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Very recently we detected the fifth outburst of KL Dra (2010 Apr Ist) which took place at least 54 cavs after the start of the previous outburst. | Very recently we detected the fifth outburst of KL Dra (2010 Apr 1st) which took place at least 54 days after the start of the previous outburst. |
Although we find no evidence [or other periods in our optical data à more comprehensive dataset is needed to test this more thoroughly. | Although we find no evidence for other periods in our optical data a more comprehensive dataset is needed to test this more thoroughly. |
During its low optical state. KWL Dra appears rather variable which is probably associated with its orbital period (Wood et. al. | During its low optical state, KL Dra appears rather variable which is probably associated with its orbital period (Wood et al. |
2002). | 2002). |
Lt also shows what appears to. be short duration brightenings (up to ~1.5 mags) on several occasions (ie M.JD55092). | It also shows what appears to be short duration brightenings (up to $\sim$ 1.5 mags) on several occasions (ie $\sim$ 55092). |
We were fortunate to be observing on the 4.2m William | We were fortunate to be observing on the 4.2m William |
Millisecond pulsars are commonly believed to be descendants of normal neutron stars that have been spun-up and recycled back as radio pulsars by acquiring angular momentum from their companion during the low-mass X-ray binary (LMXB) phase 1982).. | Millisecond pulsars are commonly believed to be descendants of normal neutron stars that have been spun-up and recycled back as radio pulsars by acquiring angular momentum from their companion during the low-mass X-ray binary (LMXB) phase \citep{ACR82,RS82}. |
There are about 20 high confidence nuclear or accretion powered (see Table 1) millisecond pulsars (MSXPs) which are thought to be the progenitors of millisecond radio pulsars (MSRPs) &vander(WijnandsKlis[1998).. | There are about $\sim$ 20 high confidence nuclear or accretion powered (see Table 1) millisecond X-ray pulsars (MSXPs) which are thought to be the progenitors of millisecond radio pulsars (MSRPs) \citep{WK98}. |
These MSXPs may become observable in radio wavelengths once accretion ceases, or the column density of the plasma from the fossil disk around the neutron star becomes thin enough to allow vacuum gap formation that leads to the production of coherent | These MSXPs may become observable in radio wavelengths once accretion ceases, or the column density of the plasma from the fossil disk around the neutron star becomes thin enough to allow vacuum gap formation that leads to the production of coherent |
remains valid where is (he average enerey per reflection and is the rate of reflections. such that and A. is the turbulent. inner scale (associated. with the parallel component of the magnetic fiekl gradient lor anisotropic turbulence). vy is the effective relative speed of the electrons and compressions within the plasma. 12, is the electron mass. and e is (he spectral index of the turbulent cascade It is assumed throughout our analvsis (hat at length scales comparable to L the strength of the magnetic fluctuations. 0/3 is comparable to that of the mean magnetic field. 2. | remains valid where is the average energy per reflection and is the rate of reflections, such that and $\lambda_s$ is the turbulent inner scale (associated with the parallel component of the magnetic field gradient for anisotropic turbulence), $v_d$ is the effective relative speed of the electrons and compressions within the plasma, $m_e$ is the electron mass, and $a$ is the spectral index of the turbulent cascade It is assumed throughout our analysis that at length scales comparable to L the strength of the magnetic fluctuations, $\delta B$ is comparable to that of the mean magnetic field, $B$. |
The magnetic spectrum of turbulence in the ISM is difficult to measure. | The magnetic spectrum of turbulence in the ISM is difficult to measure. |
Between LOMpe and 0.03pc the specirum appears to be close to Kolmogorov (a=3) (ancl flatter on larger scales) (Ilanetal.(2004):Minter&Spanegler (1996))). | Between $10^{19}{\rm pc}$ and $0.03$ pc the spectrum appears to be close to Kolmogorov $(a=3)$ (and flatter on larger scales) \cite{hfm04,ms96}) ). |
There is no direct measure of (he magnetic spectrum on scales below 0.03pc. | There is no direct measure of the magnetic spectrum on scales below $0.03$ pc. |
For the range of magnetic spectra 2<a<x the result. from our calculations (o follow. that the cascade truncation scale well exceeds the smallest scintillation scale. will not change. | For the range of magnetic spectra $2 \le a \le \infty$ the result from our calculations to follow, that the cascade truncation scale well exceeds the smallest scintillation scale, will not change. |
However. we first consider a Goldreich-Sridhar (hereafter GS) turbulent spectrum with «=2 (GoldreichandSridhar1997). (see also Matthaeusetal. (1993))). | However, we first consider a Goldreich-Sridhar (hereafter GS) turbulent spectrum with $a=2$ \citep{GS} (see also \cite{m98}) ). |
Like etal. (1998).. GS turbulence is fully anisotropic. and is based on an incompressible cascade. | Like \cite{m98}, GS turbulence is fully anisotropic, and is based on an incompressible cascade. |
It involves a more rapid cascade in the direction. perpendicular to the local mean field than parallel to it. | It involves a more rapid cascade in the direction perpendicular to the local mean field than parallel to it. |
Turbulence becomes less aud less compressible on smaller scales. aud since (he Revnolds laver seems to be modestly magnetically dominated | Turbulence becomes less and less compressible on smaller scales, and since the Reynolds layer seems to be modestly magnetically dominated |
compact companion which just left the AGB. as suggested by the presence of cool dust and. often. a planetary nebula. | compact companion which just left the AGB, as suggested by the presence of cool dust and, often, a planetary nebula. |
The fast rotation of the giant is very likely the signature of a recent mass-transfer episode (??).. | The fast rotation of the giant is very likely the signature of a recent mass-transfer episode \citep{Jorissen-Zacs-2005,Zamanov-2006}. |
Actually. the shape of the ddiagram of post-AGB stars. when compared to that of binary M and K giants. may be interpreted in two different ways. | Actually, the shape of the diagram of post-AGB stars, when compared to that of binary M and K giants, may be interpreted in two different ways. |
The first interpretation assumes that the location of post-AGB stars in the (e.logP) diagram ts controlled by the periastron envelope corresponding to 200 ((solid line in Fig. 2). | The first interpretation assumes that the location of post-AGB stars in the $(e - \log P)$ diagram is controlled by the periastron envelope corresponding to 200 (solid line in Fig. \ref{Fig:elogP_panels}) ), |
with some circular systems at shorter periods some among these which had their eccentricities pumped upwards as a result of the interaction with. the circumbinary disc (222).. | with some circular systems at shorter periods some among these which had their eccentricities pumped upwards as a result of the interaction with the circumbinary disc \citep{Artymowicz91,DeRuyter-2006,Frankowski-2008:b}. |
In this scenario. the AGB precursors were thus allowed to evolve far up the AGB. up to at least 200FR. | In this scenario, the AGB precursors were thus allowed to evolve far up the AGB, up to at least 200. |
.. The similarity between the ddiagrams of M giants and post-AGB systems would thus be purely accidental. | The similarity between the diagrams of M giants and post-AGB systems would thus be purely accidental. |
The second interpretation assumes that this similarity is not accidental. and states that mass transfer or dise interaction has not dramatically altered the location of post-AGB binaries in the (¢logP) diagram (or as suggested above. some might even be of the mass-transfer process). | The second interpretation assumes that this similarity is not accidental, and states that mass transfer or disc interaction has not dramatically altered the location of post-AGB binaries in the $(e - \log P)$ diagram (or as suggested above, some might even be of the mass-transfer process). |
This case then differs from the first interpretation by implying that the AGB precursors must have left the AGB at a rather early stage in their evolution. given their location in the («logο.P) diagram between the tidal envelopes corresponding to 85 and 200 (Fig. 1). | This case then differs from the first interpretation by implying that the AGB precursors must have left the AGB at a rather early stage in their evolution, given their location in the $(e - \log P)$ diagram between the tidal envelopes corresponding to 85 and 200 (Fig. \ref{Fig:elogP_M}) ). |
Like the M giants. post-AGB stars do not show evidence for s-process enrichment (despitetheconfusionntroducedbythedepletionpattern: ?).. | Like the M giants, post-AGB stars do not show evidence for s-process enrichment \citep[despite the confusion introduced by the depletion
pattern;][]{VanWinckel-2007}. |
We believe that the absence of s-process enrichments in both classes of stars is not a coincidence. | We believe that the absence of s-process enrichments in both classes of stars is not a coincidence. |
In fact. it may be inferred that neither M giants or post-AGB binaries reached a point on the AGB where s-process and dredge-ups wereoperating?. | In fact, it may be inferred that neither M giants nor post-AGB binaries reached a point on the AGB where s-process and dredge-ups were. |
. This argument may actually be checked by locating M giants and post-AGB stars 1 the Hertzsprung-Russell (HR) diagram (Figs. | This argument may actually be checked by locating M giants and post-AGB stars in the Hertzsprung-Russell (HR) diagram (Figs. |
5 and 6)) and comparing their location with that of Te-rich S stars (from?).. which are thermally-pulsing AGB (TP-AGB) stars enriched in s-process elements. | \ref{Fig:HR} and \ref{Fig:HR1}) ) and comparing their location with that of Tc-rich S stars \citep[from][]{VanEck-98}, which are thermally-pulsing AGB (TP-AGB) stars enriched in s-process elements. |
Luminosities of post-AGB stars are extremely scarce. | Luminosities of post-AGB stars are extremely scarce. |
When available (e.g..fromLMCobjectsandbratedonLMCstars:??) however. they are of the order of 3000 .. dogL/L..= 3.6). | When available \citep[e.g., from LMC objects and from a
luminosity-period relationship for RV Tau stars calibrated on LMC
stars;][]{Alcock-1998,DeRuyter-2006} however, they are of the order of 4000 $_\odot$ $\log L/{\rm L}_\odot = 3.6$ ). |
This luminosity lies in the middle of the range covered by Te-rich S stars (Fig. 5)). | This luminosity lies in the middle of the range covered by Tc-rich S stars (Fig. \ref{Fig:HR}) ), |
so that this argument does not really support a non-TP-AGB origin for | so that this argument does not really support a non-TP-AGB origin for |
another diffuse source. which is mareinally detected in the WEPC2 tage. | another diffuse source, which is marginally detected in the WFPC2 image. |
We show the |OIII]J ABOUT. cinission line morphology of QOBL7-383 C5 in Figure 1. | We show the $\lambda$ 5007 emission line morphology of Q0347-383 C5 in Figure \ref{fig:o3im}. |
The image inchicdes waveleugths +1 FWIHAL around the peak iutegrated cluission. | The image includes wavelengths $\pm 1$ FWHM around the peak integrated emission. |
The cussion is clearly spatially extended. over au area of ~ 0.6.1.3". correspouding to ~L5 9.8 kpc. | The emission is clearly spatially extended over an area of $\sim 0.6$ $\times1.3$, corresponding to $\sim 4.5$ $\times9.8$ kpc. |
We identify two separated. unresolved Lue cutters at a projected distance di;~0.77" oor 5.3 kpe. | We identify two separated, unresolved line emitters at a projected distance $d_{proj} \sim 0.7$ or 5.3 kpc. |
Each of the knots is spatially unresolved. (see vieht panel of Fig. | Each of the knots is spatially unresolved (see right panel of Fig. |
P. which shows the line distribution compared to a point source). aud we place upper lits ou their size from the size of the seciug disk. fucing FWITM of «2.72.9 kpe in right ascension and declination. respectively, | \ref{fig:o3im}
which shows the line distribution compared to a point source), and we place upper limits on their size from the size of the seeing disk, finding FWHM of $< 2.7 \times 2.3$ kpc in right ascension and declination, respectively. |
We do not detect amv line cinission from the diffuse. faint continuum source to the north west. | We do not detect any line emission from the diffuse, faint continuum source to the north west. |
The integrated spectrum of Q0317-383 C5 is shown in the inset of Fig. 2.. | The integrated spectrum of Q0347-383 C5 is shown in the inset of Fig. \ref{fig:spectrum}. |
The spectrum was iutegrated by sunuuiue over all spatial pixels in which the [OITIJA5007 cluission exceeds 30. | The spectrum was integrated by summing over all spatial pixels in which the $\lambda$ 5007 emission exceeds $\sigma$. |
The redshift is 170.0007. | The redshift is $\pm$ 0.0007. |
The |OITIIJAA 959.5007 doublet is very prominent. aud detected at 5o and 126. respectively, | The $\lambda\lambda$ 4959,5007 doublet is very prominent, and detected at $\sigma$ and $\sigma$ , respectively. |
Our results on the redshift. PWOAL and line fíux of [OTT] agree with those of ?.. but uulike ?.. we also detect IL) with a flux of Ej;=S8cxLN10D ere st 72 | Our results on the redshift, FWHM, and line flux of [OIII] agree with those of \citet{pettini01}, but unlike \citeauthor{pettini01}, we also detect $\beta$ with a flux of $_{H\beta} = 8.8\pm 1.8 \times 10^{-18}$ erg $^{-1}$ $^{-2}$. |
This flux is consistent with the 30 upper luit given by 7.. | This flux is consistent with the $\sigma$ upper limit given by \citet{pettini01}. |
We find a [OIII|/II./ ux ratio of [OM] το=7241.5. | We find a $/$ $\beta$ flux ratio of $_{5007}$ $H\beta= 7.2\pm $ 1.5. |
Line widths are FWIDMIso97=180+9 kan 3 for [OTLA5007, and ΕΠΙ 694311 kins | for IL}. respectively. | Line widths are $_{5007}$ $\pm$ 9 km $^{-1}$ for $\lambda$ 5007, and $_{H\beta}$ = $\pm$ 11 km $^{-1}$ for $\beta$, respectively. |
The lower ID} line width nay be due to its uulucky wavelength with respect to the telluric absorption. | The lower $\beta$ line width may be due to its unlucky wavelength with respect to the telluric absorption. |
We use the integrated spectrum to eive a rough Bos-like inetalicity estimate for QO3L7-383 Ch. | We use the integrated spectrum to give a rough $_{23}$ -like metalicity estimate for Q0347-383 C5. |
We did not measure the [OTJAA3726.3729 doublet. therefore. we use the correlation of [OII|A3727/|OTIT|A5O0 7. with [OIII|A5Q07 /TL? given by ?/— for low-1etalicity galaxies to estimate the most likely [OTJA3727 flux. | We did not measure the $\lambda\lambda3726,3729$ doublet, therefore, we use the correlation of $\lambda$ $\lambda$ 5007 with $\lambda$ $\beta$ given by \citet{kobulnicky99} for low-metalicity galaxies to estimate the most likely $\lambda$ 3727 flux. |
With the mneasured uncertainties and thelogt|OZI]/IP.j)=Ot sugeested by the ? correlation. Ro;=1.05. | With the measured uncertainties and the$\log([OII]/H\beta) = 0.4$ suggested by the \citeauthor{kobulnicky99}
correlation, $_{23}=1.05$. |
If instead we only use the measured [OTH] aud Ts values. aud neglect any contribution from [OT]. we find Rosor=0.95. | If instead we only use the measured [OIII] and $\beta$ values, and neglect any contribution from [OII], we find $_{23,OIII} = 0.95$. |
This corresponds to a highly conservative. but probably very oose lower bouud. | This corresponds to a highly conservative, but probably very loose lower bound. |
Including thelo uucertamties of our flux measurements. this corresponds to Bos20.560, or a netalicity between 8.6 aud 7.9. | Including the$1\sigma$ uncertainties of our flux measurements, this corresponds to $_{23}> 0.86$, or a metalicity between 8.6 and 7.9. |
This estimate lua appear relatively uncertain. ont we cluphasize that this is the case for any uctalicity estimate of high-redshift galaxies frou euission ines. | This estimate may appear relatively uncertain, but we emphasize that this is the case for any metalicity estimate of high-redshift galaxies from emission lines. |
Even the sample of ὃν which had measured OII[A3127. |OIII|AA 959.5007. and IL) fluxes. could rot be corrected for extinction. which will introduce considerable uncertainties. | Even the sample of \citet{pettini01}, which had measured $\lambda3727$, $\lambda\lambda$ 4959,5007, and $\beta$ fluxes, could not be corrected for extinction, which will introduce considerable uncertainties. |
With these caveats dn ual. Q0317-383 C5 has an oxygen abundance similar to those of the subsample of ? with Ro» measured. | With these caveats in mind, Q0347-383 C5 has an oxygen abundance similar to those of the subsample of \citet{pettini01} with $_{23}$ measured. |
Comparing with the solar oxygen abundance estimate of ?.. [O/I7|= £0.05. we find that QO317-383 Ch has amildly subsolar metalicity ranging LU/ITI|-- 0.1 0.7dex. | Comparing with the solar oxygen abundance estimate of \citet{allende01}, , $[O/H] = 8.69\pm 0.05$ , we find that Q0347-383 C5 has amildly subsolar metalicity ranging $[M/H]$$\sim$$-$ 0.1 – $-$ 0.7dex. |
uomentun dissipation of very low-niass stars. is described iu Delfosse ct al (19984). | momentum dissipation of very low-mass stars, is described in Delfosse et al (1998a). |
Even though this was not the main focus of the ooesran. we also realised from the start that for most of these stars we obtain radial velocity precisious which are sufficient to detect eiaut planets. if amy exists around hei. | Even though this was not the main focus of the program, we also realised from the start that for most of these stars we obtain radial velocity precisions which are sufficient to detect giant planets, if any exists around them. |
We present in this letter the first such detection. around Cl 876 1576290. LIIS 530. Ross του, TWP 113020). à V—10.2 ATL chwart (Reid et al.. | We present in this letter the first such detection, around Gl 876 $-15\degr$ 6290, LHS 530, Ross 780, HIP 113020), a V=10.2 M4 dwarf (Reid et al., |
1995) at d= LT024z0.016 pc (ESA. 1997). | 1995) at d = $\pm$ 0.046 pc (ESA, 1997). |
Delfosse et al. ( | Delfosse et al. ( |
1998a) present iu detail the observe siuuple. while Delfosse et al. ( | 1998a) present in detail the observed sample, while Delfosse et al. ( |
1998) discuss the observing anc analysis techuique at length. | 1998b) discuss the observing and analysis technique at length. |
We therefore ouly briefly sunuuarize this information in section 2. | We therefore only briefly summarize this information in section 2. |
We then procece to discuss in section 3H the radial velocity detection of the planetary companion of GL 576. | We then proceed to discuss in section 3 the radial velocity detection of the planetary companion of GL 876. |
Iu section | we consider the nuplicatious of this detection and sugees some possible follow-up observatious. | In section 4 we consider the implications of this detection and suggest some possible follow-up observations. |
The sample contains the 127 AL dwarts listed in he third edition of the nearby star catalog (CNS3 oelinunuryv version. Clese Jahres. 1991) with a distance closer than 9 pe. a DI950.0. declination above -16 degrees. brighter than V=15. aud without a close uuch brighter primary. | The sample contains the 127 M dwarfs listed in the third edition of the nearby star catalog (CNS3 preliminary version, Gliese Jahreiss, 1991) with a distance closer than 9 pc, a B1950.0 declination above -16 degrees, brighter than V=15, and without a close much brighter primary. |
Observations Lave been carried out since September 1995 with the ELODIE fiber-ed spectrograph (Darauue ct al. | Observations have been carried out since September 1995 with the ELODIE fiber-fed spectrograph (Baranne et al., |
1996) ou the 1.9314 clescope at Observatoire de Taute Provence (OTP). | 1996) on the 1.93m telescope at Observatoire de Haute Provence (OHP). |
The R=12000 spectra are wavelength calibrated through siauultaneous observations of a thorum lamp. | The R=42000 spectra are wavelength calibrated through simultaneous observations of a thorium lamp. |
Since Juue 1998 some southern stars have also been observed with the nearly ideutical CORALIE spectrograph ou the recently commissioned swiss 1.21i telescope at la Silla (Chile). | Since June 1998 some southern stars have also been observed with the nearly identical CORALIE spectrograph on the recently commissioned swiss 1.2m telescope at la Silla (Chile). |
cm CORALIE mostly differs from the older ELODIE iustruineut by its spectral resolution of R=hO000. better suupliug of the spectrograph PSF by the CCD camera pixels. aud. significantly better temperature coutrol. | CORALIE mostly differs from the older ELODIE instrument by its spectral resolution of R=50000, better sampling of the spectrograph PSF by the CCD camera pixels, and significantly better temperature control. |
The fist indications are that these modificatious together result in a substantialle improved iutriusic stability (Queloz et ab. | The first indications are that these modifications together result in a substantially improved intrinsic stability (Queloz et al., |
in preparation). | in preparation). |
The extracted AL dwarf spectra are analysed through cross-correlation with a binary (0/1) teiiplate coustructed yon an observed ELODIE spectru of Barnard’s star. C1699 (Delfosse et al.. | The extracted M dwarf spectra are analysed through cross-correlation with a binary (0/1) template constructed from an observed ELODIE spectrum of Barnard's star, Gl699 (Delfosse et al., |
1998b). | 1998b). |
For slowly rotating stars he resulting velocities have internal standard errors (photon noise plus low level uncalibrated iustrumieutal instabilities) which typically range from 10-15 + for xisbt M dwarfs 510) to ~+75 tat the magnitude indt of the sample. | For slowly rotating stars the resulting velocities have internal standard errors (photon noise plus low level uncalibrated instrumental instabilities) which typically range from 10-15 $^{-1}$ for bright M dwarfs 10) to $\sim$ 75 $^{-1}$ at the magnitude limit of the sample. |
For GL 876 (V=10.2) typical standard errors are LO to 20 1. depending on airmass and seeing conditions. | For Gl 876 (V=10.2) typical standard errors are 10 to 20 $^{-1}$, depending on airmass and seeing conditions. |
\laguetic activity is common in AL type cwarfs. and imav further degrade the measurement accuracy (Saar ct aL. | Magnetic activity is common in M type dwarfs, and may further degrade the measurement accuracy (Saar et al., |
1998). | 1998). |
This poteutial error source ds still unconpletelv characterised for AL chwarts. but for slowly rotating stars (Vsini«Menos cC) we cau already. bound it to ay, | for our cross-correlalon analysis with the ALL binary template. | This potential error source is still uncompletely characterised for M dwarfs, but for slowly rotating stars $V\,sin{i}<3km.s^{-1}$ ) we can already bound it to $\sigma_{\mathrm Vr}$ $^{-1}$ for our cross-correlation analysis with the M4 binary template. |
Within the brighter two thirds of the sample. a conservative asstion a the preseut time is that we will detect all varialles with seiui-auuplitudes lavgcr than +. | Within the brighter two thirds of the sample, a conservative assumption at the present time is that we will detect all variables with semi-amplitudes larger than $^{-1}$. |
Asstumine or illustration a 5 vears period. this corresponds to a l Jupiter uass (Mj) planet orbiting a 0.25. MIV. primary (at Ls AU). or to a 2 M planet orbiting a 0.6 AOV priuary (at 2.5 AU). | Assuming for illustration a 5 years period, this corresponds to a 1 Jupiter mass $_J$ ) planet orbiting a 0.25 M4V primary (at 1.8 AU), or to a 2 $_J$ planet orbiting a 0.6 M0V primary (at 2.5 AU). |
Since planet detection was not initially eiipliasized in the observing οστά, its sampling strategv is not optimal for detection of low amplitude variations on timescales shorter than a few vears. | Since planet detection was not initially emphasized in the observing program, its sampling strategy is not optimal for detection of low amplitude variations on timescales shorter than a few years. |
Cl 876 was observed once at cach observing seasous in 1905 axd 1996 and its velocity varlations became appareut frou he three observatious obtained in late 1997. | Gl 876 was observed once at each observing seasons in 1995 and 1996 and its velocity variations became apparent from the three observations obtained in late 1997. |
It was thor1 marked n our prograni ists as a variable. | It was then marked in our program lists as a variable. |
This low delination source however )'cane uuobservable from OUP before we could gather nore data and determine its οἱhit. | This low declination source however became unobservable from OHP before we could gather more data and determine its orbit. |
The commissioning of the swiss 1.21 telescope at la Silla aud its CORALIE spectrograph in June 1998 provied the first opportunity Oo obtain 23 additional measurenents of this southerly source. which allowed » finally determine its orbit. | The commissioning of the swiss 1.2m telescope at la Silla and its CORALIE spectrograph in June 1998 provided the first opportunity to obtain 3 additional measurements of this southerly source, which allowed to finally determine its orbit. |
These observations were obtained within two weeks of the first ight ofthis telescope. providing au encouraging indication on its potential for planet discovery, | These observations were obtained within two weeks of the first light of this telescope, providing an encouraging indication on its potential for planet discovery. |
An end of night neasurement from OIIP at a laree airmass provided a confirmation on June 22. just in time to confidently aunouuce the discovery at the LAU "Precise Stellar Radial Velocities” οςiferenice. (Victoria. Canada. June 21* to 26). | An end of night measurement from OHP at a large airmass provided a confirmation on June 22, just in time to confidently announce the discovery at the IAU “Precise Stellar Radial Velocities” conference (Victoria, Canada, June $^{st}$ to $^{th}$ ). |
At this conference we learucd from C. Marcy that his group indepeudeutly identified the orbit of CI 576. with orbital elements compatible with our own determination. | At this conference we learned from G. Marcy that his group independently identified the orbit of Gl 876, with orbital elements compatible with our own determination. |
Weather permitting. we have since then attempted to observe Cl 876 at most every three nights. and often every night. | Weather permitting, we have since then attempted to observe Gl 876 at most every three nights, and often every night. |
The 1998 data therefore donate the orbital solution. | The 1998 data therefore dominate the orbital solution. |
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