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Likewise. I perform a similar calculation or the host stars of transiting exoplanets. this time assuniüne aliguinent between the host stars spin and he orbit of its exoplanet. | Likewise, I perform a similar calculation for the host stars of transiting exoplanets, this time assuming alignment between the host star's spin and the orbit of its exoplanet. |
I identify those transiting exoplauct svstenis in which the disagreement between he esn/ predicted frou the simple enipirical model nuder the asstuption of spin-orbit aliguiament and the observed ¢sin/ is larger than the equivalent disagreement for any of the 866 stars in the SPOCS control sample. | I identify those transiting exoplanet systems in which the disagreement between the $v\sin{i}$ predicted from the simple empirical model under the assumption of spin-orbit alignment and the observed $v\sin{i}$ is larger than the equivalent disagreement for any of the 866 stars in the SPOCS control sample. |
Tn particular. those svstenis that lave anomalously stnall esu nieasureineuts relative to the measurement errors aud the width of simulated οsins distribution are possibly spin-orbit miusaliened alone the line of sight. | In particular, those systems that have anomalously small $v\sin{i}$ measurements relative to the measurement errors and the width of simulated $v\sin{i}$ distribution are possibly spin-orbit misaligned along the line of sight. |
I describe the details of 11v. calculation iu 822. | I describe the details of my calculation in 2. |
I discuss niv results in the context of close-in plauet formation m $33. | I discuss my results in the context of close-in planet formation in 3. |
I stmunarize mv findiues iu &11I. | I summarize my findings in 4. |
I combine the well defined. mass-rotation period relation iu the ITvades aud Pracsepe (e.g.bwin&Bou-vier2009) with the relation P.X£7 (eg.Weber&Davis1967:Moestel1968:Kawaler1988) to compute the expected rotation periods of Suu-like stars as a function of mass aud age. | I combine the well defined mass-rotation period relation in the Hyades and Praesepe \citep[e.g.][]{irw09} with the relation $P_{\ast} \propto
t^{1/2}$ \citep[e.g.][]{web67,mes68,kaw88} to compute the expected rotation periods of Sun-like stars as a function of mass and age. |
I use a Moute Carlo simulation to transform that functiou mto a es1u/,;,, distribution bv mnargnmalizius over the uncertain distributions of stellar mass AZ. radius AR. age ne and the applicable oouclinatiou distribution for the sample. | I use a Monte Carlo simulation to transform that function into a $v\sin{i}_{sim}$ distribution by marginalizing over the uncertain distributions of stellar mass $M_{\ast}$, radius $R_{\ast}$, age $\tau_{\ast}$, and the applicable inclination distribution for the sample. |
As a result. I ‘an compare the observed esimn/j44; lucasuremenuts with 1ο mnenni CSIs, estimate from the Monte Carlo simulation to determine the degree of agreement between 16 two. | As a result, I can compare the observed $v\sin{i}_{obs}$ measurements with the mean $\overline{v\sin{i}}_{sim}$ estimate from the Monte Carlo simulation to determine the degree of agreement between the two. |
In order to determine the degree ofdisaereemout wat can be attributed to random effects I compare re measured values of ¢sin/,p. of Sun-like stars in the uyolar Neighborhood from the SPOCS catalog with the uean CSἐκ Hom the Monte Carlo simulation οἼνοιι je stellar paraicters of those stars and the standard isotropic inclination distribution. | In order to determine the degree of disagreement that can be attributed to random effects, I compare the measured values of $v\sin{i}_{obs}$ of Sun-like stars in the Solar Neighborhood from the SPOCS catalog with the mean $\overline{v\sin{i}}_{sim}$ from the Monte Carlo simulation given the stellar parameters of those stars and the standard isotropic inclination distribution. |
Tn this way. I use the control sample of SPOCS stars to determine a threshold evel of cisaerecment that can be expected between he simple enipirical model and observation eiven the observational uncertainties aud the imperfections of the uodel. | In this way, I use the control sample of SPOCS stars to determine a threshold level of disagreement that can be expected between the simple empirical model and observation given the observational uncertainties and the imperfections of the model. |
I then do a similar calculation for the host stars of transiting exoplauets. this time assunüng aligumenut vetween the spin of the host star and the orbit of its exoplanet. | I then do a similar calculation for the host stars of transiting exoplanets, this time assuming alignment between the spin of the host star and the orbit of its exoplanet. |
I ideutify those exoplanet host stars iu which the predicted esinἐκ disagrees with the observed value οσαρε relative to the measurement errors bv an amount ereater then an equivaleut disaerecieut for any of the 866 stars in the control sample. | I identify those exoplanet host stars in which the predicted $\overline{v\sin{i}}_{sim}$ disagrees with the observed value $v\sin{i}_{obs}$ relative to the measurement errors by an amount greater then an equivalent disagreement for any of the 866 stars in the control sample. |
Systems with οσαἐν slower than οπιη bv au amount large relative to the quoted errors in οsin/,;5, aud the width of the simulated 0siu;,, distribution given the nucertaity in stellar mass. radius. and age are either | Systems with $v\sin{i}_{obs}$ slower than $\overline{v\sin{i}}_{sim}$ by an amount large relative to the quoted errors in $v\sin{i}_{obs}$ and the width of the simulated $v\sin{i}_{sim}$ distribution given the uncertainty in stellar mass, radius, and age are either |
Figure 6. shows the variation of £,. £, and s, with ;A. which corresponds to a change in the strength of the basic state magnetic field at the top of the laver (since plz=0) is kept constant). | Figure \ref{fig:emfsb0} shows the variation of $\barcalE_x$ , $\barcalE_y$ and $s_r$ with $A$, which corresponds to a change in the strength of the basic state magnetic field at the top of the layer (since $\rho (z=0)$ is kept constant). |
For low values of Dy. £, increases {his can be thought of as an increase in the field strength allowing the instability to become stronger and therefore generate more emf, | For low values of $B_0$, $\barcalE_x$ increases — this can be thought of as an increase in the field strength allowing the instability to become stronger and therefore generate more emf. |
As the strength of the field increases Burther. the growth rate continues to increase: (he influence of rotation is (hus diminished and. just as im Figure 5.. £, decreases in magnitude. | As the strength of the field increases further, the growth rate continues to increase; the influence of rotation is thus diminished and, just as in Figure \ref{fig:emfszeta}, $\barcalE_x$ decreases in magnitude. |
There is a sharp increase in £, over low Alfvénn speeds as the instability sets in. a slight decrease as the Alfvénn speed is further increased and then a sharp drop at high. Alfvénn speeds. | There is a sharp increase in $\barcalE_y$ over low Alfvénn speeds as the instability sets in, a slight decrease as the Alfvénn speed is further increased and then a sharp drop at high Alfvénn speeds. |
since lor many astrophlivsical dvnanmos it is believed that the toroidal field. generated by the differential rotation (the '«-ellect in the language of mean field electrodynanmics). dominates the poloidal component. (hen it makes sense initially to consider the nature of the emf that results from the instability of apurely toroidal field: this is the approach adopted in | Since for many astrophysical dynamos it is believed that the toroidal field, generated by the differential rotation (the $\omega$ -effect' in the language of mean field electrodynamics), dominates the poloidal component, then it makes sense initially to consider the nature of the emf that results from the instability of apurely toroidal field; this is the approach adopted in |
The successful latching of CoRoT produces data for transiting exoplauets with unprecedented accuracy compared to ground-based observations. providing a lüensure of anetarv ness and radius auc thus information on the mean density aud bulk composition of cxoplaucts. | The successful launching of CoRoT produces data for transiting exoplanets with unprecedented accuracy compared to ground-based observations, providing a measure of planetary mass and radius and thus information on the mean density and bulk composition of exoplanets. |
—Unfortunately, a remaining source of uncertainty on the planetary parameters (1nasxs and radius) is due to uucertainties ou the stellar parameters. | Unfortunately, a remaining source of uncertainty on the planetary parameters (mass and radius) is due to uncertainties on the stellar parameters. |
The precise deteruination of eravity by spectroscopy is difficult. and particularly for F stars. | The precise determination of gravity by spectroscopy is difficult, and particularly for F stars. |
The determinalon of he mass of the star is the maim wacertainty on the pauets size. | The determination of the mass of the star is the main uncertainty on the planet's size. |
This uncertainty on the mass is generaIvy abot on the radius aud on the deusivy). | This uncertainty on the mass is generally about on the radius and on the density). |
Tjs nia o0 of Muportance to disantauele a massive plane from a brown dwarf. as in the case of the recenlv. discovered “super-Jupiter” CoRoT-Exo-3b (Lecoute et al. | This may be of importance to disantangle a massive planet from a brown dwarf, as in the case of the recently discovered "super-Jupiter" CoRoT-Exo-3b (Leconte et al. |
2009). | 2009). |
It nay also prevent distinguishing au ocean planet from an carth-like plauet (Caasset et al. | It may also prevent distinguishing an ocean planet from an earth-like planet (Grasset et al. |
2009). | 2009). |
Our goal is to examine the accuracy ou he latter poriuueters namely mass. ΓΙ and radius. tha could be obtained using asteroscisimology with the expected accuracy on oscillation frequencies of CoRoT. We explore the space of stellar parameters. iu ternis of luass. effective temperature. huuinositv. metallicity arc πμπιο leneth paraneter aud we analyse the seusitivitv of predicted. spectiii of oscillation frequencies to these paraucters. | Our goal is to examine the accuracy on the latter parameters, namely mass, luminosity and radius, that could be obtained using asteroseismology with the expected accuracy on oscillation frequencies of CoRoT. We explore the space of stellar parameters, in terms of mass, effective temperature, luminosity, metallicity and mixing length parameter and we analyse the sensitivity of predicted spectrum of oscillation frequencies to these parameters. |
Adopting various evels of stellar LOISC ALE white noise. we analyse the frequency uncertainty for oscillation modes of a given amplitude and lifetime. | Adopting various levels of stellar noise and white noise, we analyse the frequency uncertainty for oscillation modes of a given amplitude and lifetime. |
This analysis is based on the performances of CoBoT in the Asteroscisinology and Planet Finder chanuels described in Auverene et al. ( | This analysis is based on the performances of CoRoT in the Asteroseismology and Planet Finder channels described in Auvergne et al. ( |
2009). | 2009). |
Among all the transiting systems discovered so far. the large differeuce between the fundamental paramcters of NO-3 eiveu by Johlus-IXrull et al. ( | Among all the transiting systems discovered so far, the large difference between the fundamental parameters of XO-3 given by Johns-Krull et al. ( |
2008) (AT=1. l= 0.08 AQ. : R22.13 + 0.21 R.) aud bx Winn et al. ( | 2008) (M=1.41 $\pm$ 0.08 $\msol$ ; R=2.13 $\pm$ 0.21 $\rsol$ ) and by Winn et al. ( |
2008) (AT=1.218 x 0.066 AL.: R=L377 x 0.083 Π. ) Was a strong motivation to exaiuine how accurately fundamental parameters of a star (nass. radius. luuinositv) can be coustrained by asteroseiuolosgv. | 2008) (M=1.213 $\pm$ 0.066 $\msol$; R=1.377 $\pm$ 0.083 $\rsol$ ) was a strong motivation to examine how accurately fundamental parameters of a star (mass, radius, luminosity) can be constrained by asteroseismology. |
Iu this preliminary study. we first analyse stars with characteristics close to the ones of NO-3 AL~ L1 M. I 1.5 Ro. Tet ~ 6100 IX). | In this preliminary study, we first analyse stars with characteristics close to the ones of XO-3 (M $\sim$ 1.4 $\msol$, R $\sim$ 1.5 $\rsol$, Teff $\sim$ 6400 K). |
We focus on modes with low order 7 (0<7x3) aud radial order s between 5 aud 20. | We focus on modes with low order $l$ $0 \leq l \leq 3$ ) and radial order $n$ between 5 and 20. |
The pulsation calculations are performed with a nonradial code originally developed bv Lee (1985) aud based ou a linear non-aciabaticstability analvsis. | The pulsation calculations are performed with a nonradial code originally developed by Lee (1985) and based on a linear non-adiabaticstability analysis. |
The equations are linearised around lvdrostatic equilibiun. and cigenfunctions are expressed with spherical harimonics Ji. | The equations are linearised around hydrostatic equilibrium, and eigenfunctions are expressed with spherical harmonics $Y_{lm}$. |
The eigeufrequeucies are defined by with o, the oscillation frequency aud 0; the cami rate Gf positive) or erowth rate Gf negative) (see details in Mulet-Marquis ot. al. | The eigenfrequencies are defined by _i, with $\sigma_r$ the oscillation frequency and $\sigma_i$ the damping rate (if positive) or growth rate (if negative) (see details in Mulet-Marquis et al. |
2007. and references therein). | 2007, and references therein). |
The system of eqiations is solved with a Ileuvev-tvpoe relaxation method. | The system of equations is solved with a Henyey-type relaxation method. |
Stellar structure models are calculated using the Livermore opacities (Iglesias. Rogers. 1996). | Stellar structure models are calculated using the Livermore opacities (Iglesias, Rogers, 1996). |
The uunber of eridpoiits of the equilixiu models (4000 eridpoiuts) is high enough for the uncertaiutv on 1C frequency determination by the linear code to be less than the accuracy expeced from CoBoTor a pulsation niocle with an iufinite lifeiue aud observec for 150 davs (60d plz). | The number of gridpoints of the equilibrium models (4000 gridpoints) is high enough for the uncertainty on the frequency determination by the linear code to be less than the accuracy expected from CoRoTfor a pulsation mode with an infinite lifetime and observed for 150 days 0.1 $\mu$ Hz). |
this band overlaps with rotation-vibration transitions of CO; present in the purge gas in our spectrometer. | this band overlaps with rotation-vibration transitions of $_2$ present in the purge gas in our spectrometer. |
The use of 190150, isotopic species does not improve this situation. | The use of $^{13}$ $^{18}$ $_2$ isotopic species does not improve this situation. |
Therefore we have focused on the strongest HCOOH band, the C=O stretching mode, at 1710 cm~!, to monitor formic acid formation. | Therefore we have focused on the strongest HCOOH band, the C=O stretching mode, at 1710 $^{-1}$, to monitor formic acid formation. |
The observed difference RAIR spectrum of 190150. bombarded by H-atoms has been compared to that of 120150) in Fig. | The observed difference RAIR spectrum of $^{13}$ $^{18}$ $_2$ bombarded by H-atoms has been compared to that of $^{12}$ $^{18}$ $_2$ in Fig. |
2aa. Both spectra show very weak features at 1730 cm™! and 1500 cm!, but none at positions typical for HCOOH. | \ref{spec}a a. Both spectra show very weak features at 1730 $^{-1}$ and 1500 $^{-1}$, but none at positions typical for HCOOH. |
The detected bands, however, occur at exactly the same positions as previously seen for HxCO when formed upon CO hydrogenation (e.g.,Watanabeetal.,2004). | The detected bands, however, occur at exactly the same positions as previously seen for $_2$ CO when formed upon CO hydrogenation \citep[e.g.,][]{watanabe2004}. |
. Furthermore, the features do not shift when a different isotopic species is used, which is expected when the formation involves CO; ice. | Furthermore, the features do not shift when a different isotopic species is used, which is expected when the formation involves $_2$ ice. |
In the TPD spectra of CO; ices bombarded with H-atoms for 3 hrs (not shown) mass 29 and 30 amu desorb in two steps around 100 K and 140 K, and the formation of other species is not detected through TPD. | In the TPD spectra of $_2$ ices bombarded with H-atoms for 3 hrs (not shown) mass 29 and 30 amu desorb in two steps around 100 K and 140 K, and the formation of other species is not detected through TPD. |
These temperatures and masses are identical to what was observed by Fuchsetal.(2007) for H;CO desorption from CO ices bombarded by H-atoms. | These temperatures and masses are identical to what was observed by \citet{fuchs2007} for $_2$ CO desorption from CO ices bombarded by H-atoms. |
During H- bombardment an increase in the mass 28 amu signal is observed, which is likely due to degassing of both CO and Ν2 from the metal parts of our experiment. | During H-atom bombardment an increase in the mass 28 amu signal is observed, which is likely due to degassing of both CO and $_2$ from the metal parts of our experiment. |
It is therefore plausible that the measured low level of H5CO formation observed in the experiment originates from hydrogenation of background gaseous CO and is not related to the CO» ice. | It is therefore plausible that the measured low level of $_2$ CO formation observed in the experiment originates from hydrogenation of background gaseous CO and is not related to the $_2$ ice. |
To test whether the presence of H5O affects the reactivity of CO; ice upon H-atom bombardment as has been observed for CO in CO:H50 mixtures (Watanabeetal.,2004; 2007),, mixtures of to CO2:H20 have been investigated (see Fig. | To test whether the presence of $_2$ O affects the reactivity of $_2$ ice upon H-atom bombardment as has been observed for CO in $_2$ O mixtures \citep{watanabe2004,fuchs2007}, , mixtures of to $_2$ $_2$ O have been investigated (see Fig. |
2bb). | \ref{spec}b b). |
Like in the experiments with pure CO», weak RAIR features of similar intensity are observed at 1730 and 1500 cm7!. | Like in the experiments with pure $_2$, weak RAIR features of similar intensity are observed at 1730 and 1500 $^{-1}$. |
Again we assign these features tothe C=O stretching mode and the bending mode of H5CO. | Again we assign these features tothe C=O stretching mode and the C-H bending mode of $_2$ CO. |
Thus, as for the pure ices, a small amount of background CO accretes and forms H2CO. | Thus, as for the pure ices, a small amount of background CO accretes and forms $_2$ CO. |
Within the sensitivity of our experiment, we conclude that CO» does not react with H-atoms even if mixed with H5O. Finally, mixtures of '*C!°O and 1390/50, are studied to determine which species is more likely to react upon H-atom bombardment: CO or CO». | Within the sensitivity of our experiment, we conclude that $_2$ does not react with H-atoms even if mixed with $_2$ O. Finally, mixtures of $^{12}$ $^{16}$ O and $^{13}$ $^{18}$ $_2$ are studied to determine which species is more likely to react upon H-atom bombardment: CO or $_2$. |
Since CO hydrogenation reactions were previously reported in the literature (Watanabeetal.,2004;Fuchsetal., 2007),, the answer to this question must be CO. | Since CO hydrogenation reactions were previously reported in the literature \citep{watanabe2004,fuchs2007}, the answer to this question must be CO. |
In Fig. | In Fig. |
2cc and d, the resulting difference spectra for CO»5:CO mixtures are shown for different ice thicknesses and with mixture concentrations of CO;:CO and 2:CO, respectively. | \ref{spec}c c and d, the resulting difference spectra for $_2$ :CO mixtures are shown for different ice thicknesses and with mixture concentrations of $_2$ :CO and $_2$ :CO, respectively. |
Similar to H- bombardment of pure CO; ices and CO?:H50 mixtures, no evidence for HCOOH formation is observed in CO2:CO mixtures. | Similar to H-atom bombardment of pure $_2$ ices and $_2$ $_2$ O mixtures, no evidence for HCOOH formation is observed in $_2$ :CO mixtures. |
In contrast, CO does react with H-atoms to form HCO and CH30OH as is evidenced by the presence of strong H5CO absorption features at 1730 and 1500 cm™! and CH3OH at 1030 cm™!. | In contrast, CO does react with H-atoms to form $_2$ CO and $_3$ OH as is evidenced by the presence of strong $_2$ CO absorption features at 1730 and 1500 $^{-1}$ and $_3$ OH at 1030 $^{-1}$. |
This is consistent with H-atom bombardment experiments for pure CO and CO:H2O mixtures by Watanabeetal.(2004) and Fuchsetal.(2007). | This is consistent with H-atom bombardment experiments for pure CO and $_2$ O mixtures by \citet{watanabe2004} and \citet{fuchs2007}. |
. The complementary TPD data show the same picture of no HCOOH formation and clear HCO and CH30H formation from CO. | The complementary TPD data show the same picture of no HCOOH formation and clear $_2$CO and $_3$ OH formation from CO. |
Other products than the precursor and product species H2CO and CH30OH are not observed. | Other products than the precursor and product species $_2$ CO and $_3$ OH are not observed. |
No hydrogenation products of CO», specifically HCOOH, are observed within the experimental sensitivity. | No hydrogenation products of $_2$, specifically HCOOH, are observed within the experimental sensitivity. |
The limit on the formation reaction rate for HCOOH from CO); is <7.0x107" cm? s! based on the limit on the column density for HCOOH after 1 min of H-atom bombardment for all ice morphologies (see 3.2 for the derivation and Table3 for the individual values for each experiment). | The limit on the formation reaction rate for HCOOH from $_2$ is $\leq$ $\times$ $^{-17}$ $^{2}$ $^{-1}$ based on the limit on the column density for HCOOH after 1 min of H-atom bombardment for all ice morphologies (see \ref{rate_sec} for the derivation and Table \ref{dest} for the individual values for each experiment). |
In Figure 3 the absorbance divided by the initial absorbance at {—0, A/Ao, is shown for the and CO5:CO ice mixtures. | In Figure \ref{cofit} the absorbance divided by the initial absorbance at $t=$0, $A/A_0$ , is shown for the and $_2$ :CO ice mixtures. |
The data are fitted as described in 3.2 and the fits are indicated in Fig. | The data are fitted as described in \ref{rate_sec} and the fits are indicated in Fig. |
3 with dotted lines. | \ref{cofit} with dotted lines. |
The resulting values for ao and fo as well as the ko are given in Table 4.. | The resulting values for $\alpha_0$ and $\beta_0$ as well as the $k_0$ are given in Table \ref{corate}. . |
Since the H2CO band strength could not be determined accurately in these experiments only o and fo are | Since the $_2$ CO band strength could not be determined accurately in these experiments only $\alpha_0$ and $\beta_0$ are |
spectrum (normalized to unitv power) is shown in figure 1. | spectrum (normalized to unity power) is shown in figure 1. |
There are three peaks present. two of which are aliases of each other al 0.124706(1) and 0.183825(1) Hz. on opposite sides of the Nv«quist frequency. | There are three peaks present, two of which are aliases of each other at 0.124706(1) and 0.183828(1) Hz, on opposite sides of the Nyquist frequency. |
These peaks have highly significant. power values of 38.4. | These peaks have highly significant power values of 38.4. |
The chance probability of such a power is 3x10.I. | The chance probability of such a power is $\sim 3\times10^{-17}$. |
The peak at 0.308 Iz is due to the digitization limit (3.241 s) of the ACIS-I and has a power value approximately equal to the nunmber of photons. | The peak at 0.308 Hz is due to the digitization limit (3.241 s) of the ACIS-I and has a power value approximately equal to the number of photons. |
The spectrum of the source is very soft. | The spectrum of the source is very soft. |
An acceptable spectral fit to the data is given by a black-body model with an absorption column density. 1440.2x107! II-atoms/eni? and a value of kT = 0.4140.01 keV. The fit to the data is shown in Figure 2. | An acceptable spectral fit to the data is given by a black-body model with an absorption column density $\pm 0.2\times 10^{21}$ $^2$ and a value of kT = $\pm$ 0.01 keV. The fit to the data is shown in Figure 2. |
A search of imaging X-rav satellite archives for data for this source has produced many observations dating from 1979. | A search of imaging X-ray satellite archives for data for this source has produced many observations dating from 1979. |
These observations are listed in Table 1. | These observations are listed in Table 1. |
The most extensive source coverage is from the ROSAT satellite. | The most extensive source coverage is from the ROSAT satellite. |
For ROSAT we have restricted the observations shown in the Table bv excluding all observations which contain fewer than ~ 50 source photons. | For ROSAT we have restricted the observations shown in the Table by excluding all observations which contain fewer than $\sim$ 50 source photons. |
This means that for the ROSAT ILRI exposures less than 7.5 ks are excluded. | This means that for the ROSAT HRI exposures less than 7.5 ks are excluded. |
For the PSPC exposures less than 1.5 ks are excluded. as are observations in which the source is more than 30J) are minutes from the center of the field of view. | For the PSPC exposures less than 1.5 ks are excluded as are observations in which the source is more than 30 arc minutes from the center of the field of view. |
This latter restriction is designed to avoid large vignetting corrections to the counting rate. | This latter restriction is designed to avoid large vignetting corrections to the counting rate. |
Results based on observations of only one narrow field are subject to biases due to large seale-structure in the direction of that [lield. | Results based on observations of only one narrow field are subject to biases due to large scale-structure in the direction of that field. |
For this reason we have compiled additional eroup/cluster detections from the deep ROSAT PSPC surveys of Hasinger (1998a). Zamorani (1999) and. Bower (1996). | For this reason we have compiled additional group/cluster detections from the deep $ROSAT$ PSPC surveys of Hasinger (1998a), Zamorani (1999) and Bower (1996). |
For consistency. we only use deep surveys based on PSPC data. | For consistency, we only use deep surveys based on PSPC data. |
In all these surveys a considerable effort has been made by the authors to identity all of the N-rav. sources. | In all these surveys a considerable effort has been made by the authors to identify all of the X-ray sources. |
The Lockman hole survey of Hasinger (19982). Schmidt (1998). Lehmann (2000) was obtained using 207 ksec of PSPC data combined with LIBI data. | The Lockman hole survey of Hasinger (1998a), Schmidt (1998), Lehmann (2000) was obtained using 207 ksec of PSPC data combined with HRI data. |
A source detection algorithm enploving a maximum likelihood fit to the PSE was used. | A source detection algorithm employing a maximum likelihood fit to the PSF was used. |
All 50 X-ray sources with [Iuxes D5.5x]10 [(0.5-2 keV) have been identilied. resulting in 3 eroup/cluster identifications. | All 50 X-ray sources with fluxes $>$ $^{-15}$ (0.5-2 keV) have been identified, resulting in 3 group/cluster identifications. |
ALL 3 (numbers 41. 58. 67) are extended: N-ray sources with multiple galaxy redshifts per system. | All 3 (numbers 41, 58, 67) are extended X-ray sources with multiple galaxy redshifts per system. |
We do not include the high redshift clusters reported in Hasinger (1998b). because these were outside the complete survey area. or below the Dux limit. of the Hasinger (1998a) PSPC survey. | We do not include the high redshift clusters reported in Hasinger (1998b) because these were outside the complete survey area, or below the flux limit, of the Hasinger (1998a) PSPC survey. |
Vhe survey of Zamorani (1999). used. a PSPC exposure of 56 ksee on the Marano field. | The survey of Zamorani (1999) used a PSPC exposure of 56 ksec on the Marano field. |
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