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The data reduction was carried out with the software package MUA+EWS-1.3. | The data reduction was carried out with the software package MIA+EWS-1.3. |
This software consists of two independent reduction programs (ATTA and EWS) which were both applied [or comparison. | This software consists of two independent reduction programs (MIA and EWS) which were both applied for comparison. |
The software. further information and manuals can be downloaded from theInternet?. | The software, further information and manuals can be downloaded from the. |
. A general description of the basic data reduction steps is also given in Leinertetal.(2004) and we refer the reader to this paper. | A general description of the basic data reduction steps is also given in \citet{leinert} and we refer the reader to this paper. |
As the results derived with MIA and EWS agreed quite well (overall differences in the correlated Πιν <8%). we decided to show only plots resulting from the MILA package. | As the results derived with MIA and EWS agreed quite well (overall differences in the correlated flux $\leq 8\%$ ), we decided to show only plots resulting from the MIA package. |
FU Ori was found to be a binary svstem by Wang | FU Ori was found to be a binary system by \citet{wang}. |
etal.(2004).. confirmed the detected companion and concluded that il was a voung star of spectral tvpe lx showing clear NIR excess. | \citet{reipurth} confirmed the detected companion and concluded that it was a young star of spectral type K showing clear NIR excess. |
In (he same paper il was also stated that il was very unlikely that the binary component (FU Ori $) triggered the outburst of FU Ori observed in the late 1930s. | In the same paper it was also stated that it was very unlikely that the binary component (FU Ori S) triggered the outburst of FU Ori observed in the late 1930s. |
Knowing about (he existence of the fainter companion the integration time of some MIDI acquisition images was increased in order to derive N-band photometry for both components. | Knowing about the existence of the fainter companion the integration time of some MIDI acquisition images was increased in order to derive N-band photometry for both components. |
In three images FU Ori 5 was clearly visible (Figure 1)) and aperture photometry could be applied to the observations. | In three images FU Ori S was clearly visible (Figure \ref{Image}) ) and aperture photometry could be applied to the observations. |
The results are summarized in Table 2.. | The results are summarized in Table \ref{fluxes}. |
Interestingly. FU Ori $ shows relatively a hieher N-band excess (han FU Ori itself. | Interestingly, FU Ori S shows relatively a higher N-band excess than FU Ori itself. |
This can be explained in terms ol differences in (he geometry of an assumed. circumstellar disks (e.g.. larger flaring angle or smaller clisk inclination). | This can be explained in terms of differences in the geometry of an assumed circumstellar disks (e.g., larger flaring angle or smaller disk inclination). |
We furthermore derive a separation between the components of 0." 4842:0."01 and a position angle of 162.5°44.1° (measured from north eastwards) which is in good agreement with the results from Wangetal.(2004) and Πορ&Aspin(2004).. | We furthermore derive a separation between the components of $''$ $\pm$ $''$ 01 and a position angle of $^\circ\pm$ $^\circ$ (measured from north eastwards) which is in good agreement with the results from \citet{wang} and \citet{reipurth}. |
Our errors are standard deviations based on Gaussian fitting of 3 independent images. | Our errors are standard deviations based on Gaussian fitting of 3 independent images. |
For an assumed distance of 450 pe the separation corresponds to 217.844.5 AU. | For an assumed distance of 450 pc the separation corresponds to $\pm$ 4.5 AU. |
This pulsar has been selected lor additional observations because of its unusually long nulling periods. | This pulsar has been selected for additional observations because of its unusually long nulling periods. |
Alter the initial. post-discovery study by Caciwell (1997). we have observed 1 on several occasions between 2000 and 2002. | After the initial, post-discovery study by Cadwell (1997), we have observed it on several occasions between 2000 and 2002. |
As shown in Fig. | As shown in Fig. |
5b. the pulsar spends of the time in a “quasi-null” state. | 5b, the pulsar spends of the time in a “quasi-null” state. |
The "on-states" occur once every 400-600. periods and last. on average. for ~LO0 periods. | The “on-states” occur once every 400-600 periods and last, on average, for $\sim$ 100 periods. |
A more detailed inspection of the on-states (Fig. | A more detailed inspection of the on-states (Fig. |
5c) reveals that these burst-like emissions are quite similar in shape and duration and thev decay into a null state in a manner (hat is quite reasonably described by a /e/" function wilh 7 ~45 seconds. | 5c) reveals that these burst-like emissions are quite similar in shape and duration and they decay into a null state in a manner that is quite reasonably described by a $te^{-t/\tau}$ function with $\tau\sim$ 45 seconds. |
A comparison of the single pulse characteristics of PSR J1752+2359 with a more tvpical nulling behavior of PSR J1649+2533 (Fig. | A comparison of the single pulse characteristics of PSR J1752+2359 with a more typical nulling behavior of PSR J1649+2533 (Fig. |
5a) clearly shows that the (wo phenomena exhibit distinctly. different morphologies. because PSR. J1752+2359 switches olf eradually. rather than suddenly. as observed in the nulling pulsars (Ritchines 1976. Deich et al. | 5a) clearly shows that the two phenomena exhibit distinctly different morphologies, because PSR J1752+2359 switches off gradually, rather than suddenly, as observed in the nulling pulsars (Ritchings 1976, Deich et al. |
1936. Biges 1992). | 1986, Biggs 1992). |
Moreover. the astonishing morphological similarity of the individual on-states aud (heir very similar Ginescales ancl repetition rates are unique among the nulline pulsars. | Moreover, the astonishing morphological similarity of the individual on-states and their very similar timescales and repetition rates are unique among the nulling pulsars. |
To our knowledge. these emission characteristics of PSR. J17524-2359 make it an exceptional member of the pulsar population. | To our knowledge, these emission characteristics of PSR J1752+2359 make it an exceptional member of the pulsar population. |
The pulsars discussed in this paper fall into three evolutionary categories as indicated in the P-P diagram shown in Fig. | The pulsars discussed in this paper fall into three evolutionary categories as indicated in the $P$ $\dot P$ diagram shown in Fig. |
6. | 6. |
When interpreted in terms of the rotating magnetic dipole model (Pacini 1968: Gold 1963). the location of a pulsar in this diagram depends on the strength of its surface magnetic field (—(PP) !2) and on its characteristic age (7=P/2P), | When interpreted in terms of the rotating magnetic dipole model (Pacini 1968; Gold 1968), the location of a pulsar in this diagram depends on the strength of its surface magnetic field $\sim (P\dot P)^{1/2}$ ) and on its characteristic age $\tau=P/2\dot P$ ). |
Within this framework. the relatively voung pulsar PSR JOS38+2817. possibly still associated wilh the supernova remnant 5147. appears close to the cluster of similarly voung objects. whereas (he 16 “normal” pulsars fall in (he area (hat is occupied by Che largest. intermediate age category, | Within this framework, the relatively young pulsar PSR J0538+2817, possibly still associated with the supernova remnant S147, appears close to the cluster of similarly young objects, whereas the 16 “normal” pulsars fall in the area that is occupied by the largest, intermediate age category. |
The two possible exceptions in this eroup are PSR J1813—1522 and PSh J19084-2351 (Table 4). which have characteristic ages well above LOO Myr. magnetic fields below LO! Gauss. and may more properly be categorized as mildly ~reeveled” pulsars which originated [rom disrupted binaries (e.g. Bailes 1989). | The two possible exceptions in this group are PSR J1813+1822 and PSR J1908+2351 (Table 4), which have characteristic ages well above 100 Myr, magnetic fields below $^{11}$ Gauss, and may more properly be categorized as mildly “recycled” pulsars which originated from disrupted binaries (e.g. Bailes 1989). |
Finally. as expected. the millisecond pulsar PSR J17094-2313 is found among other members of the old. recycled pulsar population (Phinney Kulkuni 1994). | Finally, as expected, the millisecond pulsar PSR J1709+2313 is found among other members of the old, recycled pulsar population (Phinney Kulkarni 1994). |
flucestogetherwithacailablencar$00017TIS00307TIE22221 fraredandopticaldata(Abavafianetal,2009 | fluxes together with available near-infrared and optical data \citep{aba09,far09b,skr06,mar05, mcc08,lan07,
kil06,cmc06,sts06,zac05,den05,mon03}. |
:FarihiHeLando richtargetstars. Iilieetul.(2006) tis and -WNeolors. suchthatthell andy. band fluccesofthesestarsarcanchorcdtothce2 ALAS SJ bandflic. andhingceuponitsacearacy | For a few of the metal-rich target stars, \citet{kil06} has published spectroscopically derived $J-H$ and $H-K$ colors, such that the $H$ - and $K$ -band fluxes of these stars are anchored to the 2MASS $J$ -band flux, and hinge upon its accuracy. |
ERI PIM photometry 32520pd 11225-079aretaken from Farihi(2009). | IRTF $JHK$ photometry of $-$ 3252 and $-$ 079 are taken from \citet{far09b}. |
ΤΠhere vos ue.Me"nhel(itlfeudi 2005) anduicarultraviolet photouictriefluecsarcplottedtoassisti αμnest (αμ f | Where available,; \citealt{mar05}) ) far- and near-ultraviolet photometric fluxes are plotted to assist in establishing the appropriate brightness level of the white dwarf. |
auneρα feli I f LowFat | However, because the interstellar extinction to these stars is not known, the fluxes were used with some caution. |
Five of the target white dwrfs come from the Ihuuburg Schinidt (IIS: Tagenetal. 19953) aud the Ihuuburg ESO Sclinidt (TE: Wisotzkietal. 1996)) quasar survevs. | Five of the target white dwarfs come from the Hamburg Schmidt (HS; \citealt{hag95}) ) and the Hamburg ESO Schmidt (HE; \citealt{wis96}) ) quasar surveys. |
The coordinates in SIMDBAD for these five white dwarfs lave poor precision. and are often nussing frou the white dwarf catalog [οςους&Siow(1999. 2008). | The coordinates in SIMBAD for these five white dwarfs have poor precision, and are often missing from the white dwarf catalog \citet{mcc99,mcc08}. |
. The correct positions for these white dwarts were providedby the SPY team (BR. Napiwotzki 2007. private communication). aud their coordinates m the IRAC images are given in Table L. | The correct positions for these white dwarfs were provided by the SPY team (R. Napiwotzki 2007, private communication), and their coordinates in the IRAC images are given in Table \ref{tbl4}. |
The WE aud IIS white dwarfs teud to suffer frou a lack of accurate eround-based photometry iu the literature. | The HE and HS white dwarfs tend to suffer from a lack of accurate ground-based photometry in the literature. |
À typical star has entries in a few photographic catalogs errors). variable quality +’-band CCD photometry from the Carlsberg Aleridian Catalog (CAIC: CopenhagenUniversityObservatory 2006)). and a DENIS f-haucl nieasurenüent if in the southern hemisphere. | A typical star has entries in a few photographic catalogs errors), variable quality $r'$ -band CCD photometry from the Carlsberg Meridian Catalog (CMC; \citealt{cmc06}) ), and a DENIS $I$ -band measurement if in the southern hemisphere. |
For | 0093. | O716. 1630. aud 22229| 2335. the oulv nea-πο photome trvis 40)the diyd ou ΕΠ μμ. üt:pa Meet Is jjgU08: haspublishedspectroscopicallgdο. less relable owing the faintuess of these MEShautslieetal.2006). | For $+$ 0093, $+$ 0746, $-$ 1630, and $+$ 2335, the only near-infrared photometry is the 2MASS $J$ -band value, as the $H$ - and $K_s$ -band data are substantially less reliable owing to the faintness of these stars \citep{skr06}. |
. ἀλλavailable photometric data for cach white dwarf +] Won, pure dile an14dανaIAftoruds salleatti WE 20095). | All available photometric data for each white dwarf were fitted by a pure hydrogen or pure helium white dwarf model with $\log\,[g\,({\rm cm\,s}^{-2})]=8.0$ and effective temperature to the nearest K \citep{koe09b}. |
. Because the surface eravity aud the atinosplieric conrposition ave fixed. and because the effective temperature resolution ids lanited. the models showu in the fleures do not necessarilv describe acceuratelv or stronely constrain the effective teniperature of the stars. | Because the surface gravity and the atmospheric composition are fixed, and because the effective temperature resolution is limited, the models shown in the figures do not necessarily describe accurately or strongly constrain the effective temperature of the stars. |
Rather. the atimospheric model fits predict the expected level of mud-intrared photospheric cussion. aud | Rather, the atmospheric model fits predict the expected level of mid-infrared photospheric emission, and |
al zqq4Q=20 OF ται15 and remains reionized. and a model with a "double reionization ∣↽≻≀↧↪∖⊽≼↲≼⇂∪∐⊔∐↲∐↓⋯⇂≼↲↥∪↓∙↗⋅⋅↕∐∖∖↽↥∐≺∢∐⊔∐↲↧∕∐↕∖↽≼↲↕⋅⊳∖⊽≼↲↕⋟∖⇁↕⋅≼↲↕∪∐↕∠≼↲≼⇂≀↧↴↥⊳∶↓⋅≺⋅↕⊓↓↓∶⊥↱≻⋅↱≻⋅↕⋅≼↲≺∢∪∐↓∣↽≻∐∐↲⋟∖⊽ 0»: ⇁⋅ ⋅⋅⋅ ⊏↓⊐ − ∣↽≻⋡∖↽⊳∶↓⋅≺⋅∙⊲∶↓≊↽⊰⋅↱≻⋅≀↧↴∐≼⇂≼↲⇀↸↕↽≻≼↲↕⋅↕≼↲∐≺∢≼↲⋟∖⊽≀↕↴ | at $z_{\rm reion} =20$ or $z_{\rm reion} =15$ and remains reionized, and a model with a `double reionization' based on the model of \citet{cen:02}, in which the Universe is reionized at $z^{(1)}_{\rm reion}=15.5$, recombines by $\zrec =13.5$, and experiences a second reionization at $\zion2=6$. |
⋟∖⊽≼↲≺∢∪∐≼⇂↕⋅≼↲↕∪∐↕∠≀↧↴⊔∪∐≀↧↴↥⊳∶↓⋅⊲⋅↕⊓↓↓∶↻⋅∏∐↲≼∐↕≻⋟∖⊽∐↥⊔∐↲⋟∖⊽↥≀↧↴↕⋅⋅ − ⋅⋅⋅ (2)⋅↴ ⋅⋅↽⋅⋅ formation history are due to the suppression of gas infall by the photoionizing background. and demonstrate that a substantial faction of the total star formation in the Universe would be taking place in small mass halos in the absence any. kind of feedback. | The `dips' in the star formation history are due to the suppression of gas infall by the photoionizing background, and demonstrate that a substantial fraction of the total star formation in the Universe would be taking place in small mass halos in the absence any kind of feedback. |
Models N'I-3a.c. shown in Fig. | Models MT-3a–c, shown in Fig. |
3bb. all assume the ?9 double-reionization history described above. and also include supernova feedback. | \ref{fig:sfrd}b b, all assume the \citet{cen:02} double-reionization history described above, and also include supernova feedback. |
Note that the ‘dips’ are now much ess dramatic. because much of the gas has already been removed from the small halos that are affected by photoionization squelching. | Note that the `dips' are now much less dramatic, because much of the gas has already been removed from the small halos that are affected by photoionization squelching. |
An uncertain aspect of implementing supernova eedback in semi-analvlic models is the fate of the reheated gas. | An uncertain aspect of implementing supernova feedback in semi-analytic models is the fate of the reheated gas. |
In model NI-3a. the gas reheated by SN is removed [rom the disk but retained in the halo (retention feedback). in MTI-3b the reheated gas is removed from the disk dark matter halo of all galaxies ejection feedback). and in model MT-3e (the gas is ejected from the halo only if the halo virial velocity is less than 100 km/s. This threshold for ejection of gas by super-winds is notivaled by theoretical arguments (2) ancl observations of nearby galaxies (?).. | In model MT-3a, the gas reheated by SN is removed from the disk but retained in the halo (`retention' feedback), in MT-3b the reheated gas is removed from the disk dark matter halo of all galaxies (`ejection' feedback), and in model MT-3c the gas is ejected from the halo only if the halo virial velocity is less than 100 km/s. This threshold for ejection of gas by super-winds is motivated by theoretical arguments \citep{dekel-silk} and observations of nearby galaxies \citep{martin:99}. |
The elobal star formation rate changes by as much as a [actor of six al z~3 depending on these choices. but by less than a factor of two at z215. | The global star formation rate changes by as much as a factor of six at $z\sim3$ depending on these choices, but by less than a factor of two at $z\ga15$. |
Note that the ‘dip’ following the first reionizalion is considerably more pronounced in the model with ejection feedback. as gas which has been accreted before reionization aud ejected is not allowed to re-collapse in small halos while the photoionizing background is "switched on’. | Note that the `dip' following the first reionization is considerably more pronounced in the model with ejection feedback, as gas which has been accreted before reionization and ejected is not allowed to re-collapse in small halos while the photoionizing background is `switched on'. |
We consider models NT-3a and A[T-3e to be the most realistic of the models considered here. | We consider models MT-3a and MT-3c to be the most realistic of the models considered here. |
Models with similar ingredients and parameter values have been shown to reproduce the Iuminosity function of galaxies al zo0 (SP) as well as of Lyinan break galaxies al z~3 (SPF). | Models with similar ingredients and parameter values have been shown to reproduce the luminosity function of galaxies at $z\sim0$ (SP) as well as of Lyman break galaxies at $z\sim3$ (SPF). |
In Fig. | In Fig. |
d we show the cumulative number of hydrogen-ionizing photons per hydrogen atom in the Universe. produced by Pop I stars in our fiducial models (NT-3a and NET-3€). and by our model for Pop ILE star formation. | \ref{fig:nphot} we show the cumulative number of hydrogen-ionizing photons per hydrogen atom in the Universe, produced by Pop II stars in our fiducial models (MT-3a and MT-3c), and by our model for Pop III star formation. |
? (?).. (e2).. | \citet{leitherer:99} \citep{bkl:01}. \citep[e.g.,][]{spergel:03}, |
that more luminous blazers tend to have steeper spectra in the LAT frequency range 2010a). | that more luminous blazars tend to have steeper spectra in the LAT frequency range . |
. It is usually agreed that the electrons producing the IC-component. are responsible also for at least the high frequency part of the svnchrotron component 1996). | It is usually agreed that the electrons producing the IC-component are responsible also for at least the high frequency part of the synchrotron component . |
. The origin of the seed photons lor the IC-component is somewhat more uncertain: they. could either be external photons or (he intrinsically produced svinclirotron photons (S5C-models). | The origin of the seed photons for the IC-component is somewhat more uncertain; they could either be external photons or the intrinsically produced synchrotron photons (SSC-models). |
Since (here is a rough correlation between emission line strength and Iuminositv of the IC- it has been suggested that external photons dominate the cooling in the most luminous blazars. while (hie least luminous blazars correspond to the H | Since there is a rough correlation between emission line strength and luminosity of the IC-component, it has been suggested that external photons dominate the cooling in the most luminous blazars, while the least luminous blazars correspond to the SSC-case. |
is normally assumed that the dominant contribution comes [rom first order inverse Compton scattering2009). | It is normally assumed that the dominant contribution comes from first order inverse Compton scattering. |
. The spectral properties are then determined mainly bv the energy distribution of the injected relativistic electrons. | The spectral properties are then determined mainly by the energy distribution of the injected relativistic electrons. |
Another possibility is that the source properties are such that the IC-component is due to multiple inverse Compton scatterings (ATIC). | Another possibility is that the source properties are such that the IC-component is due to multiple inverse Compton scatterings (MIC). |
In (his case. its spectral properties depend to a Iarge extent also on the scattering process itself. | In this case, its spectral properties depend to a large extent also on the scattering process itself. |
In addition to different spectral characteristics. (ime variations are expected to have properties which differ between ancl first order scattering models: (his is particular true for correlations between fIux variations in the svnchrotron and IC components. | In addition to different spectral characteristics, time variations are expected to have properties which differ between MIC-models and first order scattering models; this is particular true for correlations between flux variations in the synchrotron and IC components. |
In MIC-models. were discussed in connection with the observed properties of the radio/1m outbursts in blazars. | In MIC-models were discussed in connection with the observed properties of the radio/mm outbursts in blazars. |
Motivated by the Fermi-LAT observations of blazars. the present paper extends this discussion to the high frequency radiation. | Motivated by the -LAT observations of blazars, the present paper extends this discussion to the high frequency radiation. |
Guided bv the numerical results in DAOO. an analvtical description of MIC-mocdels is developed in 32. | Guided by the numerical results in BA00, an analytical description of MIC-models is developed in 2. |
It is emphasized that the origin of the seed photons only marginally affects the results. | It is emphasized that the origin of the seed photons only marginally affects the results. |
The consequences for the svnchrotron component is briefly. discussed in 33. im particular the effects of cooling on the svuchrotron sell-absorption frequency. | The consequences for the synchrotron component is briefly discussed in 3, in particular the effects of cooling on the synchrotron self-absorption frequency. |
The most Iuninous blazars often exhibit a distinct X-ray minimum in (heir spectral energy distribution. whieh is hard to account [or in one-zone MIC-models1996). | The most luminous blazars often exhibit a distinct X-ray minimum in their spectral energy distribution, which is hard to account for in one-zone MIC-models. |
. The extensions of standard one-zone models needed (to accomodate such minima in MIC-mocels are discussed in 44. | The extensions of standard one-zone models needed to accomodate such minima in MIC-models are discussed in 4. |
In MIC-models. the relativistic electrons radiate most of (heir energy in the IXlein-Nishina limit. | In MIC-models, the relativistic electrons radiate most of their energy in the Klein-Nishina limit. |
Hence. with a typical bulk Lorentz [actor DP~10 for the emission region in blazars. the photons detected by Ferm-LAT should be up-scattered from a [requency range overlapping that of SwifI-DAT. | Hence, with a typical bulk Lorentz factor $\Gamma \sim 10$ for the emission region in blazars, the photons detected by -LAT should be up-scattered from a frequency range overlapping that of -BAT. |
have discussed the properties of a sample of blazars detected both by FermiLAT ancl | have discussed the properties of a sample of blazars detected both by -LAT and }-BAT. |
This is taken as the starting point in | This is taken as the starting point in |
probably loosing iu the 1590 PSPC fields nüssiug frou the WGAÀ. | probably loosing in the 1590 PSPC fields missing from the WGA. |
There are 180 poiutines (LL for A and 66 for Ftype stars) in which 36 new A and 28 new Ftvpe stars could have been detected. | There are 180 pointings (114 for A– and 66 for F–type stars) in which 36 new A– and 28 new F–type stars could have been detected. |
If sve apply the detection rate that we fouud for the WGA fields. we should have missed about ~7 and ~123 stars for he À and Ftype. respectively. | If we apply the detection rate that we found for the WGA fields, we should have missed about $\sim 7$ and $\sim 13$ stars for the A– and F–type, respectively. |
If we consider that most of them would be in binary svstenis. we are really missingo very few stars. | If we consider that most of them would be in binary systems, we are really missing very few stars. |
Thus. the stars that we are presenting here are not ALL he single Av aud Fotype stars possibly etected in PSPC xntiues. | Thus, the stars that we are presenting here are not ALL the single A– and F–type stars possibly detected in PSPC pointings. |
But this docs not affect our goals. since we are rot deviving a Xrav Iuniünosity "unction. | But this does not affect our goals, since we are not deriving a X–ray luminosity function. |
Our ai is a) to ook for supposedly single A stars tha are N-rav οτους, aud b) to study the characteristics of j0 Nray enadssiou of sinele A and F stars. | Our aim is a) to look for supposedly single A stars that are X-ray emitters, and b) to study the characteristics of the X–ray emission of single A and F stars. |
Besides. sources are randomly missed roni the WCA and he results should uot be biased. | Besides, sources are randomly missed from the WGA and the results should not be biased. |
The final πολλοι, here preseuted. consists of 52 supposedly single stars. Α and carly Etype stars. associated with Xrav sources in ROSAT PSPC fields (29 fields for the Aype stars and 15 fields for the carly F type stars). | The final sample, here presented, consists of 52 supposedly single stars, A– and early F–type stars, associated with X–ray sources in ROSAT PSPC fields (29 fields for the A–type stars and 45 fields for the early F type stars). |
The saluple contains 230 iainsequence stars, 9 subgiauts. l giaut. 2 supergiants; 6 stars with uncertain spectral ype and/or Iuninositv class aud L stars without auv iuforuration about the luminosity class: there are 4 chenucally )eculiar stars aud three ó Scuti stars 1121953. 132 iux 5859119). | The sample contains 30 main–sequence stars, 9 sub–giants, 1 giant, 2 super–giants, 6 stars with uncertain spectral type and/or luminosity class and 4 stars without any information about the luminosity class; there are 5 chemically peculiar stars and three $\delta$ Scuti stars 124953, 432 and 89449). |
In Table 1 we Bist the 52 stars by their ΠΙΟ uumiber (Coluuu 1). | In Table 1 we list the 52 stars by their HD number (Column 1). |
Coluunus 2. 2. Land 5 conain the spectral type. magnitue. DV color aud distaice of the stars in parsec. | Columns 2, 3, 4 and 5 contain the spectral type, magnitude, B--V color and distance of the stars in parsec. |
Except for 1121953. all these informations were taken frou the HIPPARCOS Outit. Catalogue (Perrvinan ct al. | Except for 124953, all these informations were taken from the HIPPARCOS Output Catalogue (Perryman et al. |
1997). | 1997). |
For 1121953 the informations are from joa οἳ al. ( | For 124953 the informations are from a et al. ( |
1995). | 1995). |
Note that the IIIPPARC'OS disaliceos often are different from those previously kuown or derived. from specroscopic parallaxes. | Note that the HIPPARCOS distances often are different from those previously known or derived from spectroscopic parallaxes. |
This could imply a refinemc of the spectroscopic classification (few sub-classes) for various stars. | This could imply a refinement of the spectroscopic classification (few sub-classes) for various stars. |
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