainty of $''$ at the confidence level.
Hence. we looked for MIPSGAL point sources within 18” of the 6 methanol masers without good interferometer positions.
Hence, we looked for MIPSGAL point sources within $''$ of the 6 methanol masers without good interferometer positions.
In most cases. there was a single point source within the search radius that was defined as the counterpart for the maser.
In most cases, there was a single point source within the search radius that was defined as the counterpart for the maser.
In cases where there were more than one 24 jm source. we looked at mid-infrared colors using the GLIMPSE point source catalog. and selected counterparts that had colors consistent with those of young stellar objects.
In cases where there were more than one 24 $\mu$ m source, we looked at mid-infrared colors using the GLIMPSE point source catalog, and selected counterparts that had colors consistent with those of young stellar objects.
The positions. peak flux densities. distances and lummosities of the masers in our sample are shown in Table I.
The positions, peak flux densities, distances and luminosities of the masers in our sample are shown in Table 1.
The 3.6 em observations were carried out on December 8.) 2007 with the VLA in the B-configuration.
The 3.6 cm observations were carried out on December 8, 2007 with the VLA in the B-configuration.
The correlator was set up in the standard continuum mode. and provided a total bandwidth of 100 MHz per polarization.
The correlator was set up in the standard continuum mode, and provided a total bandwidth of 100 MHz per polarization.
The integratior time was 2.5 minutes per source.
The integration time was 2.5 minutes per source.
The source 3C286 was used as à primary and flux calibrator while 318564061. served as the phase calibrator.
The source 3C286 was used as a primary and flux calibrator while J1856+061 served as the phase calibrator.
The data were reduced using standarc procedures in the Astronomical Image Processing System (AIPS) of NRAO.
The data were reduced using standard procedures in the Astronomical Image Processing System (AIPS) of NRAO.
The full width at half-maximum (FWHM) of the synthesized beam was 0.9"x0.8”. and the Lc noise i the images range from 0.13 to 0.24 mJy/beam.
The full width at half-maximum (FWHM) of the synthesized beam was $0.9'' \times 0.8''$, and the $1\sigma$ noise in the images range from 0.13 to 0.24 mJy/beam.
The 1.3 em observations were carried out with the VLA or March 22. 2008 in the C-configuration.
The 1.3 cm observations were carried out with the VLA on March 22, 2008 in the C-configuration.
The source 3C286 was used as a primary flux calibrator while J1851+005 served as the phase calibrator.
The source 3C286 was used as a primary flux calibrator while J1851+005 served as the phase calibrator.
The observations were carried out using fast switching for good phase stability.
The observations were carried out using fast switching for good phase stability.
Each fast switching cycle comprised of 140 s on the target followed by 40 s on the phase calibrator.
Each fast switching cycle comprised of 140 s on the target followed by 40 s on the phase calibrator.
The total on-source integration time was 6 to 13 minutes per source.
The total on-source integration time was 6 to 13 minutes per source.
The data were reduced using AIPS.
The data were reduced using AIPS.
The lov noise in the maps was around 0.12 mJy/beam. the synthesized beam being 1.0"x0.9".
The $1\sigma$ noise in the maps was around 0.12 mJy/beam, the synthesized beam being $1.0'' \times 0.9''$.
A subset of the sources in the sample (see Table 2)) were observed with the VLA at 6.9 mm in the D-configuration on September 2. 2008.
A subset of the sources in the sample (see Table \ref{table2}) ) were observed with the VLA at 6.9 mm in the D-configuration on September 2, 2008.
As with the 1.3 em observations. 3C286 was used as a primary flux calibrator while J1851+005 served as the phase calibrator.
As with the 1.3 cm observations, 3C286 was used as a primary flux calibrator while J1851+005 served as the phase calibrator.
Fast switching was used. each cycle comprising of 240 s on the target followed by 60 s on the phase calibrator. with the total on-source integration times ranging from 4 to 31 minutes depending on the strength of the source.
Fast switching was used, each cycle comprising of 240 s on the target followed by 60 s on the phase calibrator, with the total on-source integration times ranging from 4 to 31 minutes depending on the strength of the source.
The data were reduced using AIPS.
The data were reduced using AIPS.
The root mean square noise level in the maps ranged from 0.4 to 1.7 mJy/beam and the FWHM synthesized beam sizes ranged from 2.7"x2.2" to 3.8"x2.3".
The root mean square noise level in the maps ranged from 0.4 to 1.7 mJy/beam and the FWHM synthesized beam sizes ranged from $2.7'' \times 2.2''$ to $3.8'' \times 2.3''$.
The 1.2 mm observations were carried out between January and March 2008 as part of "pooled observations" using the 117 element MAMBO-II bolometer camera (?) at the IRAM 30-meter telescope.
The 1.2 mm observations were carried out between January and March 2008 as part of “pooled observations” using the 117 element MAMBO-II bolometer camera \citep{krey98} at the IRAM 30-meter telescope.
Each source was mapped using the standard on-the-fly technique such that à. -2' region was uniformly sampled (although the actual extents of the maps are considerably larger).
Each source was mapped using the standard on-the-fly technique such that a $\sim 2'$ region was uniformly sampled (although the actual extents of the maps are considerably larger).
The timespent per map was around 12 minutes with pointing. focus and calibration scans being done at regular intervals.
The timespent per map was around 12 minutes with pointing, focus and calibration scans being done at regular intervals.
The data were reduced using the "mapCSF" procedure
The data were reduced using the “mapCSF” procedure
The discovery of numerous extrasolar planetary systems in the solar neighbourhood (Alavor Queloz 1995. Alarey Autler 1996: Marcy 1999) has revolutionisecl our ideas of the planet. formation process ancl how it can vary. from system, to system.
The discovery of numerous extrasolar planetary systems in the solar neighbourhood (Mayor Queloz 1995, Marcy Butler 1996; Marcy 1999) has revolutionised our ideas of the planet formation process and how it can vary from system to system.
Specifically. the fact that most of the systems found contain relatively massive planets at small separations. in contrast to our solar system. has engencdered significant research into possible orbital migration (e.g. Lin. Bodenheimer and Richarcson 1996).
Specifically, the fact that most of the systems found contain relatively massive planets at small separations, in contrast to our solar system, has engendered significant research into possible orbital migration (e.g. Lin, Bodenheimer and Richardson 1996).
More. recently. )0 discovery that there appears to. be no. such close systems in the elobular cluster 47 Tuc implies à significant dillerence in planetary formation whieh could be due to jo stellar environment (Brown et al 2000. Crillilancl οἱ al JOO).
More recently, the discovery that there appears to be no such close systems in the globular cluster 47 Tuc implies a significant difference in planetary formation which could be due to the stellar environment (Brown et al 2000, Gilliland et al 2000).
Indeed. it is possible that stellar interactions in the early stages of the globular cluster were able to clisrupt 10 circumstellar clises before any planets were able to form (Bonnell 2000) or that ie increased. radiation rom the expected number of O stars was sullicient to remove these circumstellar cdises before any planets could orm (Armitage 2000).
Indeed, it is possible that stellar interactions in the early stages of the globular cluster were able to disrupt the circumstellar discs before any planets were able to form (Bonnell 2000) or that the increased radiation from the expected number of O stars was sufficient to remove these circumstellar discs before any planets could form (Armitage 2000).
Encounters with passing stars in a dense stellar environment can lead. to disruption of the planetary system. ancl thus the ejection of the planets (see e.g. Sigurdsson. 1992).
Encounters with passing stars in a dense stellar environment can lead to disruption of the planetary system and thus the ejection of the planets (see e.g. Sigurdsson, 1992).
Phis could lead to a population of rec-Lloating planets in the cluster.
This could lead to a population of free-floating planets in the cluster.
Itecently there has been a reported. detection of a population of substellar objects in & Orionis (Zapatero-Osorio ct al. 2000) that could. be due to stellar encounters.
Recently there has been a reported detection of a population of substellar objects in $\sigma$ Orionis (Zapatero-Osorio et al, 2000) that could be due to stellar encounters.
In this letter. we investigate the formation of a population of free-Ioating planets in. various cluster environments.
In this letter, we investigate the formation of a population of free-floating planets in various cluster environments.
We pav. particular attention to the velocity distribution of this population. and the question of whether the bulk of the liberated objects could be retained in their natal environment once they are ejected (rom their parent svsten.
We pay particular attention to the velocity distribution of this population, and the question of whether the bulk of the liberated objects could be retained in their natal environment once they are ejected from their parent system.
In the next section. we discuss the properties of the initial planet population and of the various clusters.
In the next section we discuss the properties of the initial planet population and of the various clusters.
We then brielly summarise the issue of interaction cross sections. including discussion of the cdilferent. possibilities. following an interaction.
We then briefly summarise the issue of interaction cross sections, including discussion of the different possibilities following an interaction.
We then describe the simulations of the various encounters and derive velocity dispersions and other properties for both the free floating and bound planet populations.
We then describe the simulations of the various encounters and derive velocity dispersions and other properties for both the free floating and bound planet populations.
Observations indicate that YSO clisks are typically LOOAU in radius (MeCaughrean O'Dell. 1996).
Observations indicate that YSO disks are typically 100AU in radius (McCaughrean O'Dell, 1996).
Although it isn't clear to what radius in the disk planets generally form. we can estimate based on our own solar svstem that planet ancl planetessimal formation has occured at radii out to 40-50
Although it isn't clear to what radius in the disk planets generally form, we can estimate based on our own solar system that planet and planetessimal formation has occured at radii out to 40-50
spectra were obtained by scaling appropriately for the size of the two regions.
spectra were obtained by scaling appropriately for the size of the two regions.
Finally. the event arrival times were corrected to the solar svstem barycenter and to the center of the binary system. using the ephemeris of Still et al. (
Finally, the event arrival times were corrected to the solar system barycenter and to the center of the binary system, using the ephemeris of Still et al. (
2001).
2001).
The beat period. of Her. N-1 is known to vary by up to a lew days (Davkal ct al.
The beat period of Her X-1 is known to vary by up to a few days (Baykal et al.
1993).
1993).
In order to. determine how the epochs of the observations relate to the precession evele. we extracted the (2-10) keV quick-Iook light curve from the MEE web site and we determined the 35 cay phase by computing the time dillerence. between cach observation and the start of the previous rise to maximum.
In order to determine how the epochs of the observations relate to the precession cycle, we extracted the (2-10) keV quick-look light curve from the MIT web site and we determined the 35 day phase by computing the time difference between each observation and the start of the previous rise to maximum.