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We found for 52 stars the mean abundance logz(Mg)=7.67+ 0.21. | We found for 52 stars the mean abundance $\log \varepsilon({\rm Mg}) =
7.67\pm0.21$ . |
The effect of uncertainties in the microturbulent »arameter oon log(Meg) is important. especially for coolest programme stars. | The effect of uncertainties in the microturbulent parameter on $\log \varepsilon({\rm Mg})$ is important, especially for coolest programme stars. |
For 16 such stars the vvalues were derived from lines. but not from aand llines as for other stars. | For 16 such stars the values were derived from lines, but not from and lines as for other stars. |
Being close to zero these D) values are less accurate than the ID) values. | Being close to zero these ) values are less accurate than the ) values. |
When excluding these [6 stars. we obtained the mean abundance logz(Mg)=75S+0.15 for the remaining 36 stars. | When excluding these 16 stars, we obtained the mean abundance $\log
\varepsilon({\rm Mg}) = 7.59\pm0.15$ for the remaining 36 stars. |
This abundance is precisely confirmed from an analysis of the weak TISTT line for several hot B stars. | This abundance is precisely confirmed from an analysis of the weak 7877 line for several hot B stars. |
This is our recommended Meg abundance for the B-type MS stars in the solar neighbourhood (with d< 800 pe). | This is our recommended Mg abundance for the B-type MS stars in the solar neighbourhood (with $d <$ 800 pc). |
Comparing the latter value with the solar magnesium abundance loge.(Mg)=7.55+0.02. one may conclude that he metallicity of the stars is very close to the solar one. | Comparing the latter value with the solar magnesium abundance $\log \varepsilon_{\sun}({\rm Mg}) =
7.55\pm0.02$, one may conclude that the metallicity of the stars is very close to the solar one. |
Our mean Tg abundance in B stars. as well as the position of the maximum in Fig.6. log(Mg).=7.64. is also very close to the proto-Sun magnesium abundance log2,,CMg)=7.62+0.02. | Our mean Mg abundance in B stars, as well as the position of the maximum in Fig.6, $\log
\varepsilon({\rm Mg}) = 7.64$, is also very close to the proto-Sun magnesium abundance $\log \varepsilon_{ps}({\rm Mg})=7.62\pm0.02$. |
We discussed he preceding determinations of the Mg abundance in B stars by Daflon et al. ( | We discussed the preceding determinations of the Mg abundance in B stars by Daflon et al. ( |
2003). | 2003). |
Their logz(Mg) values are somewhat lower han ours. | Their $\log \varepsilon({\rm
Mg})$ values are somewhat lower than ours. |
We showed that this difference in log:(Mg) is explained by differences inYj. | We showed that this difference in $\log \varepsilon({\rm
Mg})$ is explained by differences in. |
. "Thus. our results show that the Sun is not measurably enriched in metals as compared with the neighbouring voung stars. | Thus, our results show that the Sun is not measurably enriched in metals as compared with the neighbouring young stars. |
Two of us. LSL and SIR. are grateful to the staff of the Astronomy Department and McDonald Observatory of the University of Texas for hospitality during the visit in spring 2004. | Two of us, LSL and SIR, are grateful to the staff of the Astronomy Department and McDonald Observatory of the University of Texas for hospitality during the visit in spring 2004. |
DLL acknowledges the support of the Robert A. Welch Foundation of Houston. Texas. | DLL acknowledges the support of the Robert A. Welch Foundation of Houston, Texas. |
hotspot components are described by homogeneous spheres with constant magnetic [field ancl constant properties of the relativistic electron populations. | hotspot components are described by homogeneous spheres with constant magnetic field and constant properties of the relativistic electron populations. |
The spectral energy clistributions of the emitting electrons are mocelled assuming the formalism described in Brunettietal.(2002). | The spectral energy distributions of the emitting electrons are modelled assuming the formalism described in \citet{gb02}. |
. According to this model a population of seed. electrons (with 55) ds accelerated. at the shock ancl is injected in the downstream region with a spectrum dN(z/dt x ~Fo forse,«5πρι 5. being the maximum energy of the electrons accelerated at the shock. | According to this model a population of seed electrons (with $\gamma \leq
\gamma_{*}$ ) is accelerated at the shock and is injected in the downstream region with a spectrum $\gamma$ )/dt $\propto$ $\gamma^{-p}$, for $\gamma_{*} < \gamma < \gamma_{c}$, $\gamma_{c}$ being the maximum energy of the electrons accelerated at the shock. |
Electrons accelerated at the shock are advected in the downstream region and age due to radiative losses. | Electrons accelerated at the shock are advected in the downstream region and age due to radiative losses. |
Based on Brunettietal.(2002).. the volume integrated spectrum of the electron.population in the downstream region of size LZe (E and c being the age and the advection velocity of the downstream region) is given by either a steep power-law AN(s) xsI"P for 54«*5 5,.. where 5, is the maximum cnerev of | Based on \citet{gb02}, the volume integrated spectrum of the electronpopulation in the downstream region of size $L \sim T v_{\rm adv}$ $T$ and $v_{\rm adv}$ being the age and the advection velocity of the downstream region) is given by either a steep power-law $N$ $\gamma$ ) $\propto \gamma^{-(p+1)}$ for $\gamma_{b} < \gamma < \gamma_{c}$ , where $\gamma_{b}$ is the maximum energy of |
“Hcluciall” | l” |
Schuster Nissen (1989) οΗ] calibrations.. Ht is clear rom Lig. | Schuster Nissen (1989) [Fe/H] calibrations.. It is clear from Fig. |
4 that accurate metal abundances can be derived rom ΠΟH3 photometry. using the combined calibrations as clescribed above. | 4 that accurate metal abundances can be derived from $uvby-H\beta$ photometry, using the combined calibrations as described above. |
For the empirical Fe/l1l] calibrations of Schuster dssen (1989) the estimated standard deviations of a single ohotometric determination of. Fe/1l] were £0.14 at ο] =0.5 dex. and x0.21 at. Fe/H] z1.5 dex. | For the empirical [Fe/H] calibrations of Schuster Nissen (1989) the estimated standard deviations of a single photometric determination of [Fe/H] were $\pm 0.14$ at [Fe/H] $\approx -0.5$ dex, and $\pm 0.21$ at [Fe/H] $\approx -1.5$ dex. |
Comparisons »w Leltzing et al. ( | Comparisons by Feltzing et al. ( |
(2001). of photometric— abunclances rom the Schuster Nissen (1989) calibration equations with abundances from two recent spectroscopic. studies Edvarelsson et al. ( | 2001) of photometric abundances from the Schuster Nissen (1989) calibration equations with abundances from two recent spectroscopic studies Edvardsson et al. ( |
1993) and Chen et al. ( | 1993) and Chen et al. ( |
2000) have shown hat these error estimates are overly conservative: they find a scatter of only £0.100.11 for their more metal-rich group. in very good agreement with the value of £0.12 found above rom a somewhat less homogeneous spectroscopic data set. | 2000) have shown that these error estimates are overly conservative; they find a scatter of only $\pm 0.10-0.11$ for their more metal-rich group, in very good agreement with the value of $\pm 0.12$ found above from a somewhat less homogeneous spectroscopic data set. |
In Fig. | In Fig. |
5 the distribution of metallicity. ML]. for our sample stars is presented. ancl also shown is a fit to the listogram using three Gaussians. | 5 the distribution of metallicity, [M/H], for our sample stars is presented, and also shown is a fit to the histogram using three Gaussians. |
The sample is mix of all he stellar. populations that are represented. in the solar neighbourhood. the thin and thick disks. and the halo. | The sample is mix of all the stellar populations that are represented in the solar neighbourhood, the thin and thick disks, and the halo. |
As is expected. in the solar neighbourhood. the majority of the stars belong to the disk with the thin disk dominating. à significant but smaller contribution from the thickdisk. and only a few halo stars (2S: see Fig. | As is expected, in the solar neighbourhood, the majority of the stars belong to the disk with the thin disk dominating, a significant but smaller contribution from the thickdisk, and only a few halo stars $\approx 8$; see Fig. |
6 below). | 6 below). |
In Fig. | In Fig. |
5 two main components can be seen. one with .0.95ΑΗ]0.5 dex (the disk contribution: thin plus thick) and the other with AZ/H]Z0.9 dex (mostly the halo stars). | 5 two main components can be seen, one with $-0.9 \la [M/H] \la +0.5$ dex (the disk contribution: thin plus thick) and the other with $[M/H] \la -0.9$ dex (mostly the halo stars). |
A rough Gaussian fit has been made to this metallicity histogram using the mathematical package7. | A rough Gaussian fit has been made to this metallicity histogram using the mathematical package. |
.. One disk-like component has <AH]»-—0.04 dex. a dispersion of 0.37. and is nearly symmetric. | One disk-like component has $<[M/H]> \sim -0.04$ dex, a dispersion of $0.37$, and is nearly symmetric. |
Our sample. contains an obvious contribution from the thick disk with a Gaussian waving a mean abundance of <ΑΗ>~O45 dex ancl a dispersion of 0.41. | Our sample contains an obvious contribution from the thick disk with a Gaussian having a mean abundance of $<[M/H]> \sim -0.45$ dex and a dispersion of $0.41$. |
There are eight halo stars with «ΑΗ]>L0.9 dex (see also Fig. | There are eight halo stars with $<[M/H]> \la -0.9$ dex (see also Fig. |
6). | 6). |
Nevertheless. hin and thick-cdisk stars are not clearly distinguishable from. Fig. | Nevertheless, thin and thick-disk stars are not clearly distinguishable from Fig. |
5. | 5. |
There exists no clean. straightforward: procedure or separating the stars of these different disk populations. which overlap considerably in the metallicity cistribution. | There exists no clean, straightforward procedure for separating the stars of these different disk populations, which overlap considerably in the metallicity distribution. |
]xinematies are usually invoked to improve the thin/thick disk separation. and more recently a/ Fe] abundances (sec for example. Densbv. Feltzing Lundstrómnm: 22003). | Kinematics are usually invoked to improve the thin/thick disk separation, and more recently $\alpha/Fe$ ] abundances (see for example, Bensby, Feltzing Lundströmm 2003). |
In Fig. | In Fig. |
6 the Via. ML] ciagram for our sample stars is plotted. | 6 the $V_{\rm rot}$, [M/H] diagram for our sample stars is plotted. |
“Pwo components are clearly noted: the halo component. VuaS 100 km and η]0.9 dex: the other disk-like. centered at Via~ 215 km s+ and ΔΕΗ] ~0.0 dex. | Two components are clearly noted: the halo component, $V_{\rm rot} \la $ 100 km $^{-1}$ and $[M/H] \la -0.9$ dex; the other disk-like, centered at $V_{\rm rot} \sim $ 215 km $^{-1}$ and [M/H] $\sim 0.0$ dex. |
The rotation velocity of the LSIt about the Galactic center is taken here to be 220 km s+ oso that Von=7| 220 kms + is the rest-frame rotation velocity of a given star (Ixerr LyndenBell 1986). | The rotation velocity of the LSR about the Galactic center is taken here to be 220 km $^{-1}$ so that $V_{\rm rot} = V' + $ 220 km $^{-1}$ is the rest-frame rotation velocity of a given star (Kerr LyndenBell 1986). |
More modern values for the circular speed of the Milky Way at the solar circle are somewhat larger. such as (3,=234E13 km s1 given by Pukugita Peebles (2004). | More modern values for the circular speed of the Milky Way at the solar circle are somewhat larger, such as $\Omega_{\rm 0} = 234 \pm 13$ km $^{-1}$ given by Fukugita Peebles (2004). |
Llere. V is the velocity with respect to the LSk. | Here, $V'$ is the velocity with respect to the LSR. |
It can be seen from Fig. | It can be seen from Fig. |
6 that the majority of the stars have Via~ 200 kms+ ane AMI]0.5 dex. | 6 that the majority of the stars have $V_{\rm rot} \sim $ 200 km $^{-1}$ and $[M/H] \geq -0.5$ dex. |
These stars are thin-disk stars. | These stars are thin-disk stars. |
Even though there is scarcity of stars with MI]<—0.5 dex and Vi,« 150 kms. JF. Via increases approximately linearly with M/LE] between LO<MH]—0.5 dex. | Even though there is scarcity of stars with $[M/H] < -0.5$ dex and $V_{\rm rot}<$ 150 km $^{-1}$, $V_{\rm rot}$ increases approximately linearly with [M/H] between $-1.0<[M/H]<-0.5$ dex. |
For the ΑΗ]<LO dex. it appears that there is no correlation between ML] and Via. but there are too few stars for any firm conclusion. | For the $[M/H] < -1.0$ dex, it appears that there is no correlation between [M/H] and $V_{\rm rot}$, but there are too few stars for any firm conclusion. |
This Via. MM] diagram can be used to separate out the different stellar populations. as discussed in Nissen Schuster (1991) and in Schuster et al. ( | This $V_{\rm rot}$ , [M/H] diagram can be used to separate out the different stellar populations, as discussed in Nissen Schuster (1991) and in Schuster et al. ( |
(1993). | 1993). |
In the former reference. a diagonal eut connecting (MI]. Va)=(0.8.0 | In the former reference, a diagonal cut connecting ([M/H], $V_{\rm rot}) = (-0.3,0$ |
needed to proceed with the de-correlation process. | needed to proceed with the de-correlation process. |
This is very well possible in the case of dedicated iustruinents such as EKepler. as these have been specifically designed with such high precision aud stability measurements in πιά (Doruckietal.1996:Jenkinsetal. 2010). | This is very well possible in the case of dedicated instruments such as Kepler, as these have been specifically designed with such high precision and stability measurements in mind \citep{borucki96,jenkins10}. |
. For iustraments that do not feature the calibration plau required to further de-correlate with instrumneut state paraneters. the solution is far less obvious. | For instruments that do not feature the calibration plan required to further de-correlate with instrument state parameters, the solution is far less obvious. |
We furthermore explored the de-correlation of eclipse signals observed. consecutively rather than in parallel. | We furthermore explored the de-correlation of eclipse signals observed consecutively rather than in parallel. |
We demoustrated. using [Kepler data. that despite the formal violation of the "justantanecous nixiug model. he proposed algoritlin is able to retrieve the desired signal componcut with good accuracy. | We demonstrated, using Kepler data, that despite the formal violation of the `instantaneous mixing model', the proposed algorithm is able to retrieve the desired signal component with good accuracy. |
Such an application is particularly Huportant for treating variability of the host-star which can significantly impair the quality of the final science result (eg.Czeslaetal.2009:Doisse 2011). | Such an application is particularly important for treating variability of the host-star which can significantly impair the quality of the final science result \citep[eg.][]{czesla09, boisse11, aigrain11,ballerini11}. |
. It is furthermore interesting to note that pre and post-processing steps οιο, wavelets (Carter&Winn 2009).. Fourier based techuiques (Waldiiauuetal. 2011).. de-correlatiou using iustrunienut state parameters (Swainetal. 2008))). do not break the iustautancous mixing model aud cau be run in conjunction with ICA methods. | It is furthermore interesting to note that pre and post-processing steps (e.g. wavelets \citep{carter09}, Fourier based techniques \citep{waldmann11}, de-correlation using instrument state parameters \citep{swain08}) ), do not break the instantaneous mixing model and can be run in conjunction with ICA methods. |
This lakes inclependcut conrponent analysis a very powerful and versatile tool for uou-parainetric correlation of exoplauetary data sets; | This makes independent component analysis a very powerful and versatile tool for non-parametric de-correlation of exoplanetary data sets. |
Iun the light of searching aud characterising ever smaller and fainter cxoplanctary tarects. the development of novel de-trending routines becomesincreasingly critical. | In the light of searching and characterising ever smaller and fainter exoplanetary targets, the development of novel de-trending routines becomes increasingly critical. |
Based ou the coucepts of blind source deconvolution of imstantancously nüxed signals. we lave presented a first step towards 10n-paranmetriie corrections and data filters tha do not require additional information ou the systematic noise ofthe iustruiieut or stellar activity. | Based on the concepts of blind source deconvolution of instantaneously mixed signals, we have presented a first step towards non-parametric corrections and data filters that do not require additional information on the systematic noise of the instrument or stellar activity. |
Such algoritlius have two imiportaut applications: 1) For instruments that lack a calibration plan at the accuracy of 1 in flux variatio- which is required for spectroscopy of exoplauctary atmospheres. the spectroscopic signatures become inherently eutaueled aud dependent ou the method used to correct mstrumneut aud other svstenmiaties in the data. | Such algorithms have two important applications: 1) For instruments that lack a calibration plan at the accuracy of $^{-4}$ in flux variation, which is required for spectroscopy of exoplanetary atmospheres, the spectroscopic signatures become inherently entangled and dependent on the method used to correct instrument and other systematics in the data. |
The de-correlation of spectroscopic data was demoustrated using two HST/NICMOS data sets. | The de-correlation of spectroscopic data was demonstrated using two HST/NICMOS data sets. |
2) Detections of faint cxoplanctary eclipses are often made cifieult by time-correlated activity of the host star. | 2) Detections of faint exoplanetary eclipses are often made difficult by time-correlated activity of the host star. |
We demonstrated. using a siugle Kepler tiue serics. that much of the stellar variability can be removed in time series that span several exoplauctary eclipse events. | We demonstrated, using a single Kepler time series, that much of the stellar variability can be removed in time series that span several exoplanetary eclipse events. |
The algorithun proposed is a powerful tool for lighteurve de-treudius. which can be used by its OWL or iu conjunction with anv other type of data filtering or cleaning technique. | The algorithm proposed is a powerful tool for lightcurve de-trending, which can be used by its own or in conjunction with any other type of data filtering or cleaning technique. |
This becomes an invaluable advantage for data analysis when he instruments response function is uukuownu or oorlv. eliaracterised. | This becomes an invaluable advantage for data analysis when the instrument's response function is unknown or poorly characterised. |
LDP.W. would like to thank Prof. E. Feigelson. he referee. Dr. C. Tiuctti. Dr. F. Abdalla aud Dr. ο, Fossey for comunenuts and suggestions that relped to ereatle inuprove this paper. | I.P.W. would like to thank Prof. E. Feigelson, the referee, Dr. G. Tinetti, Dr. F. Abdalla and Dr. S. Fossey for comments and suggestions that helped to greatly improve this paper. |
LP.W. is supported by au STFC Studeutship. | I.P.W. is supported by an STFC Studentship. |
To verify this qualitative picture would. however. require detailed numerical calculations. | To verify this qualitative picture would, however, require detailed numerical calculations. |
This work has been financed by the KBN erants 2P03D-01016. and. 2P03D-02117. | This work has been financed by the KBN grants 2P03D-01016, and 2P03D-02117. |
Support from Multiprocessor Systems Group at Nicholas Copernicus University’s Computer Centre in providing facilities for the time/memory-consumine Monte. Carlo caleulations is appreciated. | Support from Multiprocessor Systems Group at Nicholas Copernicus University's Computer Centre in providing facilities for the time/memory-consuming Monte Carlo calculations is appreciated. |
We are erateful to Gottfried. WKanbach and Alaurice Cos for valuable discussions on processing and analysis of high-energy. data for the Vela pulsar. | We are grateful to Gottfried Kanbach and Maurice Gros for valuable discussions on processing and analysis of high-energy data for the Vela pulsar. |
We thank the anonvmous referee for bringing our attention to the paper by Sturner et al. ( | We thank the anonymous referee for bringing our attention to the paper by Sturner et al. ( |
1995). | 1995). |
For an uuderdeuse universe. the density is described by and the ecueral solution of equation(6)) is with e» deteriiuiug peculiar motion as before: the particular solution is and is shown in Fie. € | For an underdense universe, the density is described by and the general solution of \ref{eq:freenewton}) ) is with $c_2$ determining peculiar motion as before; the particular solution is and is shown in Fig. \ref{fig:underdense}) ). |
For the zero-energv/fiat wniverse with autieravitv. the deusitv is and the ecueral solution to equation (6)) is The integral does not appear to be expressable in closed form. which makes working with it somewhat iucouveuieut. | For the zero-energy/flat universe with antigravity, the density is and the general solution to equation \ref{eq:freenewton}) ) is The integral does not appear to be expressable in closed form, which makes working with it somewhat inconvenient. |
Dowever. some characteristics of if cau be derived which are sufficient for the present purpose. | However, some characteristics of it can be derived which are sufficient for the present purpose. |
For an uudoerdeuse universe. the distauce between the free particle aud the backgrouud particle with the same asviuptotie speed is more complicated than in the critical ease. but as y becomes large approaches a constant. | For an underdense universe, the distance between the free particle and the background particle with the same asymptotic speed is more complicated than in the critical case, but as $\eta$ becomes large approaches a constant. |
That is. the free particle stavs at least this far from its correspouding backeound particle. even at infinite times. | That is, the free particle stays at least this far from its corresponding backgound particle, even at infinite times. |
Iu the situation of a fat universe with the cosmological coustaut. the peculiar motion (62) function cau be approximated at carly times by | In the situation of a flat universe with the cosmological constant, the peculiar motion $c_2$ ) function can be approximated at early times by |
supplemented: by the heat conduction equation (Maciolek-Niedzwwiecki. Wrolik Zdziarski 1997: ROO). | supplemented by the heat conduction equation }ek-Nied\'{z}wwiecki, Krolik Zdziarski 1997; R99). |
The transition from the hot to cold. solution branch is very. sharp and its location can be obtained from the general condition of radiativeconductive equilibrium (R6ezaaisska 2000). | The transition from the hot to cold solution branch is very sharp and its location can be obtained from the general condition of radiative–conductive equilibrium (R\'{o}\\.{z}aa\'{n}sska 2000). |
Phe condition gives the temperature at the transition point where 7i is the surface temperature. | The condition gives the temperature at the transition point where $\Ts$ is the surface temperature. |
For purely radiation heated disc atmosphere 72=ic. where Z1c is the temperature (Melee Begelman 1990: 16-Zaasska C'zerny20008). | For purely radiation heated disc atmosphere $\Ts = \TIC$, where $\TIC$ is the inverse-Compton temperature (McKee Begelman 1990; R\'{o}\\.{z}aa\'{n}sska Czerny2000a). |
When the al? heating is also included. 7; is somewhat higher than Zic (16 2000). | When the $\alpha P$ heating is also included, $\Ts$ is somewhat higher than $\TIC$ (R\'{o}\\.{z}aa\'{n}sska 2000). |
In a good. approximation. the transition occurs where the upper. stable branch of the = 7 curve changes to the middle. unstable branch (RCOG. R99. NIxI). The resulting structure is then basically two-lavereed: the upper. hot laver (LIL) and the lower. cold laver. although there is a thin transition laver in between the two (1199). | In a good approximation, the transition occurs where the upper, stable branch of the $\Xi$ $T$ curve changes to the middle, unstable branch (RC96, R99, NKK), The resulting structure is then basically two-layered: the upper, hot layer (HL) and the lower, cold layer, although there is a thin transition layer in between the two (R99). |
For a sulliciently hard illuminating spectra hhieh ic) iron can be considered completely: ionized in the hot [aver. while it recombines to below in the cold [aver (NIxIEx: see also 22 in Zvveki Czerny 1904). | For a sufficiently hard illuminating spectra high $\TIC$ ) iron can be considered completely ionized in the hot layer, while it recombines to below in the cold layer (NKK; see also 2 in Żyycki Czerny 1994). |
The model parameter crucial for spectral/timing predictions is the thickness of the hot laver. mua. | The model parameter crucial for spectral/timing predictions is the thickness of the hot layer, $\tauh$. |
The larger the thickness. the smaller. the fraction of primary radiation penetrating to the cold. laver anc giving rise {ο the usual reprocessecl component. | The larger the thickness, the smaller the fraction of primary radiation penetrating to the cold layer and giving rise to the usual reprocessed component. |
The reprocessed. componen is further C'omptonized as the photons escape through the hot laver. | The reprocessed component is further Comptonized as the photons escape through the hot layer. |
The result (at least in the limited energy. ban corresponding to or data) is à reprocessec component with amplitude reduced compared to the usua eeometrical factor OΕπ (Iq. 8)). | The result (at least in the limited energy band corresponding to or data) is a reprocessed component with amplitude reduced compared to the usual geometrical factor $\Omega/4\pi$ (Eq. \ref{equ:refl}) ), |
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