ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
The ability to use BW anc astrometric ueasuremenuts simultancously makes it possible to eiiploy observational timelines below the orbital ones aud stil ο able to make positive exoplanet detections. thus iclpiug to extract the inaxinui amount of information 1X1 nieasurenieuts aud increasing the time efficiencv of observations.	The ability to use RV and astrometric measurements simultaneously makes it possible to employ observational timelines below the orbital ones and still be able to make positive exoplanet detections, thus helping to extract the maximum amount of information from measurements and increasing the time efficiency of observations.
Tn a forthcoming study. we plan to study the inclusion of additional transit-photometry neasuremoenuts to further tighten the parameter probability ceusities iu transiting scenarios.	In a forthcoming study, we plan to study the inclusion of additional transit-photometry measurements to further tighten the parameter probability densities in transiting scenarios.
Also. the approach used here should be extended tosvstenis with two or more plauctary COMPaluons.	Also, the approach used here should be extended tosystems with two or more planetary companions.
Ζι=Zo cannot he mace unless the interior docs not rotate rigidy but instead the observed J, value is influcuced by deep winds (?)..	$Z_1=Z_2$ cannot be made unless the interior does not rotate rigidy but instead the observed $J_4$ value is influenced by deep winds \citep{Militzer+08}.
Ou the other laud. as reffig: Jup,oncergshouws.Z4Zo is not a favorite option for the DET-MD based Jupiter model. since that would imply Zi«Z. or even zero.	On the other hand, as \\ref{fig:Jup_converg} shows, $Z_1\ll Z_2$ is not a favorite option for the DFT-MD based Jupiter model, since that would imply $Z_1\ll Z_{\odot}$ or even zero.
This would coutracict heavy. clement abundauce measurements ni Jupiter's atinosphere. which indicate Z4=ὃνZ..	This would contradict heavy element abundance measurements in Jupiter's atmosphere, which indicate $Z_1\geq 2\times Z_{\odot}$.
Therefore. Z422Za is nota free choice for the DET-ME based model.	Therefore, $Z_1\approx Z_2$ is a free choice for the DFT-MD based model.
Consequeuth. the very. differeut. Jupiter core Inasses obtained usiug these two ab-initio EOS are not primarily due to different assumptions about the distribution of heavy elements.	Consequently, the very different Jupiter core masses obtained using these two ab-initio EOS are not primarily due to different assumptions about the distribution of heavy elements.
In order to resolve this problem. a comparison of the II/Ile adiabats is highly desirable.	In order to resolve this problem, a comparison of the H/He adiabats is highly desirable.
Furthermore. iscutropic compression experiments in the 0.5—[ Mbar regime where J» and J, are most scusitive to metallicty would ereatlv help to discriminate between compoetiug Jupiter models.	Furthermore, isentropic compression experiments in the $0.5-4$ Mbar regime where $J_2$ and $J_4$ are most sensitive to metallicty would greatly help to discriminate between competing Jupiter models.
For the other EOS considered here. both the assumptions 4=Za or Z4xZo» give acceptable Jupiter models aud hence are a iatter of free choice.	For the other EOS considered here, both the assumptions $Z_1=Z_2$ or $Z_1\not=Z_2$ give acceptable Jupiter models and hence are a matter of free choice.
As we have seen in refsec:Jupiter... Jupiter models that use oue particular EOS can have various resulting Mu.Z4.Z2} triples.	As we have seen in \\ref{sec:Jupiter}, Jupiter models that use one particular EOS can have various resulting $\{M_{core}, Z_1, Z_2\}$ triples.
Without the constraint imuposed by Z4, however. the varicty of solutions would be iuunueuse. iucludiug Moore>UsMy and Mz>38AL).	Without the constraint imposed by $J_4$ however, the variety of solutions would be immense, including $M_{core}>18\ME$ and $M_Z>38\ME$.
For extrasolar plaucts. curreutly available gravity data are M, aud Fe, only,	For extrasolar planets, currently available gravity data are $M_p$ and $R_p$ only.
A potentially observational constraint equivalent to Jy has been sugeested to be the tidal Love umber &» C2)... which measures the ability of a planet to develop an elliptic deformation in response to a tidal perturber such as the close-by parent star.	A potentially observational constraint equivalent to $J_2$ has been suggested to be the tidal Love number $k_2$ \citep{RW09}, which measures the ability of a planet to develop an elliptic deformation in response to a tidal perturber such as the close-by parent star.
In this section. we present implications of knowing a precise Ay value of the Neptune-sized extrasolar planet CJ 136b.	In this section, we present implications of knowing a precise $k_2$ value of the Neptune-sized extrasolar planet GJ 436b.
About 30 light vears away from Earth. the Hot Neptune 136b (A,=23.17Mj. RB,=£22 R4) (7) orbits the Mestar 136.	About 30 light years away from Earth, the Hot Neptune $\:$ 436b $M_p=23.17\ME$, $R_p=4.22\RE$ ) \citep{Torres+08} orbits the M-star $\:$ 436.
Placed ia a imiass-racdius diagram. this planet is located close to theoretical AZ-R relations of wari water planets.	Placed in a mass-radius diagram, this planet is located close to theoretical $M$ $R$ relations of warm water planets.
Accordingly. interior structure models assmuius au irou-silicate core. a water laver. ind a ΠΠΠο euvelope allow for a composition of water and U/Ue (?).. but also for a water-less composition.	Accordingly, interior structure models assuming an iron-silicate core, a water layer, and a H/He envelope allow for a composition of water and H/He \citep{Figueira+09}, but also for a water-less composition.
I/We cuvelope models. whether dry or not. are found to have 0.02<hy0,2 (2)..	H/He envelope models, whether dry or not, are found to have $0.02<k_2<0.2$ \citep{N-GJ436b+10}.
We here hy=0.2 and search for the core mass aud water conteut that satisfy this additional coustraint.	We here $k_2=0.2$ and search for the core mass and water content that satisfy this additional constraint.
We allow water be mixed homogencously iuto the IT/IIe euvelope (two-laver model) or confined to a deep water laver (three-laver models) aud combinations iu between (Neptune-like threc-Iaver models).	We allow water be mixed homogeneously into the H/He envelope (two-layer model) or confined to a deep water layer (three-layer models) and combinations in between (Neptune-like three-layer models).
Figure 2. shows the result in comparison with LMC-REOS based Jupiter models forced to meet the observed ο value. but uot J,	Figure \ref{fig:J2k2} shows the result in comparison with LM-REOS based Jupiter models forced to meet the observed $J_2$ value, but not $J_4$.
Iu both cases. the solutions are located along a straight line (solutions for different A5 or J5 values would span parallel lines).	In both cases, the solutions are located along a straight line (solutions for different $k_2$ or $J_2$ values would span parallel lines).
For both planets. aud hence in ecueral. the upper Iit i core mass 1s obtained for homogencous envelope models (compare	For both planets, and hence in general, the upper limit in core mass is obtained for homogeneous envelope models (compare
ancl processed the data with their iniproved. version of the WD code.	and processed the data with their improved version of the WD code.
More recently. a reanalysis of the Andersenctal.(1988). RV's was obtained by Ixaranii&Alohehi(2007) witha further improvement of the spectroscopic parameters.	More recently, a reanalysis of the \citet{and88} RV's was obtained by \citet{kar07} with a further improvement of the spectroscopic parameters.
Llowever. their solution can be considered as disputable since they seem not to have included several uncertainties (of the orbital period. for example) into the total error budget and thus presumably underestimated the final errors.	However, their solution can be considered as disputable since they seem not to have included several uncertainties (of the orbital period, for example) into the total error budget and thus presumably underestimated the final errors.
Our solution. based. on our RV data ancl ASAS photometry is compared with the ones from Andersenctal.(1088) and Miloneetal.(1992). in Table 4.	Our solution, based on our RV data and ASAS photometry is compared with the ones from \citet{and88} and \citet{mil92} in Table \ref{com_aiphe}.
In. general our results are in agreement with the previous papers.	In general our results are in agreement with the previous papers.
One should note that having only S RV measurements [ου every component. we were able to improve the spectroscopic results of Andersenetal.(1988). who hac46 datapoints for every component. by a factor of 9 to 4.	One should note that having only 8 RV measurements for every component, we were able to improve the spectroscopic results of \citet{and88}, who had 46 datapoints for every component, by a factor of 2 to 4.
Unfortunately. having only one light-curve from ASAS we cannot compete with the results of Miloneetal.(1992) who used light curves from 12 bandpasses: most of which were more precise than ours.	Unfortunately, having only one light-curve from ASAS we cannot compete with the results of \citet{mil92} who used light curves from 12 bandpasses; most of which were more precise than ours.
Still. their phase coverage is not complete and the orbital and physical parameters might. be improved. with high-precision photometry.	Still, their phase coverage is not complete and the orbital and physical parameters might be improved with high-precision photometry.
Our solution converged to a somewhat different value of the orbital inclination than Ancersen’s and Alillone’s.	Our solution converged to a somewhat different value of the orbital inclination than Andersen's and Millone's.
This is of course the reason for the discrepancy in absolute masses between the solutions.	This is of course the reason for the discrepancy in absolute masses between the solutions.
Our. AZsin’?/ is however Fu more precise.	Our $M\sin^3i$ is however far more precise.
“Phus we may conclude that this systems parameters can be derived with an unprecedented precision (possibly ~0.01 in masses and ~0.1 in radi) if only one had. more iocine RY measurements and millimagnitude photometry.	Thus we may conclude that this systems' parameters can be derived with an unprecedented precision (possibly $\sim 0.01$ in masses and $\sim 0.1$ in radii) if only one had more iodine RV measurements and millimagnitude photometry.
UX Alen. as well as Al Phe. was discovered to be à variable in Bamberg patrol plates and reported by Strohmeier (1966).	UX Men, as well as AI Phe, was discovered to be a variable in Bamberg patrol plates and reported by \citet{str66}.
. The first orbital solution was obtained by Lambert(1974) and the first full physical analysis. based on Imboerts radial velocities and photometry. was done w Clausen&CGronbech.(1976).	The first orbital solution was obtained by \citet{imb74} and the first full physical analysis, based on Imberts' radial velocities and photometry, was done by \citet{cla76}.
. Later. this photometric dataset was reanalised together with the new COILAVEL racial velocities by Andersenetal.(1989). who obtained he most up-to-date solution for UX Alen.	Later, this photometric dataset was reanalised together with the new CORAVEL radial velocities by \citet{and89} who obtained the most up-to-date solution for UX Men.
Comparison of his solution with our results is shown in Table 5..	Comparison of this solution with our results is shown in Table \ref{com_uxmen}.
Again. our results are in general in agreement with Andersen's.	Again, our results are in general in agreement with Andersen's.
As for Al Phe. having 58 radial velocity measurements of UX Alen. we succeeded to reach a better precision in the spectroscopic parameters than Andersen et al.	As for AI Phe, having 8 radial velocity measurements of UX Men, we succeeded to reach a better precision in the spectroscopic parameters than Andersen et al.
with 29 datapoints for the primary and 31 for the secondary.	with 29 datapoints for the primary and 31 for the secondary.
In our model we fixed the eccentricity to 0.	In our model we fixed the eccentricity to 0.
We found no signillicant improvement in the best-fitting mocel (in terms of rms) when e and o were set as free pramoeters and the resulting e was undistinguishable from0 within the formal errors.	We found no signifficant improvement in the best-fitting model (in terms of $rms$ ) when $e$ and $\omega$ were set as free prameters and the resulting $e$ was undistinguishable from 0 within the formal errors.
The non-zero eccentricity may be however induced by a putative third body. found in NAC'O images by 'l'okovininctal.(2006) about 0.751arcsec from the binary.	The non-zero eccentricity may be however induced by a putative third body, found in NACO images by \citet{tok06} about 0.751arcsec from the binary.
V415. Aql was first reported as a variable by HLolfmeister (1936).	V415 Aql was first reported as a variable by \citet{hof36}.
. The first ephemeris 2428670.532 \| [7 2.4628 d was determined by Guthnick&Schneller(1936	The first ephemeris 2428670.532 + $E \cdot$ 2.4628 d was determined by \citet{gut39}.
) The double-lined character was not. previously reported thus there was no racial velocity curve obtained to date and the only known ight curve analysis was done by Brancewiez&Dworak(1980) who derived. an orbital period. of 2.46273 d. Our UN data. despite having only 3 measurements for. every component. clearly shows that the actual period is (wo times onger.	The double-lined character was not previously reported thus there was no radial velocity curve obtained to date and the only known light curve analysis was done by \citet{bra80} who derived an orbital period of 2.46273 d. Our RV data, despite having only 3 measurements for every component, clearly shows that the actual period is two times longer.
This is confirmed by the ASAS light curve exhibiting unequal eclipses (see Eig. 12).	This is confirmed by the ASAS light curve exhibiting unequal eclipses (see Fig. \ref{fig_master}) ).
Η divided by 2. our period is in a good agreement with the value from Brancewiez&Dworak (1980).	If divided by 2, our period is in a good agreement with the value from \citet{bra80}.
. Unfortunately. the uncertainties of their results are unavailable.	Unfortunately, the uncertainties of their results are unavailable.
In Fable 6 we compare their results with ours.	In Table \ref{com_v415aql} we compare their results with ours.
The improvement in the derived parameters is obvious.	The improvement in the derived parameters is obvious.
Also he distance estimate given by Brancewicz&Dworak(1980 of this should be now treated with caution.	Also the distance estimate given by \citet[][see Tab. \ref{tab_info} of this should be now treated with caution.
Out of the eighteen eclipsing binaries in our sample. 15 are	Out of the eighteen eclipsing binaries in our sample, 15 are
represents illumination of gas bv the X-ray rotating beams in à region between the accretion column ancl the discs transition region.	represents illumination of gas by the X-ray rotating beams in a region between the accretion column and the disc's transition region.
Each pulse may then represent reprocessed. gas slightly further out. from. the continuum. region which is supported. by the phase delays of 0.1 spin eveles.	Each pulse may then represent reprocessed gas slightly further out from the continuum region which is supported by the phase delay of 0.1 spin cycles.
For zero phase delay. the continuum and. emission regions would lic in the same axis (racial infall).	For zero phase delay, the continuum and emission regions would lie in the same axis (radial infall).
For a yhase delav of 0.25 eveles. the corotation velocity. of the emission region would be perpencicular to the racial infall of he continuum region.	For a phase delay of 0.25 cycles, the corotation velocity of the emission region would be perpendicular to the radial infall of the continuum region.
Therefore. the emission region ies in a region between corotation and radial infall. aud most likely. tracing the magnetic field lines.	Therefore, the emission region lies in a region between corotation and radial infall, and most likely, tracing the magnetic field lines.
Such a delay is consistent with the field lines being swept back by the »erturbed Dow of disc matter.	Such a delay is consistent with the field lines being swept back by the perturbed flow of disc matter.
The nearly svmmetrie pulses in the continuum indicate accreting poles of equal brightness at anti-diamoetric ocations.	The nearly symmetric pulses in the continuum indicate accreting poles of equal brightness at anti-diametric locations.
The pulses are very asvmimetric with the irst pulse dominating over the second one (lux is higher by a factor of 2).	The pulses are very asymmetric with the first pulse dominating over the second one (flux is higher by a factor of 2).
Phe width of the pulses is identical. 318417 xIM n. which. indicates. à common OPLgL.M and sets upper imits for the broadening mechanism in the emission-region. perhaps the scale of turbulence (change from circular motion to quasi-racial) or the sound velocity (2<LO” Ix).	The width of the pulses is identical, $\pm$ 17 km $^{-1}$, which indicates a common origin, and sets upper limits for the broadening mechanism in the emission-region, perhaps the scale of turbulence (change from circular motion to quasi-radial) or the sound velocity $T < 10^{6}$ K).
While one pulse is in maximum. blueshift. the other is in maximum recshift ("corkscrew) which is entirely consistent with an accretion curtain.	While one pulse is in maximum blueshift, the other is in maximum redshift (“corkscrew”) which is entirely consistent with an accretion curtain.
“Phe maximum of the first. pulse is viewed. directly. [rom the (lower) pole whilst the curtain Mocks any visibility of theLL region from the upper »ole. resulting in the minimum of the second pulse (in the average-subtracted line. profiles).	The maximum of the first pulse is viewed directly from the (lower) pole whilst the curtain blocks any visibility of the region from the upper pole, resulting in the minimum of the second pulse (in the average-subtracted line profiles).
This varving view of the wo inner parts of the curtain explains the quaclrapole-like xutern of the pulsed. emission in Fig.	This varying view of the two inner parts of the curtain explains the quadrapole-like pattern of the pulsed emission in Fig.
10.	10.
Partial overlap of the pulses between spin. phases 0.50.7 suggests that the azimuthal extent is large enough so that the two pulse ocations can be viewed simultaneously at times.	Partial overlap of the pulses between spin phases 0.5–0.7 suggests that the azimuthal extent is large enough so that the two pulse locations can be viewed simultaneously at times.
The phase range of the pulses also suggests that the azimuthal extent of the accretion curtain in the continuum is smaller than hat of the emission regions (Lig.	The phase range of the pulses also suggests that the azimuthal extent of the accretion curtain in the continuum is smaller than that of the emission regions (Fig.
7).	7).
We associate the emission-line pulsations with material which has lost its Keplerian Dow due to the magnetic fields coupling and is most likely close to radial free-fall along he magnetic field lines (see previous discussion about phase delay).	We associate the emission-line pulsations with material which has lost its Keplerian flow due to the magnetic field's coupling and is most likely close to radial free-fall along the magnetic field lines (see previous discussion about phase delay).
Then. theLL emission region will be located within he magnetosphere of the white dwarf and will rotate with he spin period of 545 seconds.	Then, the emission region will be located within the magnetosphere of the white dwarf and will rotate with the spin period of 545 seconds.
Εις gives\|l2xdi -—gd./sn;--uAkms where Vie.-=408 km 17. 2?=545 seconds is. the spin. »eriod.	This gives$ V = V_{obs}/~\sin~i = \frac{2~\pi~R}{P} $ km $^{-1}$, where $V_{obs} = 408$ km $^{-1}$, $P =545$ seconds is the spin period.
"There is no eclipse in the X-ray and. optical light curves.	There is no eclipse in the X-ray and optical light curves.
On the other hand. the two accreting poles of the white chvarl (double-peak pulse) can be viewed only in hieh-inclination systenis. (50, 70°) and this is most likely the case for RX JO558\|5353.	On the other hand, the two accreting poles of the white dwarf (double-peak pulse) can be viewed only in high-inclination systems $50^{o}-70^{o}$ ) and this is most likely the case for RX J0558+5353.
Phe pulsations become visible only for favourable orientations of the rotating scarchlight beans to the line-of-sight.	The pulsations become visible only for favourable orientations of the rotating searchlight beams to the line-of-sight.
Truly. this is supported by a correlation xtween inclination andpulse amplitude (Ποτος and Mason 1990).	Truly, this is supported by a correlation between inclination andpulse amplitude (Hellier and Mason 1990).
For a free-fall velocity of 408 kim sn60°=471 kms +. the distance is οlot km or RY4.9RewAly? from the white chwarl. assuming an inclination of 60" and using Bag70.84totAZ,Ob" lm (Llamacla Salpeter 1901).	For a free-fall velocity of 408 km $^{-1}$ / $\sin~60^{o} = 471$ km $^{-1}$ , the distance is $ R \sim 4.1 \times 10^{4}$ km or $ R \sim 4.9~R_{wd}~M^{+0.6}_{0.6} $ from the white dwarf, assuming an inclination of $^{o}$ and using $R_{wd} \sim 0.84 \times 10^{4}~M^{-0.6}_{0.6} $ km (Hamada Salpeter 1961).
Τμ may well represent the magnetospheric radius where the disc. ds. disrupted. [I—B4. hall the Alfven radius. which gives Hoa~O26Row(ANLapiiy)dsLot km where Ly; is the luminosity in units of 107 eres s.! and pao is the magnetic moment in units of 107 € em? (Ghosh and Lamb LOTS).	This may well represent the magnetospheric radius where the disc is disrupted, $R=R_{mag}$, half the Alfven radius, which gives $ R_{mag} \sim 0.26 ~R_{wd}~ ( M_{0.6}^{0.91}~~L_{33}^{-2/7}~ \mu_{30}^{4/7}) = 4.1 \times 10^{4}$ km where $L_{33}$ is the luminosity in units of $10^{33}$ ergs $^{-1}$ and $\mu_{30}$ is the magnetic moment in units of $10^{30}$ G $^{3}$ (Ghosh and Lamb 1978).
The impact of the stream overflow on the discs transition region is slightlv further out at ⋅⋅UST ≼≻⋅↖∖↻∫↘⋟⋯∣⇀∪⊓⊽⊓↾⊽−↿∖↙∫∶≱⋅⋅↱≻⊐⋜⋯∠⇂⊔⋯∙∖⇁↓↥⋜↧∖⇁∢⊾↓∪∖∖⋎∢⊾↓⋅∖⇁⋖⋅⇂⋯∙↓↿⊓⋅≱∖∪⇂ MN ∿∶∫≻⋅↱≻∪↓∡⊔	The impact of the stream overflow on the disc's transition region is slightly further out at $6.86~R_{wd}~M_{0.6}^{-0.27}$ $q=0.5$ ) and may have lower velocities of $\sim$ 350 km $^{-1}$.
↓⊳∖⊥⊳↾∐↥∢⋅≼∙⋖⋟↓⋅∢≱∣⋜↧⇂↕⋖⋟↓↕↓⋅⋯∐⊔⊳∖∫⊐⋟⋯↿∖↓⊲↸↓⋅	The corotation radius $R_{co}$ (Frank et al.
⋜⋯↥⊔⊾↿⋜↧↓⊳↓≤⋗≤⊔⊐⋡ Peo~10RewMUTORF isH larger than the magnetosphere (slow rotator) for accretion to occur.	1992), $R_{co} \sim 10~R_{wd}~M^{-0.27}_{0.6}$, is larger than the magnetosphere (slow rotator) for accretion to occur.
αρα et al. (	Haberl et al. (
1994) derive an observed Iuminosity of 0.18107(—L300y eres + from the N-rav Hlux between 0.1-2.4. keV. Fittingpe the X-ray spectrum with a blackbody of 57 eV. (soft component) and a bremsstrahlung spectrum of 10. keV. (hard component) gives a luminosity of 1.710°ο... eres + for an absorption column density of 6107" LE em.? (LHaberl et al.	1994) derive an observed luminosity of $0.18~\times~10^{32}~ (\frac{d}{300~{\rm pc}})^{2}$ ergs $^{-1}$ from the X-ray flux between 0.1-2.4 keV. Fitting the X-ray spectrum with a blackbody of 57 eV (soft component) and a bremsstrahlung spectrum of 10 keV (hard component) gives a luminosity of $1.7~\times~10^{33}~(\frac{d}{300~{\rm pc}})^{2}$ ergs $^{-1}$ for an absorption column density of $6~\times~~10^{20}$ H $^{-2}$ (Haberl et al.
1994).	1994).
‘Phis can be used to estimate the magnetic moment of the white dwarl of 2.407 G em? (or B0.5 MG). assuming a Ao= 1. Ly;= 1. d=300 pe and a negligible softX-ray luminosity compared to the hard. X-ray component.	This can be used to estimate the magnetic moment of the white dwarf of $2.4~\times~10^{32}$ G $^{3}$ (or $B \sim 0.5$ MG), assuming a $M_{0.6} = 1$ , $L_{33}=1$ , $d = 300$ pc and a negligible softX-ray luminosity compared to the hard X-ray component.
Such a magnetic moment is not inconsistent with the existence of an accretion disc. evidence for which we found. from the periodogram	Such a magnetic moment is not inconsistent with the existence of an accretion disc, evidence for which we found from the periodogram
is Ἱπαάσσα at 2=0.157 aud hence likely to be the 21-cim absorber.	is indeed at $z = 0.437$ and hence likely to be the 21-cm absorber.
The quasar field was imaged using the EFOSC? iustrument on the ESO 3.611 telescope at La Silla duriug the night of 29th Jaunary. 2001.	The quasar field was imaged using the EFOSC2 instrument on the ESO 3.6m telescope at La Silla during the night of 29th January, 2001.
The Bessel BR-buik 1inage was obtained on CCD #110 in the uubinned mode (0.157 aresec/pix. 2060 « 2060 pixels = 5.1 « 5.1 arcmin feld).	The Bessel R-band image was obtained on CCD 40 in the unbinned mode (0.157 arcsec/pix, 2060 $\times$ 2060 pixels $\equiv$ 5.4 $\times$ 5.4 arcmin field).
We obtained a total of teu dithex inages of 150s cach.	We obtained a total of ten dithered images of 450s each.
The calibration observations iucluded twilielr sky flats aud observations of the stadia fell RUI52 (Landolt 1992)).	The calibration observations included twilight sky flats and observations of the standard field RU152 \cite{landolt92}) ).
We used the seieuce frames ο construc and subtract the fringe pattern from the frames.	We used the science frames to construct and subtract the fringe pattern from the frames.
The atmospheric extinction and the CCD colour term were calculated using observations of standard fields taken at several airimasscs and in other bands.	The atmospheric extinction and the CCD colour term were calculated using observations of standard fields taken at several airmasses and in other bands.
We expect the photometry to be accurate to about and use a conservative error of 0.1 mae.	We expect the photometry to be accurate to about and use a conservative error of 0.1 mag.
The images were reduced. calibrated. registered aud co-added m a standard mauner using IRAF.	The images were reduced, calibrated, registered and co-added in a standard manner using IRAF.
The final image obtained had a zero-point of R=32.601 with au RAIS of 10.5 counts/pix.	The final image obtained had a zero-point of R=32.601 with an RMS of 10.5 counts/pix.

The ability to use BW anc astrometric ueasuremenuts simultancously makes it possible to eiiploy observational timelines below the orbital ones aud stil ο able to make positive exoplanet detections. thus iclpiug to extract the inaxinui amount of information 1X1 nieasurenieuts aud increasing the time efficiencv of observations.

The ability to use RV and astrometric measurements simultaneously makes it possible to employ observational timelines below the orbital ones and still be able to make positive exoplanet detections, thus helping to extract the maximum amount of information from measurements and increasing the time efficiency of observations.

Tn a forthcoming study. we plan to study the inclusion of additional transit-photometry neasuremoenuts to further tighten the parameter probability ceusities iu transiting scenarios.

In a forthcoming study, we plan to study the inclusion of additional transit-photometry measurements to further tighten the parameter probability densities in transiting scenarios.

Also. the approach used here should be extended tosvstenis with two or more plauctary COMPaluons.

Also, the approach used here should be extended tosystems with two or more planetary companions.

Ζι=Zo cannot he mace unless the interior docs not rotate rigidy but instead the observed J, value is influcuced by deep winds (?)..

$Z_1=Z_2$ cannot be made unless the interior does not rotate rigidy but instead the observed $J_4$ value is influenced by deep winds \citep{Militzer+08}.

Ou the other laud. as reffig: Jup,oncergshouws.Z4Zo is not a favorite option for the DET-MD based Jupiter model. since that would imply Zi«Z. or even zero.

On the other hand, as \\ref{fig:Jup_converg} shows, $Z_1\ll Z_2$ is not a favorite option for the DFT-MD based Jupiter model, since that would imply $Z_1\ll Z_{\odot}$ or even zero.

This would coutracict heavy. clement abundauce measurements ni Jupiter's atinosphere. which indicate Z4=ὃνZ..

This would contradict heavy element abundance measurements in Jupiter's atmosphere, which indicate $Z_1\geq 2\times Z_{\odot}$.

Therefore. Z422Za is nota free choice for the DET-ME based model.

Therefore, $Z_1\approx Z_2$ is a free choice for the DFT-MD based model.

Consequeuth. the very. differeut. Jupiter core Inasses obtained usiug these two ab-initio EOS are not primarily due to different assumptions about the distribution of heavy elements.

Consequently, the very different Jupiter core masses obtained using these two ab-initio EOS are not primarily due to different assumptions about the distribution of heavy elements.

In order to resolve this problem. a comparison of the II/Ile adiabats is highly desirable.

In order to resolve this problem, a comparison of the H/He adiabats is highly desirable.

Furthermore. iscutropic compression experiments in the 0.5—[ Mbar regime where J» and J, are most scusitive to metallicty would ereatlv help to discriminate between compoetiug Jupiter models.

Furthermore, isentropic compression experiments in the $0.5-4$ Mbar regime where $J_2$ and $J_4$ are most sensitive to metallicty would greatly help to discriminate between competing Jupiter models.

For the other EOS considered here. both the assumptions 4=Za or Z4xZo» give acceptable Jupiter models aud hence are a iatter of free choice.

For the other EOS considered here, both the assumptions $Z_1=Z_2$ or $Z_1\not=Z_2$ give acceptable Jupiter models and hence are a matter of free choice.

As we have seen in refsec:Jupiter... Jupiter models that use oue particular EOS can have various resulting Mu.Z4.Z2} triples.

As we have seen in \\ref{sec:Jupiter}, Jupiter models that use one particular EOS can have various resulting $\{M_{core}, Z_1, Z_2\}$ triples.

Without the constraint imuposed by Z4, however. the varicty of solutions would be iuunueuse. iucludiug Moore>UsMy and Mz>38AL).

Without the constraint imposed by $J_4$ however, the variety of solutions would be immense, including $M_{core}>18\ME$ and $M_Z>38\ME$.

For extrasolar plaucts. curreutly available gravity data are M, aud Fe, only,

For extrasolar planets, currently available gravity data are $M_p$ and $R_p$ only.

A potentially observational constraint equivalent to Jy has been sugeested to be the tidal Love umber &» C2)... which measures the ability of a planet to develop an elliptic deformation in response to a tidal perturber such as the close-by parent star.

A potentially observational constraint equivalent to $J_2$ has been suggested to be the tidal Love number $k_2$ \citep{RW09}, which measures the ability of a planet to develop an elliptic deformation in response to a tidal perturber such as the close-by parent star.

In this section. we present implications of knowing a precise Ay value of the Neptune-sized extrasolar planet CJ 136b.

In this section, we present implications of knowing a precise $k_2$ value of the Neptune-sized extrasolar planet GJ 436b.

About 30 light vears away from Earth. the Hot Neptune 136b (A,=23.17Mj. RB,=£22 R4) (7) orbits the Mestar 136.

About 30 light years away from Earth, the Hot Neptune $\:$ 436b $M_p=23.17\ME$, $R_p=4.22\RE$ ) \citep{Torres+08} orbits the M-star $\:$ 436.

Placed ia a imiass-racdius diagram. this planet is located close to theoretical AZ-R relations of wari water planets.

Placed in a mass-radius diagram, this planet is located close to theoretical $M$ $R$ relations of warm water planets.

Accordingly. interior structure models assmuius au irou-silicate core. a water laver. ind a ΠΠΠο euvelope allow for a composition of water and U/Ue (?).. but also for a water-less composition.

Accordingly, interior structure models assuming an iron-silicate core, a water layer, and a H/He envelope allow for a composition of water and H/He \citep{Figueira+09}, but also for a water-less composition.

I/We cuvelope models. whether dry or not. are found to have 0.02<hy0,2 (2)..

H/He envelope models, whether dry or not, are found to have $0.02<k_2<0.2$ \citep{N-GJ436b+10}.

We here hy=0.2 and search for the core mass aud water conteut that satisfy this additional coustraint.

We here $k_2=0.2$ and search for the core mass and water content that satisfy this additional constraint.

We allow water be mixed homogencously iuto the IT/IIe euvelope (two-laver model) or confined to a deep water laver (three-laver models) aud combinations iu between (Neptune-like threc-Iaver models).

We allow water be mixed homogeneously into the H/He envelope (two-layer model) or confined to a deep water layer (three-layer models) and combinations in between (Neptune-like three-layer models).

Figure 2. shows the result in comparison with LMC-REOS based Jupiter models forced to meet the observed ο value. but uot J,

Figure \ref{fig:J2k2} shows the result in comparison with LM-REOS based Jupiter models forced to meet the observed $J_2$ value, but not $J_4$.

Iu both cases. the solutions are located along a straight line (solutions for different A5 or J5 values would span parallel lines).

In both cases, the solutions are located along a straight line (solutions for different $k_2$ or $J_2$ values would span parallel lines).

For both planets. aud hence in ecueral. the upper Iit i core mass 1s obtained for homogencous envelope models (compare

For both planets, and hence in general, the upper limit in core mass is obtained for homogeneous envelope models (compare

ancl processed the data with their iniproved. version of the WD code.

and processed the data with their improved version of the WD code.

More recently. a reanalysis of the Andersenctal.(1988). RV's was obtained by Ixaranii&Alohehi(2007) witha further improvement of the spectroscopic parameters.

More recently, a reanalysis of the \citet{and88} RV's was obtained by \citet{kar07} with a further improvement of the spectroscopic parameters.

Llowever. their solution can be considered as disputable since they seem not to have included several uncertainties (of the orbital period. for example) into the total error budget and thus presumably underestimated the final errors.

However, their solution can be considered as disputable since they seem not to have included several uncertainties (of the orbital period, for example) into the total error budget and thus presumably underestimated the final errors.

Our solution. based. on our RV data ancl ASAS photometry is compared with the ones from Andersenctal.(1088) and Miloneetal.(1992). in Table 4.

Our solution, based on our RV data and ASAS photometry is compared with the ones from \citet{and88} and \citet{mil92} in Table \ref{com_aiphe}.

In. general our results are in agreement with the previous papers.

In general our results are in agreement with the previous papers.

One should note that having only S RV measurements [ου every component. we were able to improve the spectroscopic results of Andersenetal.(1988). who hac46 datapoints for every component. by a factor of 9 to 4.

One should note that having only 8 RV measurements for every component, we were able to improve the spectroscopic results of \citet{and88}, who had 46 datapoints for every component, by a factor of 2 to 4.

Unfortunately. having only one light-curve from ASAS we cannot compete with the results of Miloneetal.(1992) who used light curves from 12 bandpasses: most of which were more precise than ours.

Unfortunately, having only one light-curve from ASAS we cannot compete with the results of \citet{mil92} who used light curves from 12 bandpasses; most of which were more precise than ours.

Still. their phase coverage is not complete and the orbital and physical parameters might. be improved. with high-precision photometry.

Still, their phase coverage is not complete and the orbital and physical parameters might be improved with high-precision photometry.

Our solution converged to a somewhat different value of the orbital inclination than Ancersen’s and Alillone’s.

Our solution converged to a somewhat different value of the orbital inclination than Andersen's and Millone's.

This is of course the reason for the discrepancy in absolute masses between the solutions.

Our. AZsin’?/ is however Fu more precise.

Our $M\sin^3i$ is however far more precise.

“Phus we may conclude that this systems parameters can be derived with an unprecedented precision (possibly ~0.01 in masses and ~0.1 in radi) if only one had. more iocine RY measurements and millimagnitude photometry.

Thus we may conclude that this systems' parameters can be derived with an unprecedented precision (possibly $\sim 0.01$ in masses and $\sim 0.1$ in radii) if only one had more iodine RV measurements and millimagnitude photometry.

UX Alen. as well as Al Phe. was discovered to be à variable in Bamberg patrol plates and reported by Strohmeier (1966).

UX Men, as well as AI Phe, was discovered to be a variable in Bamberg patrol plates and reported by \citet{str66}.

. The first orbital solution was obtained by Lambert(1974) and the first full physical analysis. based on Imboerts radial velocities and photometry. was done w Clausen&CGronbech.(1976).

The first orbital solution was obtained by \citet{imb74} and the first full physical analysis, based on Imberts' radial velocities and photometry, was done by \citet{cla76}.

. Later. this photometric dataset was reanalised together with the new COILAVEL racial velocities by Andersenetal.(1989). who obtained he most up-to-date solution for UX Alen.

Later, this photometric dataset was reanalised together with the new CORAVEL radial velocities by \citet{and89} who obtained the most up-to-date solution for UX Men.

Comparison of his solution with our results is shown in Table 5..

Comparison of this solution with our results is shown in Table \ref{com_uxmen}.

Again. our results are in general in agreement with Andersen's.

Again, our results are in general in agreement with Andersen's.

As for Al Phe. having 58 radial velocity measurements of UX Alen. we succeeded to reach a better precision in the spectroscopic parameters than Andersen et al.

As for AI Phe, having 8 radial velocity measurements of UX Men, we succeeded to reach a better precision in the spectroscopic parameters than Andersen et al.

with 29 datapoints for the primary and 31 for the secondary.

In our model we fixed the eccentricity to 0.

We found no signillicant improvement in the best-fitting mocel (in terms of rms) when e and o were set as free pramoeters and the resulting e was undistinguishable from0 within the formal errors.

We found no signifficant improvement in the best-fitting model (in terms of $rms$ ) when $e$ and $\omega$ were set as free prameters and the resulting $e$ was undistinguishable from 0 within the formal errors.

The non-zero eccentricity may be however induced by a putative third body. found in NAC'O images by 'l'okovininctal.(2006) about 0.751arcsec from the binary.

The non-zero eccentricity may be however induced by a putative third body, found in NACO images by \citet{tok06} about 0.751arcsec from the binary.

V415. Aql was first reported as a variable by HLolfmeister (1936).

V415 Aql was first reported as a variable by \citet{hof36}.

. The first ephemeris 2428670.532 | [7 2.4628 d was determined by Guthnick&Schneller(1936

The first ephemeris 2428670.532 + $E \cdot$ 2.4628 d was determined by \citet{gut39}.

) The double-lined character was not. previously reported thus there was no racial velocity curve obtained to date and the only known ight curve analysis was done by Brancewiez&Dworak(1980) who derived. an orbital period. of 2.46273 d. Our UN data. despite having only 3 measurements for. every component. clearly shows that the actual period is (wo times onger.

The double-lined character was not previously reported thus there was no radial velocity curve obtained to date and the only known light curve analysis was done by \citet{bra80} who derived an orbital period of 2.46273 d. Our RV data, despite having only 3 measurements for every component, clearly shows that the actual period is two times longer.

This is confirmed by the ASAS light curve exhibiting unequal eclipses (see Eig. 12).

This is confirmed by the ASAS light curve exhibiting unequal eclipses (see Fig. \ref{fig_master}) ).

Η divided by 2. our period is in a good agreement with the value from Brancewiez&Dworak (1980).

If divided by 2, our period is in a good agreement with the value from \citet{bra80}.

. Unfortunately. the uncertainties of their results are unavailable.

Unfortunately, the uncertainties of their results are unavailable.

In Fable 6 we compare their results with ours.

In Table \ref{com_v415aql} we compare their results with ours.

The improvement in the derived parameters is obvious.

Also he distance estimate given by Brancewicz&Dworak(1980 of this should be now treated with caution.

Also the distance estimate given by \citet[][see Tab. \ref{tab_info} of this should be now treated with caution.

Out of the eighteen eclipsing binaries in our sample. 15 are

Out of the eighteen eclipsing binaries in our sample, 15 are

represents illumination of gas bv the X-ray rotating beams in à region between the accretion column ancl the discs transition region.

represents illumination of gas by the X-ray rotating beams in a region between the accretion column and the disc's transition region.

Each pulse may then represent reprocessed. gas slightly further out. from. the continuum. region which is supported. by the phase delays of 0.1 spin eveles.

Each pulse may then represent reprocessed gas slightly further out from the continuum region which is supported by the phase delay of 0.1 spin cycles.

For zero phase delay. the continuum and. emission regions would lic in the same axis (racial infall).

For zero phase delay, the continuum and emission regions would lie in the same axis (radial infall).

For a yhase delav of 0.25 eveles. the corotation velocity. of the emission region would be perpencicular to the racial infall of he continuum region.

For a phase delay of 0.25 cycles, the corotation velocity of the emission region would be perpendicular to the radial infall of the continuum region.

Therefore. the emission region ies in a region between corotation and radial infall. aud most likely. tracing the magnetic field lines.

Therefore, the emission region lies in a region between corotation and radial infall, and most likely, tracing the magnetic field lines.

Such a delay is consistent with the field lines being swept back by the »erturbed Dow of disc matter.

Such a delay is consistent with the field lines being swept back by the perturbed flow of disc matter.

The nearly svmmetrie pulses in the continuum indicate accreting poles of equal brightness at anti-diamoetric ocations.

The nearly symmetric pulses in the continuum indicate accreting poles of equal brightness at anti-diametric locations.

The pulses are very asvmimetric with the irst pulse dominating over the second one (lux is higher by a factor of 2).

The pulses are very asymmetric with the first pulse dominating over the second one (flux is higher by a factor of 2).

Phe width of the pulses is identical. 318417 xIM n. which. indicates. à common OPLgL.M and sets upper imits for the broadening mechanism in the emission-region. perhaps the scale of turbulence (change from circular motion to quasi-racial) or the sound velocity (2<LO” Ix).

The width of the pulses is identical, $\pm$ 17 km $^{-1}$, which indicates a common origin, and sets upper limits for the broadening mechanism in the emission-region, perhaps the scale of turbulence (change from circular motion to quasi-radial) or the sound velocity $T < 10^{6}$ K).

While one pulse is in maximum. blueshift. the other is in maximum recshift ("corkscrew) which is entirely consistent with an accretion curtain.

While one pulse is in maximum blueshift, the other is in maximum redshift (“corkscrew”) which is entirely consistent with an accretion curtain.

“Phe maximum of the first. pulse is viewed. directly. [rom the (lower) pole whilst the curtain Mocks any visibility of theLL region from the upper »ole. resulting in the minimum of the second pulse (in the average-subtracted line. profiles).

The maximum of the first pulse is viewed directly from the (lower) pole whilst the curtain blocks any visibility of the region from the upper pole, resulting in the minimum of the second pulse (in the average-subtracted line profiles).

This varving view of the wo inner parts of the curtain explains the quaclrapole-like xutern of the pulsed. emission in Fig.

This varying view of the two inner parts of the curtain explains the quadrapole-like pattern of the pulsed emission in Fig.

10.

Partial overlap of the pulses between spin. phases 0.50.7 suggests that the azimuthal extent is large enough so that the two pulse ocations can be viewed simultaneously at times.

Partial overlap of the pulses between spin phases 0.5–0.7 suggests that the azimuthal extent is large enough so that the two pulse locations can be viewed simultaneously at times.

The phase range of the pulses also suggests that the azimuthal extent of the accretion curtain in the continuum is smaller than hat of the emission regions (Lig.

The phase range of the pulses also suggests that the azimuthal extent of the accretion curtain in the continuum is smaller than that of the emission regions (Fig.

7).

We associate the emission-line pulsations with material which has lost its Keplerian Dow due to the magnetic fields coupling and is most likely close to radial free-fall along he magnetic field lines (see previous discussion about phase delay).

We associate the emission-line pulsations with material which has lost its Keplerian flow due to the magnetic field's coupling and is most likely close to radial free-fall along the magnetic field lines (see previous discussion about phase delay).

Then. theLL emission region will be located within he magnetosphere of the white dwarf and will rotate with he spin period of 545 seconds.

Then, the emission region will be located within the magnetosphere of the white dwarf and will rotate with the spin period of 545 seconds.

Εις gives|l2xdi -—gd./sn;--uAkms where Vie.-=408 km 17. 2?=545 seconds is. the spin. »eriod.

This gives$ V = V_{obs}/~\sin~i = \frac{2~\pi~R}{P} $ km $^{-1}$, where $V_{obs} = 408$ km $^{-1}$, $P =545$ seconds is the spin period.

"There is no eclipse in the X-ray and. optical light curves.

There is no eclipse in the X-ray and optical light curves.

On the other hand. the two accreting poles of the white chvarl (double-peak pulse) can be viewed only in hieh-inclination systenis. (50, 70°) and this is most likely the case for RX JO558|5353.

On the other hand, the two accreting poles of the white dwarf (double-peak pulse) can be viewed only in high-inclination systems $50^{o}-70^{o}$ ) and this is most likely the case for RX J0558+5353.

Phe pulsations become visible only for favourable orientations of the rotating scarchlight beans to the line-of-sight.

The pulsations become visible only for favourable orientations of the rotating searchlight beams to the line-of-sight.

Truly. this is supported by a correlation xtween inclination andpulse amplitude (Ποτος and Mason 1990).

Truly, this is supported by a correlation between inclination andpulse amplitude (Hellier and Mason 1990).

For a free-fall velocity of 408 kim sn60°=471 kms +. the distance is οlot km or RY4.9RewAly? from the white chwarl. assuming an inclination of 60" and using Bag70.84totAZ,Ob" lm (Llamacla Salpeter 1901).

For a free-fall velocity of 408 km $^{-1}$ / $\sin~60^{o} = 471$ km $^{-1}$ , the distance is $ R \sim 4.1 \times 10^{4}$ km or $ R \sim 4.9~R_{wd}~M^{+0.6}_{0.6} $ from the white dwarf, assuming an inclination of $^{o}$ and using $R_{wd} \sim 0.84 \times 10^{4}~M^{-0.6}_{0.6} $ km (Hamada Salpeter 1961).

Τμ may well represent the magnetospheric radius where the disc. ds. disrupted. [I—B4. hall the Alfven radius. which gives Hoa~O26Row(ANLapiiy)dsLot km where Ly; is the luminosity in units of 107 eres s.! and pao is the magnetic moment in units of 107 € em? (Ghosh and Lamb LOTS).

This may well represent the magnetospheric radius where the disc is disrupted, $R=R_{mag}$, half the Alfven radius, which gives $ R_{mag} \sim 0.26 ~R_{wd}~ ( M_{0.6}^{0.91}~~L_{33}^{-2/7}~ \mu_{30}^{4/7}) = 4.1 \times 10^{4}$ km where $L_{33}$ is the luminosity in units of $10^{33}$ ergs $^{-1}$ and $\mu_{30}$ is the magnetic moment in units of $10^{30}$ G $^{3}$ (Ghosh and Lamb 1978).

The impact of the stream overflow on the discs transition region is slightlv further out at ⋅⋅UST ≼≻⋅↖∖↻∫↘⋟⋯∣⇀∪⊓⊽⊓↾⊽−↿∖↙∫∶≱⋅⋅↱≻⊐⋜⋯∠⇂⊔⋯∙∖⇁↓↥⋜↧∖⇁∢⊾↓∪∖∖⋎∢⊾↓⋅∖⇁⋖⋅⇂⋯∙↓↿⊓⋅≱∖∪⇂ MN ∿∶∫≻⋅↱≻∪↓∡⊔

The impact of the stream overflow on the disc's transition region is slightly further out at $6.86~R_{wd}~M_{0.6}^{-0.27}$ $q=0.5$ ) and may have lower velocities of $\sim$ 350 km $^{-1}$.

↓⊳∖⊥⊳↾∐↥∢⋅≼∙⋖⋟↓⋅∢≱∣⋜↧⇂↕⋖⋟↓↕↓⋅⋯∐⊔⊳∖∫⊐⋟⋯↿∖↓⊲↸↓⋅

The corotation radius $R_{co}$ (Frank et al.

⋜⋯↥⊔⊾↿⋜↧↓⊳↓≤⋗≤⊔⊐⋡ Peo~10RewMUTORF isH larger than the magnetosphere (slow rotator) for accretion to occur.

1992), $R_{co} \sim 10~R_{wd}~M^{-0.27}_{0.6}$, is larger than the magnetosphere (slow rotator) for accretion to occur.

αρα et al. (

Haberl et al. (

1994) derive an observed Iuminosity of 0.18107(—L300y eres + from the N-rav Hlux between 0.1-2.4. keV. Fittingpe the X-ray spectrum with a blackbody of 57 eV. (soft component) and a bremsstrahlung spectrum of 10. keV. (hard component) gives a luminosity of 1.710°ο... eres + for an absorption column density of 6107" LE em.? (LHaberl et al.

1994) derive an observed luminosity of $0.18~\times~10^{32}~ (\frac{d}{300~{\rm pc}})^{2}$ ergs $^{-1}$ from the X-ray flux between 0.1-2.4 keV. Fitting the X-ray spectrum with a blackbody of 57 eV (soft component) and a bremsstrahlung spectrum of 10 keV (hard component) gives a luminosity of $1.7~\times~10^{33}~(\frac{d}{300~{\rm pc}})^{2}$ ergs $^{-1}$ for an absorption column density of $6~\times~~10^{20}$ H $^{-2}$ (Haberl et al.

1994).

‘Phis can be used to estimate the magnetic moment of the white dwarl of 2.407 G em? (or B0.5 MG). assuming a Ao= 1. Ly;= 1. d=300 pe and a negligible softX-ray luminosity compared to the hard. X-ray component.

This can be used to estimate the magnetic moment of the white dwarf of $2.4~\times~10^{32}$ G $^{3}$ (or $B \sim 0.5$ MG), assuming a $M_{0.6} = 1$ , $L_{33}=1$ , $d = 300$ pc and a negligible softX-ray luminosity compared to the hard X-ray component.

Such a magnetic moment is not inconsistent with the existence of an accretion disc. evidence for which we found. from the periodogram

Such a magnetic moment is not inconsistent with the existence of an accretion disc, evidence for which we found from the periodogram

is Ἱπαάσσα at 2=0.157 aud hence likely to be the 21-cim absorber.

is indeed at $z = 0.437$ and hence likely to be the 21-cm absorber.

The quasar field was imaged using the EFOSC? iustrument on the ESO 3.611 telescope at La Silla duriug the night of 29th Jaunary. 2001.

The quasar field was imaged using the EFOSC2 instrument on the ESO 3.6m telescope at La Silla during the night of 29th January, 2001.

The Bessel BR-buik 1inage was obtained on CCD #110 in the uubinned mode (0.157 aresec/pix. 2060 « 2060 pixels = 5.1 « 5.1 arcmin feld).

The Bessel R-band image was obtained on CCD 40 in the unbinned mode (0.157 arcsec/pix, 2060 $\times$ 2060 pixels $\equiv$ 5.4 $\times$ 5.4 arcmin field).

We obtained a total of teu dithex inages of 150s cach.

We obtained a total of ten dithered images of 450s each.

The calibration observations iucluded twilielr sky flats aud observations of the stadia fell RUI52 (Landolt 1992)).

The calibration observations included twilight sky flats and observations of the standard field RU152 \cite{landolt92}) ).

We used the seieuce frames ο construc and subtract the fringe pattern from the frames.

We used the science frames to construct and subtract the fringe pattern from the frames.

The atmospheric extinction and the CCD colour term were calculated using observations of standard fields taken at several airimasscs and in other bands.

The atmospheric extinction and the CCD colour term were calculated using observations of standard fields taken at several airmasses and in other bands.

We expect the photometry to be accurate to about and use a conservative error of 0.1 mae.

We expect the photometry to be accurate to about and use a conservative error of 0.1 mag.

The images were reduced. calibrated. registered aud co-added m a standard mauner using IRAF.

The images were reduced, calibrated, registered and co-added in a standard manner using IRAF.

The final image obtained had a zero-point of R=32.601 with au RAIS of 10.5 counts/pix.

The final image obtained had a zero-point of R=32.601 with an RMS of 10.5 counts/pix.