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We also used visual and CCD observations from the AAVSO) International Database. | We also used visual and CCD observations from the AAVSO International Database. |
Where both Johnson D and V. data were available at the same epoch. we retained only the V maxinia. | Where both Johnson $B$ and $V$ data were available at the same epoch, we retained only the $V$ maxima. |
These data points are listed in the first column of ret'abOC.. | These data points are listed in the first column of \\ref{TabOC}. |
Column22 lists the epoch number (££=0 was arbitrarily taken at the first maximum of theAepfer data). | 2 lists the epoch number $E=0$ was arbitrarily taken at the first maximum of the data). |
Weights were assigned to individual data sets similarly to the case of V2279CCve. | Weights were assigned to individual data sets similarly to the case of Cyg. |
Phe weights are listed in 4. | The weights are listed in 4. |
The final ephemeris was derived by a weighted linear least squares fit to the preliminary residuals computed. from à formerly published (usually slightly incorrect) period. value which is an inherent step inthe OC' method. | The final ephemeris was derived by a weighted linear least squares fit to the preliminary residuals computed from a formerly published (usually slightly incorrect) period value which is an inherent step in the $O-C$ method. |
No weight was assigned to the photographie normal maxima. these ο6 residuals were omitted. from fitting procedure. | No weight was assigned to the photographic normal maxima, these $O-C$ residuals were omitted from fitting procedure. |
Phe ο.6 residuals in 33 have been calculated with the final ephemeris C=BID(2454955.726040.0008)n|(4.92545440.000001)LE. | The $O-C$ residuals in 3 have been calculated with the final ephemeris $C = {\rm BJD}\,(2\,454\,955.7260 \pm 0.0008) + (4.925454 \pm
0.000001)\,{\rm E}$. |
This period is considered to be more accurate than the one obtained by fitting the frequency. ane its harmonics. | This period is considered to be more accurate than the one obtained by fitting the frequency and its harmonics. |
The source of data is given in the last column of refTabOC.. | The source of data is given in the last column of \\ref{TabOC}. |
The οC' diagram is shown in Πο ο | The $O-C$ diagram is shown in \\ref{FigOC}. |
This plot indicates that the pulsation period has been constant during the last vy. In. principle. the moments of maximum light diller for different: passbands. | This plot indicates that the pulsation period has been constant during the last y. In principle, the moments of maximum light differ for different passbands. |
From simultaneous observations we [ind a cillerence of 0.020 dd. between the V. and. Ap maxima in the sense of VoAp. which is neglected in refFigOC.. | From simultaneous observations we find a difference of $-0.020$ d between the $V$ and $Kp$ maxima in the sense of $V - Kp$, which is neglected in \\ref{FigOC}. |
We note that this does not change our conclusions in any wer. | We note that this does not change our conclusions in any way. |
The suspected. spectroscopic. binarity of V1I54€C€veg can be investigated with the help of the new radial velocity data. most of them obtained with the Tautenburg 2.0 m telescope (see reftabspec)). | The suspected spectroscopic binarity of Cyg can be investigated with the help of the new radial velocity data, most of them obtained with the Tautenburg 2.0 m telescope (see \\ref{tabspec}) ). |
The comparison with the data taken from the literature (Gorvnyaetal.1998). and Imbert(1999). does not give a new evidence of spectroscopic binarity because the 5-velocitv (mean radial velocity averaged: over a complete pulsational evele) derived. from the new data. practically coincides with that obtained from the previous observations see refFie-RVeomp.. | The comparison with the data taken from the literature \citep{goretal98}
and \citet{imb99} does not give a new evidence of spectroscopic binarity because the $\gamma$ -velocity (mean radial velocity averaged over a complete pulsational cycle) derived from the new data practically coincides with that obtained from the previous observations – see \\ref{Fig-RVcomp}. |
Quantitativelv. the dillerence is 0.19+O.30kkmss ! and |0.15+£0.22 kkmiss1 between our new data and data of Gorvnyaetal.(1998). and Embert(1999).. respectively. | Quantitatively, the difference is $-0.19 \pm 0.30$ $^{-1}$ and $-0.15 \pm0.22$ $^{-1}$ between our new data and data of \citet{goretal98} and \citet{imb99}, respectively. |
We add that a period. derived from. all the available racial velocity data is in very good agreement with (less than 2-0 from) the photometric period derived in the previous subsection. | We add that a period derived from all the available radial velocity data is in very good agreement with (less than $\sigma$ from) the photometric period derived in the previous subsection. |
Copheids pulsate in one of the first three radial modes (funclamental (E). first COL) and second overtone (02)) or simultaneously in two or three of them. | Cepheids pulsate in one of the first three radial modes (fundamental (F), first (O1) and second overtone (O2)) or simultaneously in two or three of them. |
Some triple-niode Cepheids pulsate in the first three overtones at the same time. | Some triple-mode Cepheids pulsate in the first three overtones at the same time. |
From a pulsational ancl evolutionary. point. of view it is important to determine the pulsational mode. of a monoperiodic Cepheid. | From a pulsational and evolutionary point of view it is important to determine the pulsational mode of a monoperiodic Cepheid. |
Cepheics with pulsational periods similar to VII54CCvg may pulsate in the funcamental or the first overtone mode. | Cepheids with pulsational periods similar to Cyg may pulsate in the fundamental or the first overtone mode. |
The usual way of distinction is the use of Fourier parameters that. show characteristic »ogressionas a function of period. | The usual way of distinction is the use of Fourier parameters that show characteristic progressionas a function of period. |
Llowever. the Fourier xweanmeters of radial velocity curves are indistinguishable or fundamental ancl first overtone pulsators with periods around Selel (see 33 of Baranowskietal. 2009)). | However, the Fourier parameters of radial velocity curves are indistinguishable for fundamental and first overtone pulsators with periods around d (see 3 of \citealt{bsd09}) ). |
Light curve Fourier parameters suller from. similar problems. | Light curve Fourier parameters suffer from similar problems. |
Dased on reffps.. A», and Os, are the most promising parameters or mode discrimination. | Based on \\ref{fps}, , $R_{21}$ and $\phi_{31}$ are the most promising parameters for mode discrimination. |
However. à closer inspection reveals that there is a 2x dilference between| νε values of | However, a closer inspection reveals that there is a $2\pi$ difference between $\phi_{31}$ values of |
Following the logic of Sheltonetal.(2007).. the rims deviation between the mean of the pareut population and the mean of a sample of four observations drawn randomly from the parent population should be q/z/l times as laree as the average deviation between 2 sight lines. thus 114. aud the intrinsic fluctuation between cohunn deusities alone different directious should be a factor \/7/2 times the average deviation between 2 sight lues. thus 214. | Following the logic of \citet{shelton_etal_07}, the rms deviation between the mean of the parent population and the mean of a sample of four observations drawn randomly from the parent population should be $\sqrt{\pi}/4$ times as large as the average deviation between 2 sight lines, thus $11\%$, and the intrinsic fluctuation between column densities along different directions should be a factor $\sqrt{\pi}/2$ times the average deviation between 2 sight lines, thus $21\%$. |
Combining these terius in quadrature vields 21%. | Combining these terms in quadrature yields $24\%$. |
This is larger than the average statistical plus continui placement error and the average svsteniatic error obtained from the four individual column density measurements. | This is larger than the average statistical plus continuum placement error and the average systematic error obtained from the four individual column density measurements. |
Tere. we add the 3 sources of uncertaiuty in quadrature in order to obtain an estimate of the uncertaüntv im the estimated cobpuun density for our direction. | Here, we add the 3 sources of uncertainty in quadrature in order to obtain an estimate of the uncertainty in the estimated column density for our direction. |
Thecobluuu density and combined error are 197022«101 2. | Thecolumn density and combined error are $1.97^{+0.53}_{-0.51}\times 10^{14}$ $^{-2}$. |
The column cdeusitv of local iuaterial mist he subtracted from the Lue of sight ουν density. | The column density of local material must be subtracted from the line of sight column density. |
The Savage&Lehner(2006). survey of ttoward nearby white dwarf iuclides 9 sight lines that are within 20° of our off-cloud direction. | The \citet{savage_lehner_06}
survey of toward nearby white dwarfs includes 9 sight lines that are within $\degr$ of our off-cloud direction. |
Tous of wwere not detected ou 1 of these sight lines. | Ions of were not detected on 4 of these sight lines. |
Tn those cases. Savage&Lehner(2006) listed 20 upper linüts. | In those cases, \citet{savage_lehner_06} listed $2\sigma$ upper limits. |
We treat these cases as if the lo upper linits were half of the 20 upper limits aud the detected values were 0. so that we can fud the average for the 9 sight lines. | We treat these cases as if the $1\sigma$ upper limits were half of the $2\sigma$ upper limits and the detected values were 0, so that we can find the average for the 9 sight lines. |
The resulting average volume deusity of iis 1.5!rat«10 Sem7. | The resulting average volume density of is $1.5^{+1.2}_{-0.4} \times 10^{-8}$ $^{-3}$. |
Although there are reports (i.e. Barstowctal. 20093) claiming that only a few of the previous ddetectious for nearby sight lines can be shown to be interstellar. as opposed to plotospheric or ambiguous (generally due to the simibuities between the observed volocities and the velocities of photosphieric lines in the spectra). a culled dataset is not vet available. | Although there are reports (i.e., \citealt{barstow_etal_09}) ) claiming that only a few of the previous detections for nearby sight lines can be shown to be interstellar, as opposed to photospheric or ambiguous (generally due to the similarities between the observed velocities and the velocities of photospheric lines in the spectra), a culled dataset is not yet available. |
Thus. we use the Savage&Lehner(2006) dataset. | Thus, we use the \citet{savage_lehner_06}
dataset. |
We estimate the maximum extent of the Local Bubble from a survey of the Local Cavitv's wall. iu this case the aud sstvey of Welshetal.(2010). | We estimate the maximum extent of the Local Bubble from a survey of the Local Cavity's wall, in this case the and survey of \citet{welsh_etal_10}. |
. In Figures Ll and 17 of that paper. the Local Cavity extends 75 pc in the direction of our observation. | In Figures 14 and 17 of that paper, the Local Cavity extends 75 pc in the direction of our observation. |
Taking this aud the average volue deusity vields au cstimate of the Local Bubble’s coluun density of of BLISS«10 7. | Taking this and the average volume density yields an estimate of the Local Bubble's column density of of $3.4^{+2.8}_{-1.0} \times 10^{12}$ $^{-2}$. |
This is so small compared with the full coluuu denusitv through the Calactic disk aud halo that approximations are consequential. | This is so small compared with the full column density through the Galactic disk and halo that approximations are inconsequential. |
Subtracting this value from the full column density vields a halo οσο deusity of Novy,=19E05«10! 7, | Subtracting this value from the full column density yields a halo column density of $N_{OVI} = 1.9 \pm 0.5 \times10^{14}$ $^{-2}$. |
Again. the unrounded measurement values are used iu the calculations. but the rounded results are reported. | Again, the unrounded measurement values are used in the calculations, but the rounded results are reported. |
There are 33 archival oobservations. within⋅⋅ oot our ppointings that have at least 016 EPIC-MOS exposure. | There are 33 archival observations within of our pointings that have at least some EPIC-MOS exposure. |
All but 3 of the observations were badly affected by soft proton contamination. | All but 3 of the observations were badly affected by soft proton contamination. |
We have processed the data from these 3 observations using the method described m Henley&Shelton(2010). | We have processed the data from these 3 observations using the method described in \citet{henley_shelton_10}. |
. We removed times affected bv soft proton flares. and also times when the solar wiud proton flux exceeded 2.«105 cin2 1 (the latter step was to reduce contamination from solar wind charge exchange N-ravs: sce Section 2.1.2.below). | We removed times affected by soft proton flares, and also times when the solar wind proton flux exceeded $2 \times 10^8$ $^{-2}$ $^{-1}$ (the latter step was to reduce contamination from solar wind charge exchange X-rays; see Section 2.4.2,below). |
We extracted spectra from the blank sky regions of the EPIC-MOS chips. and measured the intensities iu the ο VII Ka riplet (569-571eV) and O VIII Ίνα line. accounting or the effects of residual soft proton contamination aud he extragalactic backeround (the spectrum of which we assumed to be 10.56(£/keV) $19 ph ὃς by 3 τον1: Chenetal. 1997)). | We extracted spectra from the blank sky regions of the EPIC-MOS chips, and measured the intensities in the O VII $\alpha$ triplet (569-574eV) and O VIII $\alpha$ line, accounting for the effects of residual soft proton contamination and the extragalactic background (the spectrum of which we assumed to be $E$ $^{-1.46}$ ph $^{-2}$ $^{-1}$ $^{-1}$ $^{-1}$; \citealt{chen_etal_97}) ). |
For the of portionsobservation POOTSOGOL (6=90.07.60 38.07) not affected. by soft xoton fares. the solar wind proton flux exceeded our standard filtering threshold of2«105 ὃς 1. | For the portions of observation 0200750601 $\ell = 90.0\degr, b = 38.4\degr$ ) not affected by soft proton flares, the solar wind proton flux exceeded our standard filtering threshold of $2 \times 10^8$ $^{-2}$ $^{-1}$ . |
However, he proton flux was not unusuallv large. aud did not exceed So2.8&4s105 2 ft. | However, the proton flux was not unusually large, and did not exceed $2.8 \times 10^8$ $^{-2}$ $^{-1}$ . |
We- therefore. did: uot cary out proton flux filteriug ou this observation. | We therefore did not carry out proton flux filtering on this observation. |
The N-rav iutensitv measurements are listed in Table [.. | The X-ray intensity measurements are listed in Table \ref{table:ovii_oviii}. . |
in L1512 are similar to those previously observed iu other low-iass star-forming cores. LL1521F and LISLLE107? cin.2 and «107? 7. respectively: 2009)). and L1527 (Sslol! ?: | in L1512 are similar to those previously observed in other low-mass star-forming cores, L1521F and L1544$\times10^{10}$ $^{-2}$ and $\times10^{10}$ $^{-2}$, respectively; ), and L1527 $\times10^{10}$ $^{-2}$; ). |
These values are somewhat less than those observed in the quicsceut molecular clouds TMC-1 «104 7: 2007)) aud Lupusz-lAÀ «1019 2: 2010)). | These values are somewhat less than those observed in the quiescent molecular clouds TMC-1 $\times10^{11}$ $^{-2}$; ) and Lupus-1A $\times10^{10}$ $^{-2}$; ). |
The aniou-to-neutral ratios ([CGIT σου. however. are approximately in the nmüddle of the previously observed distribution: 3.6+1.3% for L1251À and L241.1% for L1512. compared with for TMC-1. for Lupus-1À. for LI5LL for L1521F and for L1527. | The anion-to-neutral ratios $_6$ $^-]/[$ $_6$ H]), however, are approximately in the middle of the previously observed distribution: $3.6\pm1.3$ for L1251A and $4.2\pm1.4$ for L1512, compared with for TMC-1, for Lupus-1A, for L1544, for L1521F and for L1527. |
hypothesised that an inverse relationship between ILatoui abundance and eas density would result iu greater ΕΕ ratios iu denser gas. which may explain the larecr ΓΗratio found in L1527 where the lvdrogen nucleon deusitv sg; is 109 7. compared tonyc105 ? in TMC-1. | hypothesised that an inverse relationship between H-atom abundance and gas density would result in greater $_6$ $^-]/[$ $_6$ H] ratios in denser gas, which may explain the larger ratio found in L1527 where the hydrogen nucleon density $n_H$ is $\sim10^6$ $^{-3}$, compared to $n_H\sim10^4$ $^{-3}$ in TMC-1. |
L1251À and L1512 fit this trend both have ny7LO? ?2005).. and their anion-to-neutral ratios are intermediate between TMC-1 aud L1527. | L1251A and L1512 fit this trend – both have $n_H\sim10^5$ $^{-3}$, and their anion-to-neutral ratios are intermediate between TMC-1 and L1527. |
The detection of Cell iu L1251À is the first reported interstellar anion m a protostar outside of Taurus. aud thus shows that the tuportance of anions mst be considered iu future studies of the chemistry of star forming regions throughout the Calaxy. | The detection of $_6$ $^-$ in L1251A is the first reported interstellar anion in a protostar outside of Taurus, and thus shows that the importance of anions must be considered in future studies of the chemistry of star forming regions throughout the Galaxy. |
L1251À aud L1512 have similar [ΟΠ Hf] ratios of and aud0.5%.. respectively. | L1251A and L1512 have similar $_6$ $_4$ H] ratios of and and, respectively. |
These are within the range of values previously measured in iuterstellar clouds by(2009). | These are within the range of values previously measured in interstellar clouds by. |
. The L1251À molecular cloud is located in the Ceplieus Flare region of low-toduteriiediate mass star formation2008). | The L1251A molecular cloud is located in the Cepheus Flare region of low-to-intermediate mass star formation. |
. Our observed position is 10" SE of the embedded Class 0 protostar L1251À IRS2. which powers a molecular outflow2010). | Our observed position is $40''$ SE of the embedded Class 0 protostar L1251A IRS3, which powers a molecular outflow. |
.. The protostar ceutre is located outside of the 26" CBT beuu (see Figure 1)). so ciission from its core does not directly influence our observations. | The protostar centre is located outside of the $26''$ GBT beam (see Figure \ref{fig:map}) ), so emission from its core does not directly influence our observations. |
However. depending on the radius of ifs outer boundary (which is likely to be up to a few times 10! AU: see. citealtjor02)). eas frou inside the protostar cuvelope is xobablv responsible for πιο] of the observed molecular cluission. | However, depending on the radius of its outer boundary (which is likely to be up to a few times $10^4$ AU; see, ), gas from inside the protostar envelope is probably responsible for much of the observed molecular emission. |
It would be of interest to observe the carbon chain (aud anion) emission closer to the protostellar core. o determine whether the elevated tempcratures there lave any impact on the abuudauces of these nispecies. as was been Lypothesised by (2008b). | It would be of interest to observe the carbon chain (and anion) emission closer to the protostellar core, to determine whether the elevated temperatures there have any impact on the abundances of these species, as has been hypothesised by . |
. Although previous observations have shown relatively arge aud abundances iu various parts of the Cepheus couples. our observations constitute the larecst reported CIT column density aud the first detection of | Although previous observations have shown relatively large $_3$ N and $_5$ N abundances in various parts of the Cepheus complex, our observations constitute the largest reported $_4$ H column density and the first detection of |
to other telescope pointing positious. | to other telescope pointing positions. |
The phase solution is sullicieutly good to permit application ol oue phase solution to another poiuting. separated by many degrees aud considerable time. but the amplitude calibration does not transfer well. being in error by several percent. | The phase solution is sufficiently good to permit application of one phase solution to another pointing, separated by many degrees and considerable time, but the amplitude calibration does not transfer well, being in error by several percent. |
The longest 32'T baseline is 30010 aud siuce we expect iouosplieric displacements to be small compared to this relatively large 327T. pixel size (Erickson. we do uot apply au lonospheric solution that is a function of position within the field of view. | The longest 32T baseline is $\sim 300$ m and since we expect ionospheric displacements to be small compared to this relatively large 32T pixel size (Erickson, \nocite{eri84} we do not apply an ionospheric solution that is a function of position within the field of view. |
However. if a field coutaius multiple calibrators they each have an iudepeudeut gain and 10nospheric fit. | However, if a field contains multiple calibrators they each have an independent gain and ionospheric fit. |
By way of introduction. the imaging transformation performed by an interferometer cau be described by the vau Cittert-Zernike equation (see Clark 1973. — and Thompsou. Moran aud Swenson. 2001. for a detailed treatinent). | By way of introduction, the imaging transformation performed by an interferometer can be described by the van Cittert-Zernike equation (see Clark 1973, \nocite{cla73} and Thompson, Moran and Swenson, 2001, \nocite{tms} for a detailed treatment). |
Here the te aud w are the spacing (the distance between autenua pairs) irieasured in wavelenetlis. auc fon aud wv are direction cosines of the brightuess distribution with respect to this coordiuate [rame (i=V1-—1/P— 112). | Here the $u, v$ and $w$ are the spacing (the distance between antenna pairs) measured in wavelengths, and $l,m$ and $n$ are direction cosines of the brightness distribution with respect to this coordinate frame $n=\sqrt{1-l^{2} - m^{2}}$ ). |
The coordinate system is defined such that the w-asxis points in the direction of the phase centre. | The coordinate system is defined such that the $w$ -axis points in the direction of the phase centre. |
The van Cittert-Zernike equation indicates that the complex-valued visibility function. WO.ew). is a Fourier-like iutegral of the sky briglitness. Z(6. 7). multiplied by the primary beam response of an interferometer. A(f.i). aud L/v. | The van Cittert-Zernike equation indicates that the complex-valued visibility function, $V(u,v,w)$, is a Fourier-like integral of the sky brightness, $I(\ell,m)$ , multiplied by the primary beam response of an interferometer, $A(\ell,m)$, and $1/n$. |
I— standardimaging it has been conventional to make the van Cittert-Zernike equation indepeucent ol n and το. aud reduce it to a two-dimeusional Fourier trausform. | In imaging it has been conventional to make the van Cittert-Zernike equation independent of $n$ and $w$, and reduce it to a two-dimensional Fourier transform. |
This is conventionally achieved by the thefteld approximation (izz 1). | This is conventionally achieved by the the approximation $n \approx 1$ ). |
Which results iu (η—1)zz0 and a projection from the celestial sphere outo a plaue linear iu { and i. preseuting au orthographic projection of tle sky. | Which results in $(n-1)w \approx 0$ and a projection from the celestial sphere onto a plane linear in $l$ and $m$, presenting an orthographic projection of the sky. |
However this approximation effectively ueglects το. implying that the direction of the projection is always perpeucictlar to the f.η plane. aud uot perpeudicular to the taugeut plaue. | However this approximation effectively neglects $w$, implying that the direction of the projection is always perpendicular to the $\ell,m$ plane, and not perpendicular to the tangent plane. |
This mauilests as a phase error that can ouly be kept to a tolerable level by ensuriug the small field approximation holds. thus restricting the size of the field of view. | This manifests as a phase error that can only be kept to a tolerable level by ensuring the small field approximation holds, thus restricting the size of the field of view. |
The phase error to first order is given by: The small lield approximation doesnothold for the MWA aud. therefore a mechanisin must be | The phase error to first order is given by: The small field approximation doesnothold for the MWA and therefore a mechanism must be |
One of the goals of the asteroseismology programme (Michel et 22006) of the CoRoT satellite (Baglin et 220006) is to explore the Hertzsprung-Russell diagram (HRD) through uninterrupted time series of white-light photometry of unprecedented precision. | One of the goals of the asteroseismology programme (Michel et 2006) of the CoRoT satellite (Baglin et 2006) is to explore the Hertzsprung-Russell diagram (HRD) through uninterrupted time series of white-light photometry of unprecedented precision. |
In this context. numerous non-radial pulsators of various kind have been observed and analysed. among which massive stars on the main sequence (e.g. Degroote et 22009: Neiner et 22009). | In this context, numerous non-radial pulsators of various kind have been observed and analysed, among which massive stars on the main sequence (e.g., Degroote et 2009; Neiner et 2009). |
With the goals of mapping the uppermost part of the HRD and understanding the role of oscillations in the mass loss of evolved massive stars. a hot supergiant was observed in the seismology programme of the satellite. | With the goals of mapping the uppermost part of the HRD and understanding the role of oscillations in the mass loss of evolved massive stars, a hot supergiant was observed in the seismology programme of the satellite. |
The B-type supergiant 550064 (V. mag of 8.21) has not been studied in detail. | The B-type supergiant 50064 $V$ mag of 8.21) has not been studied in detail. |
Its spectral type assignments range from Bila (Jacoby Hunter 1984) to Bola (Blanco et 11970). | Its spectral type assignments range from B1Ia (Jacoby Hunter 1984) to B6Ia (Blanco et 1970). |
It was monitored by CoRoT. whose performance not Cftnly delivers two orders of magnitude better precision than anye& ground-based photometry. but. even more importantly for supergiant stars. also guarantees uninterrupted data curing several months. | It was monitored by CoRoT, whose performance not only delivers two orders of magnitude better precision than any ground-based photometry, but, even more importantly for supergiant stars, also guarantees uninterrupted data during several months. |
This is essential for progress in understanding massive evolved stars. because ground-based onata for supergiants have so far suffered severely from very low Cuty cycles. | This is essential for progress in understanding massive evolved stars, because ground-based data for supergiants have so far suffered severely from very low duty cycles. |
The oscillations of evolved massive stars known so far essentially come in two faveours. | The oscillations of evolved massive stars known so far essentially come in two flavours. |
Classical. gravity mode oscillations. with periods of a few days. excited by the x mechanism. have recently beer found from space photometry (Sato et 22006; Lefever et 220074). | Classical gravity mode oscillations with periods of a few days, excited by the $\kappa\,$ mechanism, have recently been found from space photometry (Saio et 2006; Lefever et 2007a). |
On the other hand. theory predicts so-called strange modes with periods between roughly 10 and dd in stars with masses above MM... | On the other hand, theory predicts so-called strange modes with periods between roughly 10 and d in stars with masses above $_\odot$. |
These strange modes. which can be both radial and non-radial in nature. are modes trapped in a cavity caused by a density inversion in the very outer. highly non-adiabatic envelope of stars with a high L/M ratio whose radiation pressure dominates over the gas pressure (Glatzel Kiriakidis 1993; Sato et 11998; Glatzel et 11999: Dorfi Gautschy 2000). | These strange modes, which can be both radial and non-radial in nature, are modes trapped in a cavity caused by a density inversion in the very outer, highly non-adiabatic envelope of stars with a high L/M ratio whose radiation pressure dominates over the gas pressure (Glatzel Kiriakidis 1993; Saio et 1998; Glatzel et 1999; Dorfi Gautschy 2000). |
This type of oscillation has been claimed to be responsible for the mass-loss episodes of luminous stars. such as luminous blue variables (LBVs). GGlatzel Kirtakidis (1993). but observational proof of the occurrence of strange modes has not been established so far. | This type of oscillation has been claimed to be responsible for the mass-loss episodes of luminous stars, such as luminous blue variables (LBVs), Glatzel Kiriakidis (1993), but observational proof of the occurrence of strange modes has not been established so far. |
Our data of 550068 550064 was observed by CoRoT during a long run in the anticentre direction (LRaOI) for 136.9 days. | Our data of 50064 50064 was observed by CoRoT during a long run in the anticentre direction (LRa01) for 136.9 days. |
It is the only hot supergiant among the seismology targets so far. | It is the only hot supergiant among the seismology targets so far. |
The CoRoT light curve contains 3199913 datapoints. with a time sampling of 32ss. after deleting the measurements suffering from hot pixels during the passage through the South Atlantic Anomaly. | The CoRoT light curve contains 913 datapoints, with a time sampling of s, after deleting the measurements suffering from hot pixels during the passage through the South Atlantic Anomaly. |
In order to compare the space photometry behaviour of 550064 with the one reported in the literature (Halbedel 1990). we transferred the CoRoT fluxes to. magnitudes. | In order to compare the space photometry behaviour of 50064 with the one reported in the literature (Halbedel 1990), we transferred the CoRoT fluxes to magnitudes. |
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