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We plot the MCD bolometric flux versus its temperature at the inner disk radius kTycp in Figure 4 (the upper panel)
We plot the MCD bolometric flux versus its temperature at the inner disk radius $kT_{\rm MCD}$ in Figure \ref{fig:spfits2} (the upper panel).
We can see that the evolution of the MCD luminosity is consistent with the L«T track (the solid line), which implies a constant inner disk radius with the change in luminosity.
We can see that the evolution of the MCD luminosity is consistent with the $L\propto T^4$ track (the solid line), which implies a constant inner disk radius with the change in luminosity.
We note that this is based on the only two observations available.
We note that this is based on the only two observations available.
The disk temperature is relatively low, only KTucp=65.8 and 93.1 eV for XMMI and XMM2, respectively.
The disk temperature is relatively low, only $kT_{\rm MCD}=65.8$ and $93.1$ eV for XMM1 and XMM2, respectively.
The PL component is weak, and its parameter values have relatively large uncertainties.
The PL component is weak, and its parameter values have relatively large uncertainties.
Its index is consistent between XMM1 and XMM2 and is relatively steep, with ΤΡΙ, about 3.5.
Its index is consistent between XMM1 and XMM2 and is relatively steep, with $\Gamma_{\rm PL}$ about 3.5.
Forcing XMM1 and XMM2 to have the same value of ΓΡΙ, in the fit, we see a change of the PL normalization Npy, by a factor of 2.6 (4.5 c).
Forcing XMM1 and XMM2 to have the same value of $\Gamma_{\rm PL}$ in the fit, we see a change of the PL normalization $N_{\rm PL}$ by a factor of 2.6 (4.5 $\sigma$ ).
The model SIMPL(MCD) gives results very similar to the model MCD+PL (Table 3 and Figure 4)), in terms of the MCD temperature, the thermal fraction, etc.
The model SIMPL(MCD) gives results very similar to the model MCD+PL (Table \ref{tbl:mcd+pl} and Figure \ref{fig:spfits2}) ), in terms of the MCD temperature, the thermal fraction, etc.
It infers that only about fgc=3% and 196 of the thermal disk emission is Comptonized to the hard emission in XMM1 and XMM2, respectively.
It infers that only about $f_{\rm SC}$$=$$3\%$ and $1\%$ of the thermal disk emission is Comptonized to the hard emission in XMM1 and XMM2, respectively.
This model, with a natural cutoff at low energies for the hard component, infers luminosities similar to those of the model MCD+PL obtained by integrating the PL flux down to 0.2 keV (Table 3)).
This model, with a natural cutoff at low energies for the hard component, infers luminosities similar to those of the model MCD+PL obtained by integrating the PL flux down to 0.2 keV (Table \ref{tbl:mcd+pl}) ).
The spectra can also be fitted almost equally well using the models BB+PL and SIMPL(BB) (Table 3)).
The spectra can also be fitted almost equally well using the models BB+PL and SIMPL(BB) (Table \ref{tbl:mcd+pl}) ).
The BB component dominates in both ΧΜΜΙ and XMM2, contributing 280% of the 0.2-10 keV flux, similar to the MCD component in the models MCD+PL and SIMPL(MCD).
The BB component dominates in both XMM1 and XMM2, contributing $\gtrsim$ of the 0.2–10 keV flux, similar to the MCD component in the models MCD+PL and SIMPL(MCD).
Its effective temperature kTpp is also low, about 58 and 78 eV for XMM1 and XMM2, respectively.
Its effective temperature $kT_{\rm BB}$ is also low, about 58 and 78 eV for XMM1 and XMM2, respectively.
The MCD and BB models have very similar spectral shapes at high energies (Makishimaetal. 1986),, but their differences become large at low energies.
The MCD and BB models have very similar spectral shapes at high energies \citep{mamami1986}, but their differences become large at low energies.
We estimate their differences in the UV.
We estimate their differences in the UV.
We have measurements from two UV filters, i.e., UVWI and UVM2, from XMM1.
We have measurements from two UV filters, i.e., UVW1 and UVM2, from XMM1.
The flux densities of the MCD component in the model MCD+PL from XMM1 are (2.15+0.69) and (3.6831.18) x10" erg s! em? À-! at the effective wavelengths of UVWI (2910 À)) and UVM2 (2310 À)), respectively.
The flux densities of the MCD component in the model MCD+PL from XMM1 are $2.15$$\pm$$0.69$ ) and $3.68$$\pm$$1.18$ $\times$ $^{-17}$ erg $^{-1}$ $^{-2}$ $^{-1}$ at the effective wavelengths of UVW1 (2910 ) and UVM2 (2310 ), respectively.
The corresponding values for the BB component in the model BB+PL from XMM1 are (0.40+0.16) and (1.00--0.40)x 1071? erg s~4 cm~? Á-1, respectively.
The corresponding values for the BB component in the model BB+PL from XMM1 are $0.40$$\pm$$0.16$ ) and $1.00$$\pm$$0.40$ $\times$ $^{-19}$ erg $^{-1}$ $^{-2}$ $^{-1}$, respectively.
The corresponding flux densities measured with UVW1 and UVW2 are (3.0040.47) and 10719 erg s7! cm? À-1, respectively, after the Galactic (5.08+1.45)xdust extinction correction using a reddening value of E(g.vy,=0.098 (Schlegeletal.1998) and assuming a spectral shape of a MCD model at low frequencies (ie., a power law with a photon index of 2/3).
The corresponding flux densities measured with UVW1 and UVW2 are $3.00$$\pm$$0.47$ ) and $5.08$$\pm$$1.45$ $\times$ $^{-16}$ erg $^{-1}$ $^{-2}$ $^{-1}$, respectively, after the Galactic dust extinction correction using a reddening value of $E_{\rm (B-V)}=0.098$ \citep{scfida1998} and assuming a spectral shape of a MCD model at low frequencies (i.e., a power law with a photon index of $2/3$ ).
We see that the UV flux from the OM detection is much higher than the BB flux in the UV.
We see that the UV flux from the OM detection is much higher than the BB flux in the UV.
It is closer to the MCD flux in the UV, but still there is aboutan order of magnitude difference, which will be discussed in Section 4..
It is closer to the MCD flux in the UV, but still there is aboutan order of magnitude difference, which will be discussed in Section \ref{sec:discussion}. .
The flux of the PL component in the UV is
The flux of the PL component in the UV is
We] have obtained. NIR⇁ broad-band images⊀ in⊀ the 4 (1.255)" ff (1.65 yam) and. dy (2.2 jim) bands of the 12 barred galaxies that make up our sample. during three observing runs: 1994 September26. 1995 November56. and 1996 February 3:5.
We have obtained NIR broad-band images in the $J$ $\mu$ m), $H$ (1.65 $\mu$ m) and $K$ (2.2 $\mu$ m) bands of the 12 barred galaxies that make up our sample, during three observing runs: 1994 September 26, 1995 November 5–6, and 1996 February 3–5.
The observations were made with the 36m CanadaFranceHawaii Telescope (CIT).
The observations were made with the 3.6 m Canada–France–Hawaii Telescope ).
We used the Montreal NER camera (ALONICA: Nadeau et al.
We used the Montreal NIR camera (MONICA; Nadeau et al.
L994). equipped with a 256 x 256 pixel HoCd'Te array detector with a projected pixel size of07248.
1994), equipped with a 256 x 256 pixel HgCdTe array detector with a projected pixel size of.
Since the sky background in the NIB. is bright. and changes rapidly. it is essential to obtain sky frames frequenthy.
Since the sky background in the NIR is bright and changes rapidly, it is essential to obtain sky frames frequently.
We observed. the sky usually just. before. and alter cach galaxy observation.
We observed the sky usually just before and after each galaxy observation.
Frames typically consisted of four co-adcded: individual exposures.
Frames typically consisted of four co-added individual exposures.
Separate frames with the nucleus of the galaxy at slightly offset positions were co-added in the reduction process to produce the final mosaic images.
Separate frames with the nucleus of the galaxy at slightly offset positions were co-added in the reduction process to produce the final mosaic images.
The weather was clear and generally photometric thoughout these nights. and the seeing values as measured from our final images were οτ 170.
The weather was clear and generally photometric thoughout these nights, and the seeing values as measured from our final images were 0.7 – 1.0.
The main steps in the reduction of the NUR data include subtracting the sky background. interpolating across known bad. pixels. and registering and combining the images.
The main steps in the reduction of the NIR data include subtracting the sky background, interpolating across known bad pixels, and registering and combining the images.
The data reduction was done partly. with. private. programmes and partly with standard routines.
The data reduction was done partly with private programmes and partly with standard routines.
We first combined four sky exposures taken at clillerent positions on the skv immediately before and after a series of galaxy exposures by median-averaging then. while iteratively rejecting those values in the averaged sky exposure that ceviate by more than 360 [rom the average value in that frame.
We first combined four sky exposures taken at different positions on the sky immediately before and after a series of galaxy exposures by median-averaging them, while iteratively rejecting those values in the averaged sky exposure that deviate by more than $\sigma$ from the average value in that frame.
This procedure allows the automated rejection of any star images that might have been present in the sky exposures. which were observed by olfsetting the telescope blindly to a position a few arcmin from the galaxy centre.
This procedure allows the automated rejection of any star images that might have been present in the sky exposures, which were observed by offsetting the telescope blindly to a position a few arcmin from the galaxy centre.
We then subtracted the averaged sky image [rom the relevant galaxy images.
We then subtracted the averaged sky image from the relevant galaxy images.
A close inspection of the sky-subtracted galaxy iniages of the 1996 run showed the presence of periodic. horizontal lines which could be traced back to electronic crosstalk in the svstem.
A close inspection of the sky-subtracted galaxy images of the 1996 run showed the presence of periodic, horizontal lines which could be traced back to electronic crosstalk in the system.
This raises the noise level in these images.
This raises the noise level in these images.
In order to correct the problem. we used a Fourier transform technique to locate the of frequenciesthe maximum intensity in the images. and filtered. them out.
In order to correct the problem, we used a Fourier transform technique to locate the frequencies of the maximum intensity in the images, and filtered them out.
This resulted. in a considerable improvement in most images. although in a few images interference stripes are still notable at low levels.
This resulted in a considerable improvement in most images, although in a few images interference stripes are still notable at low levels.
After Uat-liclding with dome Hats. anc masking the unreliable. pixels. in. the array. we combined. the skv-subtracted and. where appropriate. ce-stripecl images.
After flat-fielding with dome flats, and masking the unreliable pixels in the array, we combined the sky-subtracted and, where appropriate, de-striped images.
Ideally. one would like to locate one or more common field
Ideally, one would like to locate one or more common field
task which has been optimized to process ACS imaging data.
task which has been optimized to process ACS imaging data.
The task removes cosmic rav events and geometric distortions. and it drizzles the citherecl frames together into a final. high-resolution. photometric image.
The task removes cosmic ray events and geometric distortions, and it drizzles the dithered frames together into a final, high-resolution, photometric image.
We aligned the final {11 images to the LGS coordinate svstem wilhecimap using stars common to both data sets.
We aligned the final $HST$ images to the LGS coordinate system with using stars common to both data sets.
The stars positions in both the LGS and ACS images were determined withnmeentroid.
The stars positions in both the LGS and ACS images were determined with.
The resulting alignment had rms errors of 0.02" (less than 1 ACS pixel).
The resulting alignment had rms errors of $\leq$ $''$ (less than 1 ACS pixel).
The ACS images. independently aligned with the LGS coordinate system. are shown in Figure 1..
The ACS images, independently aligned with the LGS coordinate system, are shown in Figure \ref{ims}.
Resolved stellar photometry was then performed on the relevant sections of the final images with DAOPIIOT II and ALLSTAR (Stetsonetal.1990).
Resolved stellar photometry was then performed on the relevant sections of the final images with DAOPHOT II and ALLSTAR \citep{stetson}.
. The count rates measured from our images were converted to VEGA magnitudes using the conversion techniques provided in the ACS DataHandbook?.
The count rates measured from our images were converted to VEGA magnitudes using the conversion techniques provided in the ACS Data.
. source 13-127 was clearly detected 4 times by our monitoring program.
Source r3-127 was clearly detected 4 times by our monitoring program.
These observations provided a lighteurve. precise positional constraints. ancl X-ray spectral measurements.
These observations provided a lightcurve, precise positional constraints, and X-ray spectral measurements.
The X-awv lighteurve of 13-127 is shown in Figure 2.. and the measured. [Iuxes. and hardness ratios are provided in Table 2..
The X-ray lightcurve of r3-127 is shown in Figure \ref{lc}, and the measured fluxes and hardness ratios are provided in Table \ref{flux}.
Observations before the [ist detection allowed reliable upper limits to the X-ray {lis from that position to be measured.
Observations before the first detection allowed reliable upper limits to the X-ray flux from that position to be measured.
Our final observation vielded a marginal detection with a signal-to-noise of 2.
Our final observation yielded a marginal detection with a signal-to-noise of 2.
We treated (his measurement as a detection keeping in mind (hat if it was spurious the high end of the error range gives the 3c upper-limit of the flux during the final observation (2004-December-05).
We treated this measurement as a detection keeping in mind that if it was spurious the high end of the error range gives the $\sigma$ upper-limit of the flux during the final observation (2004-December-05).
The lighteurve is complex. including a double-peak. as has been observed for several Galactic XRNe (see lighteurves in MeClintock&Remillard 2004)).
The lightcurve is complex, including a double-peak, as has been observed for several Galactic XRNe (see lightcurves in \citealp{mcclintock2004}) ).
There was a [actor of 3 drop in flux followed
There was a factor of 3 drop in flux followed
A different evolutionary. path is followed in the case of high mass haloes. however. since centrifugal ellects become important in halting radial collapse.
A different evolutionary path is followed in the case of high mass haloes, however, since centrifugal effects become important in halting radial collapse.
We may. estimate the halo mass at which this occurs by equating the maximum density attained in quasi-spherical collapse (7ny/A7. where Ais the usual spin parameter) to no derived above. and find a critical halo mass of llere. the spin parameter is normatisecl to a ficlucial value of 0.1. close to what is found in cosmological simulations (e.g. Padmanabhan 1993).
We may estimate the halo mass at which this occurs by equating the maximum density attained in quasi-spherical collapse $\sim n_1/\lambda^3$, where $\lambda$ is the usual spin parameter) to $n_2$ derived above, and find a critical halo mass of Here, the spin parameter is normalised to a fiducial value of 0.1, close to what is found in cosmological simulations (e.g. Padmanabhan 1993).
In such massive haloes. once radial contraction is slowed by centrifugal support. shocks develop as infalling gas joins the incipient disc.
In such massive haloes, once radial contraction is slowed by centrifugal support, shocks develop as infalling gas joins the incipient disc.
The introduction. of entropy in shocks causes the evolution in the ».7] plane to stecpen with respect to the slope of adiabatie evolution. ancl we estimate. the approximate postshock parameters to be: ncn,fA? and doce T
The introduction of entropy in shocks causes the evolution in the $[n,T]$ plane to steepen with respect to the slope of adiabatic evolution, and we estimate the approximate postshock parameters to be: $n \simeq n_1/\lambda^3$ and $T \simeq T_{\rmn vir}$.
hus. to summarise. prior to the onset of cooling. the gas arrives at a state with temperature 7;=i; and clensity n,—minns.nifd? where the relevant. density depends on whether the halo mass is less than or greater than AZ.
Thus, to summarise, prior to the onset of cooling, the gas arrives at a state with temperature $T_v=T_{\rmn vir}$ and density $n_v = {\rm min} \biggl(n_2,n_1/\lambda^3\biggr)$, where the relevant density depends on whether the halo mass is less than or greater than $M_c$.
For M<M, rotational support is negligible in the virial state. whereas in the opposite limit. rotation. thermal auc gravitational effects are all comparable prior to the onset of cooling.
For $M<M_c$, rotational support is negligible in the virial state, whereas in the opposite limit, rotation, thermal and gravitational effects are all comparable prior to the onset of cooling.
The subsequent evolution depends on the elficiency. of cooling at conditions corresponding to n...
The subsequent evolution depends on the efficiency of cooling at conditions corresponding to $n_v,T_{v}$.
In general. the gas in a virialised dark matter (DM) halo will continue to collapse and fragment. if the eriterion. ρω<<fap ds satisfied (Rees Ostriker 19)
In general, the gas in a virialised dark matter (DM) halo will continue to collapse and fragment if the criterion $t_{\rmn cool} < t_{\rmn ff}$ is satisfied (Rees Ostriker 1977).
Systems with halo masses in excess of ~LOPAL. fulfil the Reos-Ostriker criterion either if ο cooling is effective (implying the absence of a sulliciently. strong photoclissociating UV field) or else. in the absence of photoionisation heating. if the metallicity exceeds c10°°Z. (Omukai 2000: Beomm et al.
Systems with halo masses in excess of $\sim 10^6 M_\odot$ fulfil the Rees-Ostriker criterion either if $_2$ cooling is effective (implying the absence of a sufficiently strong photodissociating UV field) or else, in the absence of photoionisation heating, if the metallicity exceeds $\simeq 10^{-3.5} Z_\odot$ (Omukai 2000; Bromm et al.
2001).
2001).
More massive systems. with Al2107A.(1|S10t. are able to collapse via atomic hydrogen lines (c.g. Macau. Lerrara. Rees 2001: Oh LHaiman 2002).
More massive systems, with $M\ga 10^{8}M_{\odot}[(1+z)/10]^{-1.5}$, are able to collapse via atomic hydrogen lines (e.g. Madau, Ferrara, Rees 2001; Oh Haiman 2002).
In each. case. we assume that the gas is compressed roughlyisobericeally as it cools.
In each case, we assume that the gas is compressed roughly as it cools.
Below. we ciscuss how the results of numerical simulations lend support to this idealized model.
Below, we discuss how the results of numerical simulations lend support to this idealized model.
During the isobaric phase the pressure is A similar dependence of gas pressure on mass and redshift compared. to our high-mass case has been found bv Norman Spaans (1997).
During the isobaric phase the pressure is A similar dependence of gas pressure on mass and redshift compared to our high-mass case has been found by Norman Spaans (1997).
“Phese authors have studied the properties of. protogalactic dises. and d ds therefore not surprising that our high-mass case. where centrifugal support is important. leads to a similar overall scaling.
These authors have studied the properties of protogalactic discs, and it is therefore not surprising that our high-mass case, where centrifugal support is important, leads to a similar overall scaling.
We assume that the immecdiate progenitor of a star is a centrally concentrated. scll-eravitating core (Motte. André. Neri 1998) with a mass close to the Bonnor-Ebert value (e.g. Palla 2002) To evaluate the minimum possible fragment. mass. we need to determine Zi, and μι
We assume that the immediate progenitor of a star is a centrally concentrated, self-gravitating core (Motte, André,, Neri 1998) with a mass close to the Bonnor-Ebert value (e.g. Palla 2002) To evaluate the minimum possible fragment mass, we need to determine $T_{\rm min}$ and $n_{\rm max}$.
Lhe minimum temperature is given hy where Note that although throughout the paper we refer to the case that temperatures of ~10 Ix are obtained as οΟἱ cooling. dust may also provide a means of achieving similarly low temperatures even in the absence of molecular coolants (sec. e.g. Whitworth. Bollin Francis 1998).
The minimum temperature is given by where Note that although throughout the paper we refer to the case that temperatures of $\sim 10$ K are obtained as `CO' cooling, dust may also provide a means of achieving similarly low temperatures even in the absence of molecular coolants (see, e.g. Whitworth, Boffin Francis 1998).
A lower Iloor on the gas temperature is set. by the cosmic microwave background (CMD) with We find the maximum density by assuming that the clissipative infall proceeds along an isobar. resulting in lnsertinge these expressions for Zi, and mua into equ. (
A lower floor on the gas temperature is set by the cosmic microwave background (CMB) with We find the maximum density by assuming that the dissipative infall proceeds along an isobar, resulting in Inserting these expressions for $T_{\rm min}$ and $n_{\rm max}$ into equ. (
7). we [lind for the low-mass case and for the high-mass case Depending on the details of how the gas is accreted onto the nascent protostar in the centre of the collapsing core. the resulting stellar mass is expected to be somewhat smaller than the Bonnor-Ebert mass.
7), we find for the low-mass case and for the high-mass case Depending on the details of how the gas is accreted onto the nascent protostar in the centre of the collapsing core, the resulting stellar mass is expected to be somewhat smaller than the Bonnor-Ebert mass.
We take this uncertainty into account by expressing the final characteristic mass as and choose the elliciencv to be a20.5.
We take this uncertainty into account by expressing the final characteristic mass as and choose the efficiency to be $\alpha\simeq 0.5$.
This value is close to that inferred. for the formation of stars in the present-day Universe (Melxee. Tan 2002).
This value is close to that inferred for the formation of stars in the present-day Universe (McKee Tan 2002).
This. choice of ellicieney. reproduces the high redshift (Pop LLL) point. estimating the typical mass of such a star to be a [ων hundred AZ.. collapsing in a halo of total mass ~10AL. at ~20 (Dromm et al.
This choice of efficiency reproduces the high redshift (Pop III) point, estimating the typical mass of such a star to be a few hundred $M_{\odot} $, collapsing in a halo of total mass $\sim 10^{6}M_{\odot}$ at $z\sim 20$ (Bromm et al.
1999. 2002: Nakamura Umemura 2001: Abel ct al.
1999, 2002; Nakamura Umemura 2001; Abel et al.
2002).
2002).
Clearly. this model is highly idealised. and serves only to provide an order of magnitude estimate for typical densities ancl pressures to be expected in regions of gas collapsing in proto-galactic potentials.
Clearly, this model is highly idealised, and serves only to provide an order of magnitude estimate for typical densities and pressures to be expected in regions of gas collapsing in proto-galactic potentials.
Some support for this schematic evolution in the 6».2] phase diagram: is provided bv the results of numerical simulations of the formation of dwarf galaxies at high redshifts.
Some support for this schematic evolution in the $[n,T]$ phase diagram is provided by the results of numerical simulations of the formation of dwarf galaxies at high redshifts.
For example. Figure 1 depicts the
For example, Figure 1 depicts the
uuuber ⋯∐∖⊼≻↕∪↥⋅↸∖≺↧⋅↕↕⋡↴⋝↕∐⋜∐⋅↕↸∖↴∖↴⋜⋯∖↸⊳∪∐⊔⊔∪∐⋜∐⊔∪∐∶↴∙⊾⋯⋜↧↴∖↴↴∖↴↕↖↽↸∖ ↙⋅↙ : (7). 30DDor (?).
\cite{1999PhDT.........7B} \citep{1999A&AS..137...21B} Dor \citep{2001A&A...380..137B}.
. The most important results of that analysis can be outlined as follows: Although these results are consistent with a binary population. it is necessary to confir this asstuuption with direct evidence.
The most important results of that analysis can be outlined as follows: Although these results are consistent with a binary population, it is necessary to confirm this assumption with direct evidence.
The issue of {binarybinary frequency is stilstill subjectstbice to debate.Iebate. as nost of. the statistics⋅⋅ on binary⋅ stars are biased⋅ owards later type stars :(?)..
The issue of binary frequency is still subject to debate, as most of the statistics on binary stars are biased towards later type stars \citep{2006ApJ...640L..63L}.
There is. evidence. o»ufius towards a relatively high παπα of üuaries amoug early-type stars (7).. ere we present a new set of observations of 22070 obtained with the Cemini Multi Object Spectrograph (CALOS) at Gemini South.
There is evidence pointing towards a relatively high number of binaries among early-type stars \citep{1998AJ....115..821M}, Here we present a new set of observations of 2070 obtained with the Gemini Multi Object Spectrograph (GMOS) at Gemini South.
These comprise multi object optical spectroscopy of] 50 carly-type stars observed at leastim. six
These comprise multi object optical spectroscopy of 50 early-type stars observed at leastin six
These comprise multi object optical spectroscopy of] 50 carly-type stars observed at leastim. six.
These comprise multi object optical spectroscopy of 50 early-type stars observed at leastin six
due to instrumeut drift. the Eartlis baryceutrie motion aud orbiting planets are removed. RV RMS error is 17.13 ms.| but has not reached the predicted photon-limited RV uncertainty (L8 m.s.f. S/N-80 with a half wavelength coverage from 535 to 565 nm).
due to instrument drift, the Earth's barycentric motion and orbiting planets are removed, RV RMS error is 17.13 $\rm{m}\cdot\rm{s}^{-1}$ but has not reached the predicted photon-limited RV uncertainty (4.8 $\rm{m}\cdot\rm{s}^{-1}$, S/N=80 with a half wavelength coverage from 535 to 565 nm).
RV RMS error is expected to be further reduced alter the data pipeline is improved iu the future.
RV RMS error is expected to be further reduced after the data pipeline is improved in the future.
We also examine the reference star GD calibration method.
We also examine the reference star GD calibration method.
We use oue spectral block within Ineasurement range centering at 550 ΤΗ» (510-560 THz) aud set GD! to be an arbitrary value of -23.873 ps.
We use one spectral block within measurement range centering at 550 THz (540-560 THz) and set $^\prime$ to be an arbitrary value of -23.873 ps.
The measured RVs (barycentric velocity not corrected) are shown in Fig. 7..
The measured RVs (barycentric velocity not corrected) are shown in Fig. \ref{fig:HIP_14810_before}.
After applying correction according to Equation (6)). we find that the CD is -25.107T40.027 ps.
After applying correction according to Equation \ref{eq:GD_cal_ref}) ), we find that the GD is $\pm$ 0.027 ps.
lu comparison. GD measurement result of fiber number 51 (the fiber for HIP 11310) using the WLC method gives -25.09140.005 ps (refer to Table 1)).
In comparison, GD measurement result of fiber number 51 (the fiber for HIP 14810) using the WLC method gives $\pm$ 0.005 ps (refer to Table \ref{tab:GD_err_fre}) ).
We coulfiriu that the GDs measured by these two methods are consistent with each other at significance level.
We confirm that the GDs measured by these two methods are consistent with each other at significance level.
The PV scale is an important parameter in the DEDI method that translates a measurecl phase shift to an RV shift. aud is determined by the group delay (GD) of au interferometer.
The PV scale is an important parameter in the DFDI method that translates a measured phase shift to an RV shift, and is determined by the group delay (GD) of an interferometer.