Datasets:
ReadingTimeMachine
/
rtm-sgt-ocr-v1

Dataset card Data Studio Files Community
1
source
stringlengths
1
2.05k
⌀
target
stringlengths
1
11.7k
‘The significance of this result remains to be evaluated.
The significance of this result remains to be evaluated.
The effect. of drzüning of He from the ionization zone Is to reduce the width of the instability strip. the blue edge moving towards the red edge. eventually leading to the disappearance of the instability strip when Le is sulliciently depleted. (Coxctal. 1979)...
The effect of draining of He from the ionization zone is to reduce the width of the instability strip, the blue edge moving towards the red edge, eventually leading to the disappearance of the instability strip when He is sufficiently depleted \citep{cox79}. .
“Pureottectal.(2000). has discussed. the effect of diffusion on pulsations in Am/Pm stars using the models by Iücher.Michaud&Turcotte (2000).
\citet{turcotte00} has discussed the effect of diffusion on pulsations in Am/Fm stars using the models by \citet{richer00}.
. One significant cilference with earlier models is that a substantial amount of He remains in the ionization zone.
One significant difference with earlier models is that a substantial amount of He remains in the ionization zone.
The blue edge of the instability strip for AmPim stars is sensitive to the magnitude of the abundance /variations and is thus indicative of the depth of mixing by turbulence.
The blue edge of the instability strip for Am/Fm stars is sensitive to the magnitude of the abundance variations and is thus indicative of the depth of mixing by turbulence.
‘Tureottectal.(2000). predict that pulsating mIm stars should lie in a confined region of the LR diagram close to the red edge of the 9 Set instability strip.
\citet{turcotte00} predict that pulsating Am/Fm stars should lie in a confined region of the HR diagram close to the red edge of the $\delta$ Sct instability strip.
However. Balonaetal.(2011). show that there is no relationship between the predicted Am/Em instability strip and the actual location of these stars in the Lt diagram.
However, \citet{balona11} show that there is no relationship between the predicted Am/Fm instability strip and the actual location of these stars in the HR diagram.
Ao particularly. interesting result of the pulsation analysis of Turcottectal.(2000) is the prediction of long-period ο mocles in A-type stars.
A particularly interesting result of the pulsation analysis of \citet{turcotte00} is the prediction of long-period g modes in A-type stars.
As the star evolves. the driving regions shift deeper into the star and the ο modes become gradually more and more excited.
As the star evolves, the driving regions shift deeper into the star and the g modes become gradually more and more excited.
Whereas p modes are stabilized through cillusion. ο modes tend to be excited as a result of that. process.
Whereas p modes are stabilized through diffusion, g modes tend to be excited as a result of that process.
Lt appears that. diffusion may act to enhance driving of long-period ο modes due to à significant increase in opacity due to iron-group elements.
It appears that diffusion may act to enhance driving of long-period g modes due to a significant increase in opacity due to iron-group elements.
This may have a bearing on the fact that nearly all A-type stars observed. byWepler have unexplainecl low-Lrequencics (Dalona2011)..
This may have a bearing on the fact that nearly all A-type stars observed by have unexplained low-frequencies \citep{balona11a}.
When dealing with objects with non-standard chemical composition. such as Am stars. it is crucial that the opacities are correctly caleulated.
When dealing with objects with non-standard chemical composition, such as Am stars, it is crucial that the opacities are correctly calculated.
This question has been investigated by several authors in recent vears.
This question has been investigated by several authors in recent years.
These stucies show that à non-standarc chemical composition of the stellar atmosphere alters the flux. distribution of the star or mocifies the profiles of the Balmer lines (Leone&Manfré(1007).. Catanzaro.Leone&Dall(2004))).
These studies show that a non-standard chemical composition of the stellar atmosphere alters the flux distribution of the star or modifies the profiles of the Balmer lines \citet{leo97}, \citet{cat04}) ).
Therefore a ctermination of Tr and logg based on a comparison between observed and computed Balmer-line profiles will not be correct unless one takes into account the metallicity of 10 star.
Therefore a determination of $_{\rm eff}$ and $\log g$ based on a comparison between observed and computed Balmer-line profiles will not be correct unless one takes into account the metallicity of the star.
Thus. even estimates based on standard analysis of the spectra may be in error when applied to Am/Im stars.
Thus, even estimates based on standard analysis of the spectra may be in error when applied to Am/Fm stars.
In this paper. we investigate the determination. of effective. temperature and. surface. gravity of the Ani star LbDe2e27411 (IR 1353. A8m) using spectra in the ESO archives.
In this paper, we investigate the determination of effective temperature and surface gravity of the Am star 27411 (HR 1353, A3m) using spectra in the ESO archives.
The purpose is to determine whether the stellar parameters of this star agree with those obtained. from Strommercn photometry and hence to test the reliability of the ellective temperature calibration applied to Amym stars.
The purpose is to determine whether the stellar parameters of this star agree with those obtained from Strömmgren photometry and hence to test the reliability of the effective temperature calibration applied to Am/Fm stars.
“Phe star was used bv Reyabchikova.Wochukhov&Bagnulo(2008) as a comparison in their study on the calcium stratification in Ap stars.
The star was used by \citet{ryab08} as a comparison in their study on the calcium stratification in Ap stars.
227411 is not known to pulsate.
27411 is not known to pulsate.
However. as we know fromAepler observations. pulsations in A and E stars with amplitudes too low to be detected from the ground are common.
However, as we know from observations, pulsations in A and F stars with amplitudes too low to be detected from the ground are common.
Atmospheric models obtained. with INELAASO. (Ixurucz1903) use precomputed line opacities in the form of opacity distribution functions (ODEs).
Atmospheric models obtained with ATLAS9 \citep{kur93} use precomputed line opacities in the form of opacity distribution functions (ODFs).
These are tabulatec for multiples of the solar metallicity. anc for various microturbulent: velocities.
These are tabulated for multiples of the solar metallicity and for various microturbulent velocities.
This approach allows very [as computation of model atmospheres. but with very. little Ilexibilitv in choice of chemical profile and microturbulen velocity.
This approach allows very fast computation of model atmospheres, but with very little flexibility in choice of chemical profile and microturbulent velocity.
While this is satisfactory for most. applications. it fails for chemically peculiar stars where a non-stanclare chemical composition profile is required.
While this is satisfactory for most applications, it fails for chemically peculiar stars where a non-standard chemical composition profile is required.
This can be done with XPLASI2 (Ixurucz1997).. which is essentially identica to ΑΕΛο. but uses the opacity sampling (OS) metho to evaluate line opacities.
This can be done with ATLAS12 \citep{kur97}, which is essentially identical to ATLAS9, but uses the opacity sampling (OS) method to evaluate line opacities.
In this study. we compare. the abundances of 11D227411. obtained with both codes. to determine ifthe use of XPLASI2 is essential.
In this study we compare the abundances of 27411 obtained with both codes to determine if the use of ATLAS12 is essential.
The result that Am stars are not confined to particular region of the 9 Set instability strip depends. to a large extent. on ellective temperatures and Lumiinosities estimated from Strómmegren photometry (Smalleyetal.2011).
The result that Am stars are not confined to particular region of the $\delta$ Sct instability strip depends, to a large extent, on effective temperatures and luminosities estimated from Strömmgren photometry \citep{smalley11}.
.. Le is not clear whether the calibration. derived from normal AL stars. can be applied to πιαι stars.
It is not clear whether the calibration, derived from normal AF stars, can be applied to Am/Fm stars.
In this paper we use synthetic Strómmgren photometry applied to mocels of Am/Em stars to investigate the reliability of fundamental parameters estimated. from. the photometry.
In this paper we use synthetic Strömmgren photometry applied to models of Am/Fm stars to investigate the reliability of fundamental parameters estimated from the photometry.
Finally. we discuss the relative numbers of pulsating ancl non-pulsating Am stars and compare these to the relative numbers of o Seuti and constant stars in the instability strip.
Finally, we discuss the relative numbers of pulsating and non-pulsating Am stars and compare these to the relative numbers of $\delta$ Scuti and constant stars in the instability strip.
From this comparison. onecan deduce the ellectiveness of pulsational driving in the ionization zone and comparethe ο abundance to that expected from cilfusion caleulations.
From this comparison, onecan deduce the effectiveness of pulsational driving in the ionization zone and comparethe He abundance to that expected from diffusion calculations.
two cones are let to vary independently.
two cones are let to vary independently.
The hydrogen column density and density of the material have been frozen at log(Ny)=20 and log(ng)—3 for simplicity.
The hydrogen column density and density of the material have been frozen at $\rm{N_H}$ )=20 and $\rm{n_{H}}$ )=3 for simplicity.
We recall that the simulated spectra by our Cloudy models are not sensitive to the value of these parameters within the explored range.
We recall that the simulated spectra by our Cloudy models are not sensitive to the value of these parameters within the explored range.
The best fit show ionization parameters of log(U)=0.71+0.07 and log(U)=0.13+0.09 for NW and SE cones, respectively.
The best fit show ionization parameters of $\rm{\pm}$ 0.07 and $\rm{\pm}$ 0.09 for NW and SE cones, respectively.
Nevertheless, this fit failed to simultaneously reproduce the two spectra, being the statistics rather poor (x?— 1.6).
Nevertheless, this fit failed to simultaneously reproduce the two spectra, being the statistics rather poor $\rm{\chi^2_r=1.6}$ ).
A much better fit is obtained (x?= 1.1) if we add to this phase (hereinafter NLR1) a second one (hereinafter NLR2) with NLR conditions but a lower value of the ionization parameter U. We note that due to convergence problems within XSPEC we had to freeze the ionization parameter of the second phase.
A much better fit is obtained $\rm{\chi^2_r=1.1}$ ) if we add to this phase (hereinafter NLR1) a second one (hereinafter NLR2) with NLR conditions but a lower value of the ionization parameter U. We note that due to convergence problems within XSPEC we had to freeze the ionization parameter of the second phase.
We checked that the best fit is obtained for a value of log(U)~--ᾱ, in both regions.
We checked that the best fit is obtained for a value of $\simeq\,-3$, in both regions.
The ionization parameters for the NLR1 phases are log(U)—0.92-0.2 and log(U)=0.340.2 for NW and SE cones, respectively.
The ionization parameters for the NLR1 phases are $\rm{\pm}$ 0.2 and $\rm{\pm}$ 0.2 for NW and SE cones, respectively.
These values are consistent with the parameters of the previous model.
These values are consistent with the parameters of the previous model.
The final fit can be seen in Figure 10 and the best fit parameters and fluxes for each model component are given in Table 4..
The final fit can be seen in Figure \ref{fig:spectrumCloudy} and the best fit parameters and fluxes for each model component are given in Table \ref{tab:cloudyfit}.
In order to show the confidence level of the fluxes, Fig.
In order to show the confidence level of the fluxes, Fig.
12 includes the iso-chi-squared flux contours of the normalizations of the two phases (i.e. NLR1 and NLR2) of the reflected components for the NW and SE (bottom) cones.
\ref{fig:isocontours} includes the iso-chi-squared flux contours of the normalizations of the two phases (i.e. NLR1 and NLR2) of the reflected components for the NW (top) and SE (bottom) cones.
Similar result is obtained (top)for the transmitted components.
Similar result is obtained for the transmitted components.
The NLR1 phase mostly contributes to the X-ray spectrum at the region between 0.8-0.9 keV and 0.53-0.7 keV, whereas the NLR2 phase contributes at 0.52 and 0.7-0.85 keV. The lowest value of the ionization parameter is very similar to that required to fit the optical spectrum of Mrk 573 (Kraemeretal.
The NLR1 phase mostly contributes to the X-ray spectrum at the region between 0.8-0.9 keV and 0.53-0.7 keV, whereas the NLR2 phase contributes at 0.52 and 0.7-0.85 keV. The lowest value of the ionization parameter is very similar to that required to fit the optical spectrum of Mrk 573 \citep{Kraemer09}.
In that work, the authors claimed a three phase 2009).component to explain the optical emission line spectrum.
In that work, the authors claimed a three phase component to explain the optical emission line spectrum.
The low-ionization gas accounts for the [OIIJAA3727À and [NIIJAA6548, emission, whereas the moderately ionized phase accounts for the
The low-ionization gas accounts for the $\rm{\lambda\lambda}$ and $\rm{\lambda\lambda}$ 6548, emission, whereas the moderately ionized phase accounts for the $\rm{\lambda\lambda}$.
Nonetheless, our NLR1 exhibits a value of U higher [OIITJAA5007À..than their highly ionized phase.
Nonetheless, our NLR1 exhibits a value of U higher than their highly ionized phase.
In addition, we expect a contribution of collisionaly ionized plasma Section ??)) to be present in the extended emission.
In addition, we expect a contribution of collisionaly ionized plasma (see Section \ref{sec:origin}) ) to be present in the extended emission.
(seeHowever, given the fact that its contribution to the nuclear spectrum is about696,, we assumed its contribution to the extended emission to be equal or smaller than in the nuclear case and did not try any fit given the low count level in these regions.
However, given the fact that its contribution to the nuclear spectrum is about, we assumed its contribution to the extended emission to be equal or smaller than in the nuclear case and did not try any fit given the low count level in these regions.
In terms of flux, the NLR1 phase is the of the extended emission flux.
In terms of flux, the NLR1 phase is the of the extended emission flux.
Its contribution in the NW cone is higher than in the SE cone.
Its contribution in the NW cone is higher than in the SE cone.
Moreover, the reflection component dominates the NW cone whereas the transmission component dominates the SE one.
Moreover, the reflection component dominates the NW cone whereas the transmission component dominates the SE one.
Tentatively, one could attribute this result to an orientation effect, being the NW cone located in the farthest side and the opposite for the SE one.
Tentatively, one could attribute this result to an orientation effect, being the NW cone located in the farthest side and the opposite for the SE one.
Indeed, 2-D spectroscopic observations etal1999) indicates that at least part of the(Ferruit NW cone ionized gas is red-shifted with respect to the systemic velocity.
Indeed, 2-D spectroscopic observations \citep{Ferruit99} indicates that at least part of the NW cone ionized gas is red-shifted with respect to the systemic velocity.
This could indicates that NW cone axis could be oriented behind the sky plane, although close to it.
This could indicates that NW cone axis could be oriented behind the sky plane, although close to it.
This orientation apparently contradicts the orientation proposed by Tsvetanov&Walsh(1992) based on reddening measurements, although we remark that their findings refer to the orientation with respect to the galaxy disk, not to the sky plane.
This orientation apparently contradicts the orientation proposed by \citet{Tsvetanov92} based on reddening measurements, although we remark that their findings refer to the orientation with respect to the galaxy disk, not to the sky plane.
Remarkably, the NLR2 phase accounts for the same flux in the NW and SE cones and nearly equal relative contribution of the reflection and transmission components.
Remarkably, the NLR2 phase accounts for the same flux in the NW and SE cones and nearly equal relative contribution of the reflection and transmission components.
Thus, this low ionization phase seems to be uniformly distributed along the cone area.
Thus, this low ionization phase seems to be uniformly distributed along the cone area.
At this point, we may conclude that soft extended X-ray emission is then powered by features coming from ionized gas in two phases under NLR-like conditions.
At this point, we may conclude that soft extended X-ray emission is then powered by features coming from ionized gas in two phases under NLR-like conditions.
Furthermore, one may question: is the [O extended emission structure powered by the same III]mechanism?
Furthermore, one may question: is the [O III] extended emission structure powered by the same mechanism?
The similarity between the [O III] structures and the soft X-ray emission points to a common origin for both
The similarity between the [O III] structures and the soft X-ray emission points to a common origin for both
of the LEC. whereas the scattering between the states below he resonance can explain the origin of the PR. component.
of the LFC, whereas the scattering between the states below the resonance can explain the origin of the PR component.
As the scattered radiation. concentrates along the ocal magnetic field. direction in the scattering region. the scattering in dilleren regimes. which hold at dilferent altitudes. does form different components in the pulse profile cause of the magnetosphere rotation.
As the scattered radiation concentrates along the local magnetic field direction in the scattering region, the scattering in different regimes, which hold at different altitudes, does form different components in the pulse profile because of the magnetosphere rotation.
The observed. LEC separation from the ÀPAA ~30. implies that the first-larmonic scattering takes place close to the light exlinder.
The observed LFC separation from the MP, $\Delta\lambda\sim 30^\circ$, implies that the first-harmonic scattering takes place close to the light cylinder.
It is important to note that the formation of the LEC at large enough altiude in the magnetosphere and. its orientation along the ambient magnetic field are strongly supported by he observed shift of the position angle of linear polarization with respec to that of the ALP.
It is important to note that the formation of the LFC at large enough altitude in the magnetosphere and its orientation along the ambient magnetic field are strongly supported by the observed shift of the position angle of linear polarization with respect to that of the MP.
The position angle shift appears approximately equal to the LEC-ALP separation in »ulse longitude.
The position angle shift appears approximately equal to the LFC-MP separation in pulse longitude.
This is also expected from an analysis of the rav-magnetic field geometry in the rotating magnetosphere.
This is also expected from an analysis of the ray-magnetic field geometry in the rotating magnetosphere.
In the two regimes considered. in the present paper. he scattering mainly results in the waves of the ordinary xolarization. whereas both the ordinary and extraordinary waves are believed to be present in the original radio beam.
In the two regimes considered in the present paper, the scattering mainly results in the waves of the ordinary polarization, whereas both the ordinary and extraordinary waves are believed to be present in the original radio beam.
llence. the PR and LEC should be strongly polarized. which is consistent with the observational data.
Hence, the PR and LFC should be strongly polarized, which is consistent with the observational data.
As long as the scattering efficiency is [large enough. [or the erowth of the scattered component to reach the stage of saturation. the [frequency dependence ofP and the variations ob E due to the pulse-to-pulse [uctuations of the inciden racio beam intensity and of the parameters of the scattering plasma do not alfect the scattered component substantially.
As long as the scattering efficiency is large enough for the growth of the scattered component to reach the stage of saturation, the frequency dependence of $\Gamma$ and the variations of $\Gamma$ due to the pulse-to-pulse fluctuations of the incident radio beam intensity and of the parameters of the scattering plasma do not affect the scattered component substantially.
This is believed to. be the case Lor the PR and. LEC o the Crab pulsar.
This is believed to be the case for the PR and LFC of the Crab pulsar.
Llowever. the variations of the SCALLCLEC component can be dramatic provided that P is somewha smaller and the scattering is about to reach the saturation stage.
However, the variations of the scattered component can be dramatic provided that $\Gamma$ is somewhat smaller and the scattering is about to reach the saturation stage.
Therefore we conclude that a number of pulsars may have the transient. components. which precede the SIP ane are analogous to the PR and LEC in the Crab pulsar.
Therefore we conclude that a number of pulsars may have the transient components, which precede the MP and are analogous to the PR and LFC in the Crab pulsar.
The presence of the transient components resulting from the higher-harmonic scattering is not excluded as well.
The presence of the transient components resulting from the higher-harmonic scattering is not excluded as well.
epoch of the Ilpparcos observation was 1901.25.
epoch of the Hipparcos observation was 1991.25.
The Ilipparcos proper motion of à —15.I5-Xx1.58.ó0 =1.62 1nas Ἐ was applied to correct tιο T Tai Nv» position to the epoch of our observation.
The Hipparcos proper motion of $\alpha=$ $\pm$ 1.88, $\delta=-12.48\pm1.62$ mas $^{-1}$ was applied to correct the T Tau N position to the epoch of our observation.
The saudard error in th T Tau NS position at the epoch of the observations is then An -]ίX.IE unas. AÓ =!l.l amas aud djs dominaed 1w the error arising from the proper motion correction.
The standard error in the T Tau N position at the epoch of the observations is then $\Delta \alpha=$ 16.4 mas, $\Delta\delta=$ 14.1 mas, and is dominated by the error arising from the proper motion correction.
The IR observatiois of Duchenne et al. (
The IR observations of Duchênne et al. (
2002) eive positions for Sa and Sb as offsets from the IR position of T Tau N at epoch 2000.2.
2002) give positions for Sa and Sb as offsets from the IR position of T Tau N at epoch 2000.2.
We have used these offsets to calculate the positions for Sa and Sb which are giveu iu Table 1..
We have used these offsets to calculate the positions for Sa and Sb which are given in Table \ref{positions}.
The orbital motion of T Tau S about T Tau N was estimated by Roddier et al. (
The orbital motion of T Tau S about T Tau N was estimated by Roddier et al. (
2000) who found. that the positiou aicle of T Tau S relative to T Tau N chauged by 0ου
2000) who found that the position angle of T Tau S relative to T Tau N changed by $0^{\circ}.66$ $^{-1}$.
This translates into 8 mas per vear ib a wesward direction.
This translates into 8 mas per year in a westward direction.
Since Sa is the bright«st component iu 1ο IR.(by about 2.5 maguitudes in IK). we asse that fus represeuts motion of T Tau Sa arotnd T Tau N. althouch the motion of the photocentre could be iuflueiced. by he motion of T Tau Sb around T Tau Sa.
Since Sa is the brightest component in the IR (by about 2.5 magnitudes in K), we assume that this represents motion of T Tau Sa around T Tau N, although the motion of the photocentre could be influenced by the motion of T Tau Sb around T Tau Sa.
Since any motion between t16 epoch of our observation and that of Ducheune et al.
Since any motion between the epoch of our observation and that of Duchênne et al.
would be small. we have not attempted to account for it in the position of Sa.
would be small, we have not attempted to account for it in the position of Sa.
This asstuuption is also made by Johuston et al. axd their ¢jbservatious are consistent with orbital motion of T Tau Sb aro:ud a stationary (with respect to T Tau N) T Tau Sa.
This assumption is also made by Johnston et al, and their observations are consistent with orbital motion of T Tau Sb around a stationary (with respect to T Tau N) T Tau Sa.
The observation of the T Tau S system by WkWohhler et al. (
The observation of the T Tau S system by Köhhler et al. (
2000) is closest in time to our observation (epoch 2000.17 as opposed to 1999.96). aud so this is f position for T Tau Sb with which the VLBI source position should xO priwcipally compared.
2000) is closest in time to our observation (epoch 2000.17 as opposed to 1999.96), and so this is the position for T Tau Sb with which the VLBI source position should be principally compared.
Fiπα. in Table 1 we list the position of the soitLer1 2em SOTUECE «jbserved by Johuston et al.
Finally, in Table \ref{positions} we list the position of the southern 2cm source observed by Johnston et al.
at epoch 2001.0531. which has been proper motion corrected to epoch 1999.958 usine the proper motion reported for T Tau ο by t1050 althors.
at epoch 2001.0531, which has been proper motion corrected to epoch 1999.958 using the proper motion reported for T Tau S by these authors.
All the positions ave listed in Table 1 and are shown i Figure 1.. with the exception of T Tau N which lies ouside the field.
All the positions are listed in Table \ref{positions} and are shown in Figure \ref{posnfig}, with the exception of T Tau N which lies outside the field.
The T Tau Sa aud Sb positions are marked with crosses. he sizes of which iudicate the uncertainty i the positions which are dominated by the proper motion correction to the T Tau N position as previously noted.
The T Tau Sa and Sb positions are marked with crosses, the sizes of which indicate the uncertainty in the positions which are dominated by the proper motion correction to the T Tau N position as previously noted.
This component of the uncertainty is tlhe same for Sa and Sb. and so the uucertaimties shown for these sources are correlated.
This component of the uncertainty is the same for Sa and Sb, and so the uncertainties shown for these sources are correlated.
The VEDI source in our map is closest to the position of T Tau Sb and we uake this identification. although there is a significant offset from the Ióohlhler position.
The VLBI source in our map is closest to the position of T Tau Sb, and we make this identification, although there is a significant offset from the Köhhler position.
One wav this welt arise is because the T Tar N position is from the Vband Uipparcos observations. whereas he position of T Tau Sa relative to T Tau N bv Duchenne ct al.
One way this might arise is because the T Tau N position is from the -band Hipparcos observations, whereas the position of T Tau Sa relative to T Tau N by Duchênne et al.
was determined in the near infrared.
was determined in the near infrared.