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For a one-component plasma (OCT) where Z; is the atomic number. € is the electron charge and r;=(4a.N,/31)|? is the ion-sphere radius. | For a one-component plasma (OCP) where $Z_i$ is the atomic number, $e$ is the electron charge and $r_i=(4\pi N_i/3V)^{-1/3}$ is the ion-sphere radius. |
For a mixture. we replace this by where the sums refer to a fully ionized plasma mixture with mass fractions νι | For a mixture, we replace this by where the sums refer to a fully ionized plasma mixture with mass fractions $X_i$. |
Again. for a one component plasma. where zy; is the plasma frequency. | Again, for a one component plasma, where $\omega_{pi}$ is the plasma frequency. |
For à mixture. we replace this by where N4 is Avogadro's number. | For a mixture, we replace this by where $N_A$ is Avogadro's number. |
The expression for Focp takes different forms for the eas-liquid aud for the solid phases. aud is based on the work of ?.. | The expression for $\FOCP$ takes different forms for the gas-liquid and for the solid phases, and is based on the work of \citet{1992ApJ...388..521I}. . |
Noting Chat the OCP form of the translational part (that is. setting Q;=1 and omilting the V; s) of the ionic [ree energv 35ΕΤΤ.Α} is | Noting that the OCP form of the translational part (that is, setting $Q_i=1$ and omitting the $\chi_i$ 's) of the ionic free energy $\sum F(T,V,N_I)$ is |
energies and ejecta masses typical of classical novae (i.e. in the ranges 105«Μα«107 Mo, 1076<Eo10%” erg; e.g. ?)), the peak of X-ray emission can be reached at later times. | energies and ejecta masses typical of classical novae (i.e. in the ranges $10^{-5}
< M\rs{ej} < 10^{-4}$ $M_\odot$, $10^{46}< E_0 < 10^{47}$ erg; e.g. \citealt{2010AN....331..160B}) ), the peak of X-ray emission can be reached at later times. |
Among these models, that which best reproduces the X-ray peak at day 30 is E46.6-NW9 (orange line in the left panel in Fig. 6)), | Among these models, that which best reproduces the X-ray peak at day 30 is E46.6-NW9 (orange line in the left panel in Fig. \ref{fig6}) ), |
although it fails in reproducing the descending slope of the observed lightcurve. | although it fails in reproducing the descending slope of the observed lightcurve. |
This model predicts outburst energy of the order of zz5x10*° erg and mass of ejecta c5x107° Mo that seem to be too high in the case of V407 Cyg that shows many similarities with the novae RS Oph and U Sco. | This model predicts outburst energy of the order of $\approx 5\times 10^{46}$ erg and mass of ejecta $\approx 5\times
10^{-5}$ $M_\odot$ that seem to be too high in the case of V407 Cyg that shows many similarities with the novae RS Oph and U Sco. |
In the presence of a CDE, the X-ray luminosity mainly depends on the physical characteristics of the density enhancement: the greater the size and gas density of the CDE, the larger Lx (central and right panels in Fig. 5)). | In the presence of a CDE, the X-ray luminosity mainly depends on the physical characteristics of the density enhancement: the greater the size and gas density of the CDE, the larger $L\rs{X}$ (central and right panels in Fig. \ref{fig5}) ). |
In these models, the blast reaches its maximum X-ray luminosity later than in NOCDE models, and the time of maximum depends on the size and gas density of the CDE and on the explosion energy (and ejecta mass). | In these models, the blast reaches its maximum X-ray luminosity later than in NOCDE models, and the time of maximum depends on the size and gas density of the CDE and on the explosion energy (and ejecta mass). |
In particular, the time delay of maximum X-ray luminosity is longer for lower gas density and greater size of the CDE, and for higher explosion energy. | In particular, the time delay of maximum X-ray luminosity is longer for lower gas density and greater size of the CDE, and for higher explosion energy. |
Among the SPHCDE models, that which best reproduces the observations is E44.3-NWL40 (blue line in the right panel of Fig. 6)): | Among the SPHCDE models, that which best reproduces the observations is E44.3-NW7-CDE6.3-L40 (blue line in the right panel of Fig. \ref{fig6}) ): |
it describes a continuous rise of X-ray luminosity up to a maximum of Lx£:2x10?4 erg s! at day 30 and then a subsequent decay phase until day 60, with the same slope as the observed lightcurve. | it describes a continuous rise of X-ray luminosity up to a maximum of $L\rs{X} \approx 2\times
10^{34}$ erg $^{-1}$ at day 30 and then a subsequent decay phase until day 60, with the same slope as the observed lightcurve. |
It is worth noting that this model assumes outburst energies and ejecta masses in the range of values typically observed for recurrent novae, while the “best-fit” NOCDE model assumes values typical of classical novae. | It is worth noting that this model assumes outburst energies and ejecta masses in the range of values typically observed for recurrent novae, while the “best-fit” NOCDE model assumes values typical of classical novae. |
For E44.3-NW7-CDE6.3-L40, we extended the simulation in order to cover the whole period in which the observed lightcurve is defined. | For E44.3-NW7-CDE6.3-L40, we extended the simulation in order to cover the whole period in which the observed lightcurve is defined. |
We found that, after day 60, the synthetic lightcurve fades faster than the observations which seem to be characterized by a plateau between days 60 and 90. | We found that, after day 60, the synthetic lightcurve fades faster than the observations which seem to be characterized by a plateau between days 60 and 90. |
This discrepancy could be explained as being due to the details of the density structure of the CDE; for instance, a radial density profile of the CDE decreasing away from the companion star slower than that modelled here may change the descending slope of the lightcurve. | This discrepancy could be explained as being due to the details of the density structure of the CDE; for instance, a radial density profile of the CDE decreasing away from the companion star slower than that modelled here may change the descending slope of the lightcurve. |
Another possible cause of the discrepancy may be some departures from the simple wind density profile adopted here. | Another possible cause of the discrepancy may be some departures from the simple wind density profile adopted here. |
'The effect of the shape of the CDE on the synthetic X-ray image and lightcurve was explored by performing a 3D simulation describing a disk-like CDE (run 3D-E44.3- see right panels in Fig. 3)) | The effect of the shape of the CDE on the synthetic X-ray image and lightcurve was explored by performing a 3D simulation describing a disk-like CDE (run 3D-E44.3-NW7-CDE6.3; see right panels in Fig. \ref{fig3}) ) |
with the same parameters of run E44.3-NW7-CDE6.3-L40 but with lg=53 and l3=27 in Eq. | with the same parameters of run E44.3-NW7-CDE6.3-L40 but with $l_1=l_2=53$ and $l_3=27$ in Eq. |
1 (see Table 1)). | \ref{csm}
(see Table \ref{t:params}) ). |
In this case, again most of the X-ray emission originates in the high-temperature shocked plasma in the wake of the companion star (see right panels in Fig. 5)), | In this case, again most of the X-ray emission originates in the high-temperature shocked plasma in the wake of the companion star (see right panels in Fig. \ref{fig5}) ), |
although the X-ray source here is more compact than in run E44.3-NW7-CDE6.3-L40. | although the X-ray source here is more compact than in run E44.3-NW7-CDE6.3-L40. |
The X-ray lightcurve derived from run 3D-EA4.3-NW7-CDE6.3 is illustrated in the right panel of Fig. | The X-ray lightcurve derived from run 3D-E44.3-NW7-CDE6.3 is illustrated in the right panel of Fig. |
6 (magenta line): the curve is similar to that derived from run E44.3-NW7-CDE6.3-L40 and is in good agreement with the observed lightcurve. | \ref{fig6} (magenta line): the curve is similar to that derived from run E44.3-NW7-CDE6.3-L40 and is in good agreement with the observed lightcurve. |
We have modeled the first 60 days of evolution of the 2010 blast wave of the nova V407 Cyg. | We have modeled the first 60 days of evolution of the 2010 blast wave of the nova V407 Cyg. |
The models explore the 3D structure of the blast wave originating from the nova outburst, taking into account simultaneously the radiative cooling and thermal conduction (including heat flux saturation). | The models explore the 3D structure of the blast wave originating from the nova outburst, taking into account simultaneously the radiative cooling and thermal conduction (including heat flux saturation). |
Our analysis indicates the following: ? note that the 2010 outburst of V407 Cyg closely resembles the spectroscopic development of RS Oph which is considered the prototype of symbiotic-like recurrent novae. | Our analysis indicates the following: \cite{2011A&A...527A..98S} note that the 2010 outburst of V407 Cyg closely resembles the spectroscopic development of RS Oph which is considered the prototype of symbiotic-like recurrent novae. |
Long before massive black holes were accepted to be present iu the centers of galaxies. argued tiab a binary black hole ereated it the mereer of two galaxies would eject sals from the center of the newly created system as the binary slowly hardeued. | Long before massive black holes were accepted to be present in the centers of galaxies, \citet{bbr} argued that a binary black hole created in the merger of two galaxies would eject stars from the center of the newly created system as the binary slowly hardened. |
N-bocly simulatiOlls support this hypothesis (Ebisuzalietal.1991:Makino|LOOT:Milosavljevi6&Merritt2001). | N-body simulations support this hypothesis \citep{ebi, mak97, mnm}. |
. We LOW believe that nearly every eliptical galaxy or spiral bulee has a lack liole at its center 1998).. | We now believe that nearly every elliptical galaxy or spiral bulge has a black hole at its center \citep{mag}. |
We also )elieve that he most massive elliptical gaaxies were formed by mereine galaxies. | We also believe that the most massive elliptical galaxies were formed by merging pre-existing galaxies. |
Ht is possible tluuo their central structure still bears wituess to such events. | It is possible that their central structure still bears witness to such events. |
The “cores” seen al dje centers of the 1ost luminous elliptical galaxies nav indeed be the signatures of gravitational stirring and heating ("core scouring"} N binary black holes (Faberetal.1997). | The “cores” seen at the centers of the most luminous elliptical galaxies may indeed be the signatures of gravitational stirring and heating (“core scouring”) by binary black holes \citep{f97}. |
. Cores were initially seen in ο‘ound-basecl observations of δν Luminous eliptical galaxies as a central region of uearly constaut surface brightsess (Lauer1985:Ixormendy1985). | Cores were initially seen in ground-based observations of nearby luminous elliptical galaxies as a central region of nearly constant surface brightness \citep{l85, k85}. |
. The cores were seen to beion-isothermal. but their true foru 1wero 0 was uuknuown. | The cores were seen to be non-isothermal, but their true form as $r\rightarrow0$ was unknown. |
AST ijages. however. later showed tha uearly all galaxies have sluglar 5arleht distributions in the seise that surface brightness diverges as X(r)~r (Laueretal.19911992a.b:Craneal.1993:Konjeidsy1991:Ferrareseetal.Lauer1995 ). | images, however, later showed that nearly all galaxies have singular starlight distributions in the sense that surface brightness diverges as $\Sigma(r)\sim r^{-\gamma}$ \citep{l91, l92a, l92b, crane, k94, f94, l95}. |
. Tydically. in low luminosity early-type galaxies. decreases ouly slowly as the ceuter is approached aud a steep >0.5 Cusp continues into the resolution limit: Laueretal.(1995). classi[ied tjese systetus as “power-law” galaxies. | Typically, in low luminosity early-type galaxies, $\gamma$ decreases only slowly as the center is approached and a steep $\gamma>0.5$ cusp continues into the resolution limit; \citet{l95} classified these systems as “power-law” galaxies. |
It1 more luminous galaxies. however. the steep euvelope profile trausitious to a shallow iuner cusp witli uwQ.: Sata “break radius. rg. Tus behavior was seen lu galaxies that had cores [rol «ybservatious: {1e break radius rouglly correspouded to what hac| been measurec as tle core radius. | In more luminous galaxies, however, the steep envelope profile transitions to a shallow inner cusp with $\gamma<0.3$ at a “break radius,” $r_b.$ This behavior was seen in galaxies that had cores from ground-based observations; the break radius roughly corresponded to what had been measured as the core radius. |
We thus chose to continue o call hese "core galaxies." even though the shallow cusps iu projected brightuess in of these systems imply steep aud sineular cusps in Duinosity deusity (Laueretal.1995). | We thus chose to continue to call these “core galaxies,” even though the shallow cusps in projected brightness in of these systems imply steep and singular cusps in luminosity density \citep{l95}. |
. A strong justificaion of thisclassification schema was that tle disribution of +" over a sample ol early-(ype galaxies was seen to be stroiglv bimocal (Faberetal.1997).. where >’ is the local cusp slope at theAST resolution liinit. | A strong justification of thisclassification schema was that the distribution of $\gamma'$ over a sample of early-type galaxies was seen to be strongly bimodal \citep{f97}, where $\gamma'$ is the local cusp slope at the resolution limit. |
Lhitlally uo galaxies were seei1 {ο have Q0: Jit was thiis sensible to apply the separate core and power-law classificatious to the two distinct and cleanly separated peaks in the +! distribution. | Initially no galaxies were seen to have $0.3<\gamma'<0.5.$ It was thus sensible to apply the separate core and power-law classifications to the two distinct and cleanly separated peaks in the $\gamma'$ distribution. |
The bimodal distributioi was initially inferred frou. fits to he surface photometry usit& the "Nuker law” (Laueretal.1995:: see equation 2) . which has the asvinptotic form X(r)er "Las {το0. however. Gebhardteta.(1996) demoustraed that uon-pa"metric treatment of tlie 5ulace photometry recovered au klentical bimodal distribution. and futher that bimocality was <| feature of the luminosity density distribuious of ea‘ly type ealaxies. not just the surface briehness distributions. | The bimodal distribution was initially inferred from fits to the surface photometry using the “Nuker law” \citealt{l95}; see equation \ref{eqn:nuker}) ), which has the asymptotic form $\Sigma(r)\sim r^{-\gamma},$ as $r\rightarrow0,$ however, \citet{g96} demonstrated that non-parametric treatment of the surface photometry recovered an identical bimodal distribution, and further that bimodality was a feature of the luminosity distributions of early type galaxies, not just the surface brightness distributions. |
ΤΙe core aud power-law classifications are fundamental because tvey correlate with global | The core and power-law classifications are fundamental because they correlate with global |
PUDMy) for the two models. which permit the evaluation of the Daves factor and hence (he oclds ratio via (8)). | $P(D|{\cal M}_{\rm pl})$ for the two models, which permit the evaluation of the Bayes factor and hence the odds ratio via \ref{eq:odds_ratio_body}) ). |
For the 56 GOES peak [Iuxes associated. with active region. 11029 and above size 54 the odds ratio evaluates to νι)2:220. assuming a unity. prior odds ratio (pyp;pi=1). | For the 56 GOES peak fluxes associated with active region 11029 and above size $S_1$ the odds ratio evaluates to $r_{\rm plr/pl}(D)\approx 220$, assuming a unity prior odds ratio $\rho_{\rm plr/pl}=1$ ). |
Oclds ratios are often presented in decibels (dD) and in this case the odds ratio in decibels is 101οςrg;(2)z23dB. | Odds ratios are often presented in decibels (dB) and in this case the odds ratio in decibels is $10\log_{10} r_{\rm plr/pl}(D)\approx 23\,{\rm dB}$. |
This result implies that. if both models are assumed a priori to be equally likely. then the data favors the power-law plus rollover model by a factor οἱ more than 200. | This result implies that, if both models are assumed a priori to be equally likely, then the data favors the power-law plus rollover model by a factor of more than 200. |
An odds ratio of this magnitude may be interpreted as strong evidence for the favored model JJeffrevs 1961: Javnes 2003: see also the discussion at the end of Appendix C). and the result confirms the qualitative impression given by the lower panel ol Figure 4. | An odds ratio of this magnitude may be interpreted as strong evidence for the favored model Jeffreys 1961; Jaynes 2003; see also the discussion at the end of Appendix C), and the result confirms the qualitative impression given by the lower panel of Figure 4. |
Of course. if the simple power-law model is thought a priori to be strongly prelerable (pipi< 1). then the odds ratio is reduced. according to equation (3)). | Of course, if the simple power-law model is thought a priori to be strongly preferable $\rho_{\rm plr/pl} \ll 1$ ), then the odds ratio is reduced, according to equation \ref{eq:odds_ratio_body}) ). |
However. al [ace value the data for active region 11029 provides strong evidence for the power-law plus rollover model over the simple power-law moclel. | However, at face value the data for active region 11029 provides strong evidence for the power-law plus rollover model over the simple power-law model. |
The specific value of the odds ratio depends on the choice of the threshold 53. | The specific value of the odds ratio depends on the choice of the threshold $S_1$. |
However. the result does not depend strongly on this choice in the sense that the model with the rollover is always [avored. | However, the result does not depend strongly on this choice in the sense that the model with the rollover is always favored. |
If (his choice is made too large the odds ratio is reduced due to the resulting small event numbers. | If this choice is made too large the odds ratio is reduced due to the resulting small event numbers. |
The power-law plus rollover model] is chosen for simplicity. and it is worthwhile to consider other possible models. for example a broken power law. | The power-law plus rollover model is chosen for simplicity, and it is worthwhile to consider other possible models, for example a broken power law. |
lt is unlikely that the small set of data for AR. 11029 allows distinction between a broken power-law model and a power law with an exponential rollover. and (he calculation is not attempted. | It is unlikely that the small set of data for AR 11029 allows distinction between a broken power-law model and a power law with an exponential rollover, and the calculation is not attempted. |
Llowever. because (he Davesian model comparison involves integration over all possible values of the rollover parameter a. the specific form of the model is likely to be less important in the model comparison than the choice of a model with a departure trom power-law behavior. | However, because the Bayesian model comparison involves integration over all possible values of the rollover parameter $\sigma$, the specific form of the model is likely to be less important in the model comparison than the choice of a model with a departure from power-law behavior. |
I1 is expected Chat a model comparison between a simple power law and a broken power law will eive very comparable results for this data an odds ratio strongly in favor of the broken power law). | It is expected that a model comparison between a simple power law and a broken power law will give very comparable results for this data an odds ratio strongly in favor of the broken power law). |
The waiting-time distribution is also constructed for the 56 events with subtracted peak flux larger than 54=10*Wim7. to investigate the flaring rate. | The waiting-time distribution is also constructed for the 56 events with background-subtracted peak flux larger than $S_1=10^{-7}\,{\rm W}\,{\rm m}^{-2}$, to investigate the flaring rate. |
Figure 5 ilustrates (he analvsis. | Figure 5 illustrates the analysis. |
The upper panel shows the cumulative number of events versus lime. the middle panel shows the Bavesian blocks analvsis of (he rate versus time (Scarele | The upper panel shows the cumulative number of events versus time, the middle panel shows the Bayesian blocks analysis of the rate versus time (Scargle |
Sulphur is an a-element (like O, Ne, Mg, Si, Ar, and Ca) and these are believed to be produced mainly in SNe type II by additions of a-particles. | Sulphur is an $\alpha$ -element (like O, Ne, Mg, Si, Ar, and Ca) and these are believed to be produced mainly in SNe type II by additions of $\alpha$ -particles. |
The SNe type Ia, on the other hand, mainly produce the iron peak elements (Fe, Co, and Ni). | The SNe type Ia, on the other hand, mainly produce the iron peak elements (Fe, Co, and Ni). |
Due to the different life times of the two groups of productions sites, the interstellar medium (ISM) in the early Galaxy is believed to contain higher abundance a-elements compared to Fe than later on when the SNe type Ia start expelling iron. | Due to the different life times of the two groups of productions sites, the interstellar medium (ISM) in the early Galaxy is believed to contain higher abundance $\alpha$ -elements compared to Fe than later on when the SNe type Ia start expelling iron. |
Thus, stars that form at a certain time serve as markers of the [α/Εε[] ratio in the ISM at that particular stage in the evolution of the Galaxy. | Thus, stars that form at a certain time serve as markers of the $[\alpha / \mathrm{Fe}]$ ratio in the ISM at that particular stage in the evolution of the Galaxy. |
This requires that no extra a-elements nor Fe has been produced in the star and contaminated the photosphere during the star’s life. | This requires that no extra $\alpha$ -elements nor Fe has been produced in the star and contaminated the photosphere during the star's life. |
The theory of nucleosynthesis and stellar structure and evolution suggest this to be the case. | The theory of nucleosynthesis and stellar structure and evolution suggest this to be the case. |
Since the iron abundance, [Fe/H], roughly increases with time, an [@/Fe] vs. [Fe/H] plot is expected to show plateau for the lowest metallicities and a negative slope froma the time when the SNe type Ia start producing large amounts of iron. | Since the iron abundance, $[ \mathrm{Fe} / \mathrm{H}]$ , roughly increases with time, an $[\alpha / \mathrm{Fe}]$ vs. $[\mathrm{Fe} / \mathrm{H}]$ plot is expected to show a plateau for the lowest metallicities and a negative slope from the time when the SNe type Ia start producing large amounts of iron. |
Regarding sulphur in particular, it is believed to be produced via oxygen burning just like Si and Ca and therefore these three elements are expected to vary in lock-step with each other. | Regarding sulphur in particular, it is believed to be produced via oxygen burning just like Si and Ca and therefore these three elements are expected to vary in lock-step with each other. |
The expected flat behavior in the [@/Fe] vs. [Fe/H] plot is however not unambiguously observed for sulphur in all previous works and hence there is no agreement regarding its evolution. | The expected flat behavior in the $[\alpha / \mathrm{Fe}]$ vs. $[\mathrm{Fe} / \mathrm{H}]$ plot is however not unambiguously observed for sulphur in all previous works and hence there is no agreement regarding its evolution. |
This mismatch with the rest of the a-elements makes the behavior of sulphur interesting to investigate for its own sake, but the Galactic evolution of sulphur is also important to determine for at least two other reasons: Despite its importance, the Galactic evolution of sulphur is not a well studied subject (compared to, e.g., the Galactic evolution of oxygen). | This mismatch with the rest of the $\alpha$ -elements makes the behavior of sulphur interesting to investigate for its own sake, but the Galactic evolution of sulphur is also important to determine for at least two other reasons: Despite its importance, the Galactic evolution of sulphur is not a well studied subject (compared to, e.g., the Galactic evolution of oxygen). |
This is mainly due to the lack of suitable sulphur diagnostics. | This is mainly due to the lack of suitable sulphur diagnostics. |
During the years several diagnostics have been used with different strengths and weaknesses. | During the years several diagnostics have been used with different strengths and weaknesses. |
Sulphur, being situated directly below oxygen in the periodic table, has a similar set of energy levels and transitions as oxygen. | Sulphur, being situated directly below oxygen in the periodic table, has a similar set of energy levels and transitions as oxygen. |
For example the 1082 nm i] line and the 1045 nm triplet used in this work are analogous to the widely used i] line at 630 nm and the 845 nm lines. | For example the 1082 nm ] line and the 1045 nm triplet used in this work are analogous to the widely used ] line at 630 nm and the 845 nm lines. |
It turns out that different works, using different sulphur diagnostics, have resulted in several different scenarios for the Galactic evolution of sulphur. | It turns out that different works, using different sulphur diagnostics, have resulted in several different scenarios for the Galactic evolution of sulphur. |
From measurements of the 869 nm doublet ? and ? found a negative slope in the [S/Fe] vs. [Fe/H] plot for all observed iron abundances (-3<[Fe/H]x +0.5). | From measurements of the 869 nm doublet \citet{Israelian2001} and \citet{Takada-Hidai2002} found a negative slope in the $[\mathrm{S} / \mathrm{Fe}]$ vs. $[\mathrm{Fe} / \mathrm{H}]$ plot for all observed iron abundances $-3\leq [\mathrm{Fe} / \mathrm{H}] \leq +0.5$ ). |
To explain this they proposed either a scenario for sulphur production involving hypernovae (e.g.,?) or time-delayed iron deposition into the ISM as compared to sulphur (see ? for a description of the mechanism in connection to [O/Fe] vs. [Fe/H]). | To explain this they proposed either a scenario for sulphur production involving hypernovae \citep[e.g., ][]{Nakamura2001} or time-delayed iron deposition into the ISM as compared to sulphur (see \citet{Ramaty2000} for a description of the mechanism in connection to $[\mathrm{O} / \mathrm{Fe}]$ vs. $[\mathrm{Fe} / \mathrm{H}]$ ). |
A majority of the stars observed by ? and ? were later re-analyzed using the roughly ten times stronger triplet around 923 nm (??),, allowing a more precise abundance determination for metal-poor stars. | A majority of the stars observed by \citet{Israelian2001} and \citet{Takada-Hidai2002} were later re-analyzed using the roughly ten times stronger triplet around 923 nm \citep{Ryde2004,Korn2005}, allowing a more precise abundance determination for metal-poor stars. |
These new results, contradicting the earlier findings, suggest an evolution analogous to the rest of the a-elements. | These new results, contradicting the earlier findings, suggest an evolution analogous to the rest of the $\alpha$ -elements. |
Neither diagnostic is, however, optimal; the 923 nm triplet is stronger than the 869 nm doublet, but on the other hand it is situated in a spectral region heavily plagued with telluriclines and the middle line of the triplet is situated in the wing of the strong Paschen £ line. | Neither diagnostic is, however, optimal; the 923 nm triplet is stronger than the 869 nm doublet, but on the other hand it is situated in a spectral region heavily plagued with telluriclines and the middle line of the triplet is situated in the wing of the strong Paschen $\zeta$ line. |
vas discovered duriug the search for supersoft X-ray sources in theROSAT all-sky survey data (Cireiner. Remillard aud Motch 1995). | was discovered during the search for supersoft X-ray sources in the all-sky survey data (Greiner, Remillard and Motch 1995). |
Cremer. Remullare aud Motch (1995. 1998) lave analyzed all theROSAT data taken between 1990 September aud 1993 ober. and found a coliereut period of 0.078[7077(11) « ( = Lss35115z0.000003. hr). | Greiner, Remillard and Motch (1995, 1998) have analyzed all the data taken between 1990 September and 1993 ber, and found a coherent period of 0.07847977(11) d ( = $\pm$ 0.000003 hr). |
The pulse profile in the uuddo«0.5 keV is characterized bv a deep. inteusitv uinimuiun. with basically no X-ray flux. lasting 0.1 orbita yhase. | The pulse profile in the band $<0.5$ keV is characterized by a deep intensity minimum, with basically no X-ray flux, lasting 0.1 orbital phase. |
The X-ray spectrum is characterized bw strong dackbody cussion with a teniperature of 20+15 eV. with a clear excess cuiission above l keV. which has )en approximated by thermal brenisstraliluus with a cluperature of 20 το, The absorptiou-correced 2.1 keV fluxes of the blackbody aud the thermal sstrahhme components are τνlOHP erg em78 ! and s∖↽10DI!⊳⋅ ore ,7s |l respectively, sugecsting a huge soft excess of nearly 90 in he 0.1.2.1 keV. band. | The X-ray spectrum is characterized by strong blackbody emission with a temperature of $20\pm 15$ eV, with a clear excess emission above 1 keV, which has been approximated by thermal bremsstrahlung with a temperature of 20 keV. The absorption-corrected 0.1--2.4 keV fluxes of the blackbody and the thermal strahlung components are $7 \times 10^{-11}$ erg $^{-2}$ $^{-1}$ and $8 \times 10^{-13}$ erg $^{-2}$ $^{-1}$, respectively, suggesting a huge soft excess of nearly 90 in the 0.1–2.4 keV band. |
Szkody (1995) carrie out photometry. spectroscopy aud oolhudnetre iu the optical band. and found. several characteristics of polars such as emission Lue stronger hau IL? aud circular polarization of | Szkody (1995) carried out photometry, spectroscopy and polarimetry in the optical band, and found several characteristics of polars such as emission line stronger than $\beta$ and circular polarization of. |
All these oxoperties strouegly indicate that is a »olar. | All these properties strongly indicate that is a polar. |
Iu this paper. data of taken in 1996 Sep-Oct are presented. | In this paper, data of taken in 1996 Sep-Oct are presented. |
Iu 2 we describe how he observation was carried out. | In 2 we describe how the observation was carried out. |
Iu 5 light curve aud spectral analysis are preseuted. | In 3 light curve and spectral analysis are presented. |
We discuss these Opertics m liu combination with the ROSATaud data. | We discuss these properties in 4 in combination with the and data. |
In 5 we summarize our results. | In 5 we summarize our results. |
The observation of was carried out between 1996 September 30.75 and October 2.96 (UT). | The observation of was carried out between 1996 September 30.75 and October 2.96 (UT). |
is equipped with four equivalent N-ray Telescopes (XRT: Serlemitsos 1995). | is equipped with four equivalent X-ray Telescopes (XRT: Serlemitsos 1995). |
Iu the common focal plane. two Solid-state Tnagine Spectrometers (SIS: Burke 1991. Yamashita 1997) aud two Cas huaegiug Spectrometers (CIS: Makishima 1996. Ohashi 1996) are mounted. | In the common focal plane, two Solid-state Imaging Spectrometers (SIS: Burke 1994, Yamashita 1997) and two Gas Imaging Spectrometers (GIS: Makishima 1996, Ohashi 1996) are mounted. |
The SIS has high sensitivity iu the lower energy bandpass and hieh energy resolution of AE/E~0.02 (at the time of lamuch). whereas the CIS has high throughput in the higher cucrey bandpass and lieh time resolution. | The SIS has high sensitivity in the lower energy bandpass and high energy resolution of $\Delta E/E \simeq 0.02$ (at the time of launch), whereas the GIS has high throughput in the higher energy bandpass and high time resolution. |
Throughout the observation. the CIS operated in Pulse Ueight uormal mode in which the band 0.7-10 keV is covered by 1021 pulse height channels. | Throughout the observation, the GIS operated in Pulse Height normal mode in which the band 0.7-10 keV is covered by 1024 pulse height channels. |
The SIS mode was switched between the 1-CCD FAINT mode in high bit rate aud the 1-CCD DRIGIIT mode iu iiedium bit rate. which cover 0.1-10 keV with 096 and 2018 pulse height channels. respectively, | The SIS mode was switched between the 1-CCD FAINT mode in high bit rate and the 1-CCD BRIGHT mode in medium bit rate, which cover 0.4-10 keV with 4096 and 2048 pulse height channels, respectively. |
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