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Iu his work. we have focussed ou one particularly interesing example: a leptophilie 2HDM. in which ¢ifferent Higesao bosous are respotisible for giving masses to tlie quark aud lepton sectors.
In this work, we have focussed on one particularly interesting example: a leptophilic 2HDM, in which different Higgs bosons are responsible for giving masses to the quark and lepton sectors.
We have exaljiined t1e effect of such a modification on the collider j»Xheuomenology ol a light HiOOees boson lun a «ecoupliug regime in which the ouly light scalar is a HigesOO boson. aud have shown tiat a number of collider processes involving tie direct. decay ol the Higgs to a pair of charged leptous cau play a cruc‘ial role in its discovery.
We have examined the effect of such a modification on the collider phenomenology of a light Higgs boson in a decoupling regime in which the only light scalar is a Standard-Model-like Higgs boson, and have shown that a number of collider processes involving the direct decay of the Higgs to a pair of charged leptons can play a crucial role in its discovery.
lu particular. we have shown that there are regious of parameter space in which the Hgesoonbosou couplings to leptous can be greatly enlauced.
In particular, we have shown that there are regions of parameter space in which the Higgs-boson couplings to leptons can be greatly enhanced.
Thi scan have a potentially dramatic effect ou the Higes discovery
This can have a potentially dramatic effect on the Higgs discovery
radius.
radius.
We now give an argument (ο show that this is abwavs the case. even when the Boncli racius is smaller than the Roche lobe radius.
We now give an argument to show that this is always the case, even when the Bondi radius is smaller than the Roche lobe radius.
It is usually assumed that a protoplanet embedded in a protoplanetary dise cannot expand bevond the Bondi radius defined as: where c is the sound speed in the protoplanetary disc (e.g.. Bodenheimer et al.
It is usually assumed that a protoplanet embedded in a protoplanetary disc cannot expand beyond the Bondi radius defined as: where $c$ is the sound speed in the protoplanetary disc (e.g., Bodenheimer et al.
2000).
2000).
The argument is (hat the gas located at a distance from (he protoplanet larger (han rj has a thermal energy larger (han (he gravitational energv that would bind it to the protoplanet. and therefore cannot be accreted by it.
The argument is that the gas located at a distance from the protoplanet larger than $r_B$ has a thermal energy larger than the gravitational energy that would bind it to the protoplanet, and therefore cannot be accreted by it.
However. ifa molecule of gas located in (he protoplanetary dise bevond. ry is accelerated toward the protoplanet and collides with it. it may become bound (and therefore be accreted) like any other molecule coming from within the Roche lobe as described in section 3.1..
However, if a molecule of gas located in the protoplanetary disc beyond $r_B$ is accelerated toward the protoplanet and collides with it, it may become bound (and therefore be accreted) like any other molecule coming from within the Roche lobe as described in section \ref{sec:accretion}.
Accretion of this molecule only requires that i( loses at least part of its (kinetic plus thermal) energv into shocks during the collision.
Accretion of this molecule only requires that it loses at least part of its (kinetic plus thermal) energy into shocks during the collision.
If the molecule hits and settles into a circumplanetary disce before falling onto the planet. accretion will happen as energv is radiated away from the disc.
If the molecule hits and settles into a circumplanetary disc before falling onto the planet, accretion will happen as energy is radiated away from the disc.
When the molecule becomes bound to the protoplanet. the pressure gradient in the atmosphere adjusts itself to balance the eravitational attraction of the protoplanet. ie. hvdrostatie equilibrium is maintained.
When the molecule becomes bound to the protoplanet, the pressure gradient in the atmosphere adjusts itself to balance the gravitational attraction of the protoplanet, i.e. hydrostatic equilibrium is maintained.
At lixed Iuminositv. the protoplanet has to expand (o accommodate this extra mass. and its final surface radius does not have to be limited by ry.
At fixed luminosity, the protoplanet has to expand to accommodate this extra mass, and its final surface radius does not have to be limited by $r_B$ .
For à given mass Mj and surface temperature Z7; of the atmosphere. an equilibrium atmosphere can be constructed. with PainUB.
For a given mass $M_{\rm atm}$ and surface temperature $T_s$ of the atmosphere, an equilibrium atmosphere can be constructed with $r_{\rm atm}>r_B$.
This is exactly similar to (he process of star formation.
This is exactly similar to the process of star formation.
In a molecular cloud (hat collapses onto a star. when particles hit the surface of the star (μον have a positive energy.
In a molecular cloud that collapses onto a star, when particles hit the surface of the star they have a positive energy.
This is because they have some thermal energy. in addition to the kinetic energy. due to the acceleration by the gravitational potential of the forming star.
This is because they have some thermal energy in addition to the kinetic energy due to the acceleration by the gravitational potential of the forming star.
Accretion is possible because al least part of this enerev is dissipated into shocks.
Accretion is possible because at least part of this energy is dissipated into shocks.
The particle Chat is accreted may or mav not bring some entropy into the stellar envelope. depending on whether all or only some of the energy is racdiated away. and the star readjusts to equilibrium (Palla Stabler 1991. 1992. IIosokawaet al.
The particle that is accreted may or may not bring some entropy into the stellar envelope, depending on whether all or only some of the energy is radiated away, and the star readjusts to equilibrium (Palla Stahler 1991, 1992, Hosokawaet al.
2010).
2010).
strength of clustering within a eiven redshift interval as a fuuction of intrinsic luminosity.
strength of clustering within a given redshift interval as a function of intrinsic luminosity.
This type oL evolution might be naively expected if the Iumiuosity of a galaxy. uniquely mapped to the mass ol the dark matter halo at the given redshift in whic[un resides more luminous galaxies cluster more strongly within a given redshilt interval).
This type of evolution might be naively expected if the luminosity of a galaxy uniquely mapped to the mass of the dark matter halo at the given redshift in which it resides more luminous galaxies cluster more strongly within a given redshift interval).
Most likely. either we have too simall of a sample to place siguificant limits ou the variation of clustering with liuuinosity. or else the relatively coustaut clustering amplitude as a functiou of lunuinosity is incicative of luminosity evolution complicating the analysis.
Most likely, either we have too small of a sample to place significant limits on the variation of clustering with luminosity, or else the relatively constant clustering amplitude as a function of luminosity is indicative of luminosity evolution complicating the analysis.
Future surveys. both photometric aud spectroscopic SSDSS) wiIzll provide extremely useful datasets with which we cau explore these ideas in greater detail.
Future surveys, both photometric and spectroscopic SDSS) will provide extremely useful datasets with which we can explore these ideas in greater detail.
In the uear future. this area will wituess a mergiug of observations. seimi-aualytic theory. aud N-body simulatious. finally providiug hope that we will be able to unambiguously quautify the clustering evolutiou of galaxies.
In the near future, this area will witness a merging of observations, semi-analytic theory, and N-body simulations, finally providing hope that we will be able to unambiguously quantify the clustering evolution of galaxies.
First we wish to acknowledge Gyula Szokoly [or assistance in obtaining the cata.
First we wish to acknowledge Gyula Szokoly for assistance in obtaining the data.
We also would like to thank Barry Lasker. Gretchen Creene. aud. Brian MeLeau for allowing us access to an early version of the GSC IL.
We also would like to thank Barry Lasker, Gretchen Greene, and Brian McLean for allowing us access to an early version of the GSC II.
We also wish to acknowledge useful discussions with Pat Cote. Rich Ixrou. Lori Lubin. aud Ray Wesimanu.
We also wish to acknowledge useful discussions with Pat Cote, Rich Kron, Lori Lubin, and Ray Weymann.
We thank the anonymous referee [or valuable suggestionsMOD ou improving this work.
We thank the anonymous referee for valuable suggestions on improving this work.
This research bas made use of NASA’s LEM Data System Abstract Service.
This research has made use of NASA's Astrophysical Data System Abstract Service.
AS acknowledges support from NASA LTSA (NAC53503) aud. Grant 96A). AJC acknowledges partial support from (C1O-O7817-02-964) and LTSA (NRA-98-03-LTSA-039).
AS acknowledges support from NASA LTSA (NAG53503) and Grant (GO-07817-04-96A), AJC acknowledges partial support from (GO-07817-02-96A) and LTSA (NRA-98-03-LTSA-039).
Remarkable progress has been made over the last ten years in putting together an increasingly detailed picture of galaxy evolution since 2<4.
Remarkable progress has been made over the last ten years in putting together an increasingly detailed picture of galaxy evolution since $z<4$.
In particular. having established with some accuracy the star-formation history of the Universe (e.g.Lillyetal.1996:Hopkins&Beacom2006:Reddyetal. 2008).. the next objective is to establish how these stars were assembled over time (e.g.Dick-Barger2008:Marchesinietal. 2008).
In particular, having established with some accuracy the star-formation history of the Universe \citep[e.g.][]{Lilly,HB06,Reddy}, the next objective is to establish how these stars were assembled over time \citep[e.g.][]{dick03,cons07,bell07,cowi08,March08}.
. This mass-assembly history is. in principle. an observable quantity that can provide a robust. direct constraint on theoretical models (e.g... Boweretal. 2006).
This mass-assembly history is, in principle, an observable quantity that can provide a robust, direct constraint on theoretical models (e.g., \citealt{bowe06}) ).
One of the most generic predictions of all galaxy formation models is that the total mass in the Universe. dominated by cold dark matter (CDM. Blumenthaletal. 19843). assembles by building up progressively larger structures with time (e.g.White&Frenk 1991).
One of the most generic predictions of all galaxy formation models is that the total mass in the Universe, dominated by cold dark matter (CDM, \citealt{blum84}) ), assembles by building up progressively larger structures with time \citep[e.g.][]{wf91}.
. Observations have long shown that the most massive galaxies today actually have the stellar populations (e.g.Gal-Nelanetal.2005:Smith2008:Rettura 2008).. but this alone does not pose much «difficulty for theory if these massive galaxies were assembled early from smaller lumps of matter in which stellar populations were already established.
Observations have long shown that the most massive galaxies today actually have the stellar populations \citep[e.g.][]{gall84,ble,vD98,Nelan,Smith,rettura}, but this alone does not pose much difficulty for theory if these massive galaxies were assembled early from smaller lumps of matter in which stellar populations were already established.
More puzzling have been direct observations of high-redshift galaxies. which show that the majority of massive galaxies were already in place by >=1. and that they stopped forming new stars sooner than galaxies of lower mass (Cowieetal.1996.. Juneauetal.2005 (hereafter JOS). Fontanaetal.2004.. Bundyetal. 2006.. etal. 2008.. Tayloretal. 2008).
More puzzling have been direct observations of high-redshift galaxies, which show that the majority of massive galaxies were already in place by $z=1$, and that they stopped forming new stars sooner than galaxies of lower mass \citealt{cowi96}, \citealt{june05} (hereafter J05), \citealt{font04}, \citealt{bundy}, \citealt{moba08}, \citealt{taylor}) ).
Thus. it is of key interest to obtain a direct measurement of star-formation rate as a function of stellar mass in galaxies at different redshifts.
Thus, it is of key interest to obtain a direct measurement of star-formation rate as a function of stellar mass in galaxies at different redshifts.
To date. high-redshift measurements have been limited to the most massive or the most highly-star forming galaxies.
To date, high-redshift measurements have been limited to the most massive or the most highly-star forming galaxies.
Recent near-infrared selected spectroscopic surveys such as the Gemini Deep Deep Survey (GDDS. Abrahametal. 20043) and K20 (Fontanaetal.2004) have pushed as deep as 22.5.
Recent near-infrared selected spectroscopic surveys such as the Gemini Deep Deep Survey (GDDS, \citealt{abra04}) ) and K20 \citep{font04} have pushed as deep as $K_{AB}\simeq 22.5$ .
These are desirable as the A -band allows a clean selection to be made on approximate stellar mass out to high-redshifts. and stellar mass is a relatively robust quantity to compare with simulations (e.g.Marchesinietal.2008).
These are desirable as the $K$ -band allows a clean selection to be made on approximate stellar mass out to high-redshifts, and stellar mass is a relatively robust quantity to compare with simulations \citep[e.g.][]{March08}.
. GDDS and K20 select 105A7. (stellar mass) galaxies tos&2 and >107737. galaxies at 2.cl.
GDDS and K20 select $>10^{11}\Msun$ (stellar mass) galaxies to $z\simeq 2$ and $>10^{10}\Msun$ galaxies at $z\simeq 1$.
Spectroscopic surveys serve to provide accurate redshifts and also to measure fluxes in nebular lines such as [OII] and Ha which can be used to estimate star- rates.
Spectroscopic surveys serve to provide accurate redshifts and also to measure fluxes in nebular lines such as [OII] and $\alpha$ which can be used to estimate star-formation rates.
However. spectroscopy is generally not attempted for fainter continuum objects due to the much longer integration
However, spectroscopy is generally not attempted for fainter continuum objects due to the much longer integration
excess emission at 4.5 from hot dust was found mostly in high-surface-brightness HII regions, implying that massive stars are the primary source of heating 2006).
excess emission at 4.5 from hot dust was found mostly in high-surface-brightness HII regions, implying that massive stars are the primary source of heating .
. For IRAS 01250+2832, multiple dusty star clusters are less likely to be present, because most of dusty star clusters are surrounded by diffuse emission of PAHs as well as optical emission lines.
For IRAS 01250+2832, multiple dusty star clusters are less likely to be present, because most of dusty star clusters are surrounded by diffuse emission of PAHs as well as optical emission lines.
IRAS 012504-2832 does not show any evidence for such activities.
IRAS 01250+2832 does not show any evidence for such activities.
If the assumption that the dust emission is optically thick is not correct, the NIR red continua of LEDA 84274 can be explained by a dust temperature of ~480— 580K and emissivity of A? with 6=1—2.
If the assumption that the dust emission is optically thick is not correct, the NIR red continua of LEDA 84274 can be explained by a dust temperature of $\sim 480-580$ K and emissivity of $\lambda^{-\beta}$ with $\beta=1-2$.
In spiral galaxies that normally display the thermal emission of optically thin dust, detected non-stellar NIR excess continuum with a temperature of ~ 1000K(8= 2) and concluded that the NIR excess continuum originates in the interstellar matter of the galaxies based on the linear correlation between emission from aromatic carbon and the excess.
In spiral galaxies that normally display the thermal emission of optically thin dust, detected non-stellar NIR excess continuum with a temperature of $\sim 1000$ $\beta=2$ ) and concluded that the NIR excess continuum originates in the interstellar matter of the galaxies based on the linear correlation between emission from aromatic carbon and the excess.
In this case, the NIR excess has a luminosity of only a few percent of the FIR luminosity.
In this case, the NIR excess has a luminosity of only a few percent of the FIR luminosity.
However, the ratio of the NIR luminosity to the FIR luminosity is 0.15 and 0.75 for LEDA 84274 and IRAS 01250+2832, respectively.
However, the ratio of the NIR luminosity to the FIR luminosity is 0.15 and 0.75 for LEDA 84274 and IRAS 01250+2832, respectively.
Thus, the emission from the hot dust in LEDA 84274 and IRAS 0125042832 seems to be different from those of spiral galaxies.
Thus, the emission from the hot dust in LEDA 84274 and IRAS 01250+2832 seems to be different from those of spiral galaxies.
We suspect that the hot dust associated with obscured AGNs.
We suspect that the hot dust associated with obscured AGNs.
But the possibility of multiple dusty star clusters in both objects, especially LEDA 84274, cannot be ruled out.
But the possibility of multiple dusty star clusters in both objects, especially LEDA 84274, cannot be ruled out.
Observations at other wavelengths can resolve the question.
Observations at other wavelengths can resolve the question.
In particular, X-ray observations can provide important information on the energy source, but both galaxies are not detected in the ROSAT all-sky survey faint source catalogue
In particular, X-ray observations can provide important information on the energy source, but both galaxies are not detected in the ROSAT all-sky survey faint source catalogue
In this paper, we have (re-)determined the reddening values for 443 clusters and cluster candidates in M31, as well as metallicities for 209 sample objects without spectroscopic observations.
In this paper, we have (re-)determined the reddening values for 443 clusters and cluster candidates in M31, as well as metallicities for 209 sample objects without spectroscopic observations.
We have followed the methods described by Barmbyetal.(2000), who found that the M31 and Galactic extinction laws are the same within the observational errors, and that the M31 and Galactic GC C-M relations are also consistent with each other.
We have followed the methods described by \citet{bh00}, who found that the M31 and Galactic extinction laws are the same within the observational errors, and that the M31 and Galactic GC C-M relations are also consistent with each other.
The sample of spectroscopic and photometric data used in this paper is the newest and largest to date.
The sample of spectroscopic and photometric data used in this paper is the newest and largest to date.
The spectroscopic data were obtained from the most recent references currently available and the photometric data are from the most comprehensive catalogue of M31 clusters available at present, which includes 337 confirmed GCs and 688 GC candidates.
The spectroscopic data were obtained from the most recent references currently available and the photometric data are from the most comprehensive catalogue of M31 clusters available at present, which includes 337 confirmed GCs and 688 GC candidates.
Using the metallicities of the largest sample of clusters and cluster candidates at hand, we studied the properties of the M31 clusters.
Using the metallicities of the largest sample of clusters and cluster candidates at hand, we studied the properties of the M31 clusters.
Our main conclusions are summarised below: We reiterate that in using the method of Barmbyetal.(2000)., there are two major unavoidable assumptions (acknowledged by these authors), i.e. that in the Milky Way and in M31 both the extinction law and the intrinsic colours of the GCs are the same.
Our main conclusions are summarised below: We reiterate that in using the method of \citet{bh00}, there are two major unavoidable assumptions (acknowledged by these authors), i.e. that in the Milky Way and in M31 both the extinction law and the intrinsic colours of the GCs are the same.
The latter assumption seems reasonable, since there is no evidence that GCs in different galaxies have different intrinsic colours.
The latter assumption seems reasonable, since there is no evidence that GCs in different galaxies have different intrinsic colours.
Regarding the former assumption, there is inconsistent evidence as to whether or not this is a valid assumption.
Regarding the former assumption, there is inconsistent evidence as to whether or not this is a valid assumption.
For example, WalterbosKennicutt(1988) found that the extinction law in M31 is very similar to that in the Milky Way, by analysing the two major dust lanes on the near side of M31; however, several studies have suggested that the reddening in M31 appears to be peculiar: with E(U—B)/E(BV)1.01X0.11 dye&Richter1985) and E(U—B)/E(BV)~0.5 (Masseyetal.1995),, compared to 0.72 for the same ratio in the Milky Way.
For example, \citet{wk88} found that the extinction law in M31 is very similar to that in the Milky Way, by analysing the two major dust lanes on the near side of M31; however, several studies have suggested that the reddening in M31 appears to be peculiar: with $E(U-B)/E(B-V)=1.01\pm0.11$ \citep{ir85} and $E(U-B)/E(B-V)\sim 0.5$ \citep{Massey95}, compared to 0.72 for the same ratio in the Milky Way.
Based on a large sample of GCs with optical and near-infrared photometric data, Barmbyetal.(2000) demonstrated that the U- and K-band extinction curve of M31 is consistent with that of the Milky Way, with total-to-selective extinction coefficient Ry=3.1.
Based on a large sample of GCs with optical and near-infrared photometric data, \citet{bh00} demonstrated that the $U$ - and $K$ -band extinction curve of M31 is consistent with that of the Milky Way, with total-to-selective extinction coefficient $R_V=3.1$.
In fact, the former assumption is plausible because in the M31 disc the composition and size distribution of the large normal grains which dominate the dust mass may be similar to those in the Milky Way (see&Helou1996)..
In fact, the former assumption is plausible because in the M31 disc the composition and size distribution of the large normal grains which dominate the dust mass may be similar to those in the Milky Way \citep[see for details,][]{xh96}.
As an example, we will discuss in some detail the reddening value of the M31 GC B037 (a.k.a.
As an example, we will discuss in some detail the reddening value of the M31 GC B037 (a.k.a.
037-B327), which is known to be an extremely red object.
037-B327), which is known to be an extremely red object.
There are a few references that discuss this GC, including Barmbyetal.(2002b),, Maetal.(2006a),, Maetal.(2006c) and Cohen (2006)..
There are a few references that discuss this GC, including \citet{bk02}, \citet{Ma06}, \citet{ma06} and \citet{cohen06}. .
Kron&Mayall(1960) first noticed an extremely red colour in photographic (P) and visual (V) bands for B037, and determined its absorption to be Αν=3.90 mag.
\citet{km60} first noticed an extremely red colour in photographic $P$ ) and visual $V$ ) bands for B037, and determined its absorption to be $A_V=3.90$ mag.
Based on the photometric data for M31 star clusters in U,B, and V of VeteSnik(1962a),, VeteSnik(1962b) studied the reddening values for these objects and found that Β037 was the most highly reddened in his sample, with E(B—V)=1.28 mag (Av=4.10 mag).
Based on the photometric data for M31 star clusters in $U, B$ , and $V$ of \citet{ve62}, \citet{ve62b} studied the reddening values for these objects and found that B037 was the most highly reddened in his sample, with $E(B-V)=1.28$ mag $A_V=4.10$ mag).
Cramptonetal.(1985) calibrated (B—V)o as a function of spectroscopic slope parameter S of the continuum between ~4000 and 5000Á,, and then determined the intrinsic colours for about 40 GCs and GC candidates, including B037.
\citet{Crampton85} calibrated $(B-V)_{\rm 0}$ as a function of spectroscopic slope parameter $S$ of the continuum between $\sim 4000$ and 5000, and then determined the intrinsic colours for about 40 GCs and GC candidates, including B037.
Cramptonetal.(1985) presented a reddening value for B037 of E(B— mag.
\citet{Crampton85} presented a reddening value for B037 of $E(B-V)=1.48$ mag.
Armed with a large database of multicolour photometry, Barmbyetal.(2000) determined the reddening value for each individual M31 GC, including B037, using the correlations between optical and infrared colours and metallicity based on various “reddening-free” parameters, and derived E(B—V)=1.38+0.02 mag for B037.
Armed with a large database of multicolour photometry, \citet{bh00} determined the reddening value for each individual M31 GC, including B037, using the correlations between optical and infrared colours and metallicity based on various “reddening-free” parameters, and derived $E(B-V)=1.38\pm0.02$ mag for B037.
Using spectroscopic metallicities to predict the intrinsic colours, Barmbyetal.(2002b) rederived the reddening value for this GC, E(B—V)=1.30+0.04 mag.
Using spectroscopic metallicities to predict the intrinsic colours, \citet{bk02} rederived the reddening value for this GC, $E(B-V)=1.30\pm0.04$ mag.
Recently, Maetal.(2006a) determined the reddening and age of the B037 by comparing multicolour photometry with theoretical stellar population synthesis models.
Recently, \citet{Ma06} determined the reddening and age of the B037 by comparing multicolour photometry with theoretical stellar population synthesis models.
The reddening towards B037 determined by Maetal.(2006a) is E(B—V)=1.360+0.013 mag.
The reddening towards B037 determined by \citet{Ma06} is $E(B-V)=1.360\pm0.013$ mag.
The reddening value for B037 determined in this paper is E(B—V)=1.21+0.03 mag.
The reddening value for B037 determined in this paper is $E(B-V)=1.21\pm0.03$ mag.
It is clear that the consistent reddening values for B037 from different references confirm that this cluster suffers from very large extinction.
It is clear that the consistent reddening values for B037 from different references confirm that this cluster suffers from very large extinction.
In fact, Ma(2006c) showed the dust lane across theface of the cluster using anHST/ACS image, which may partiallyaccount for its very large reddening value (seealsoCohen 2006)..
In fact, \citet{ma06} showed the dust lane across theface of the cluster using an/ACS image, which may partiallyaccount for its very large reddening value \citep[see also][]{cohen06}. .
We would like to thank the referee, Terry Bridges, for providing rapid and thoughtful report that helped improve the original manuscript greatly.
We would like to thank the referee, Terry Bridges, for providing rapid and thoughtful report that helped improve the original manuscript greatly.
This work has been supported by the Chinese
This work has been supported by the Chinese
component.
component.
This is significantly larger than the radius of the photo-ionising disk found by Moraes and Diaz (2009)).
This is significantly larger than the radius of the photo-ionising disk found by Moraes and Diaz \cite{Mora09}) ).
The reason for this difference is not clear at present; parts of the disk at intermediate radii may not produce much ionising radiation and/or may only contribute to the profile wings.
The reason for this difference is not clear at present; parts of the disk at intermediate radii may not produce much ionising radiation and/or may only contribute to the profile wings.
It must be noted that the radius of the emitting region appears somewhat too large when compared with the equatorial radius for an accretion disk around the white dwarf according to the two possible system parameters of Kürrster and Barwig (1988)), but the distance from the centre of the white dwarf to the inner Lagrangian point is close to the calculated value of the outer radius of the disk, using expressions on page 33 of Warner (1995)).
It must be noted that the radius of the emitting region appears somewhat too large when compared with the equatorial radius for an accretion disk around the white dwarf according to the two possible system parameters of Kürrster and Barwig \cite{Kurs88}) ), but the distance from the centre of the white dwarf to the inner Lagrangian point is close to the calculated value of the outer radius of the disk, using expressions on page 33 of Warner \cite{Warn95}) ).
The simplest explanation of these results is that some radiation is emitted by regions of higher radial velocity at the outer edge of the disk than that directly deduced from the observations, but that it is occulted because of a particular geometry of the line emitting region.
The simplest explanation of these results is that some radiation is emitted by regions of higher radial velocity at the outer edge of the disk than that directly deduced from the observations, but that it is occulted because of a particular geometry of the line emitting region.
High velocity components have been observed in in the past, in particular in the high velocity resonance line P Cygni absorption component in the ultraviolet, which had an edge radial velocity of -5000 km s! (Selvelli and Friedjung, 2003)).
High velocity components have been observed in in the past, in particular in the high velocity resonance line P Cygni absorption component in the ultraviolet, which had an edge radial velocity of -5000 km $^{-1}$ (Selvelli and Friedjung, \cite{Selv03}) ).
Such a component could have been produced in a high velocity jet.
Such a component could have been produced in a high velocity jet.
No evidence for it is seen today in our Ha profiles, nor in older optical profiles such as that of the 1990 spectrum (Ringwald et al. 1996)).
No evidence for it is seen today in our $\alpha$ profiles, nor in older optical profiles such as that of the 1990 spectrum (Ringwald et al. \cite{Ring96}) ).
One can thus wonder to which extent the ultraviolet spectrum of HR Del has changed between the lifetime of IUE and the epoch of our optical spectra, but no UV spectra could be obtained since the end of IUE in 1996.
One can thus wonder to which extent the ultraviolet spectrum of HR Del has changed between the lifetime of IUE and the epoch of our optical spectra, but no UV spectra could be obtained since the end of IUE in 1996.
The only recent, UV data available are two photometric points from the Galex all-sky survey, which give the following AB magnitudes (observations of September 2006): 12.78 in the Near-UV (central wavelength 2315 À) and 12.96 in the UV (A 1539 A).
The only recent, UV data available are two photometric points from the Galex all-sky survey, which give the following AB magnitudes (observations of September 2006): 12.78 in the Near-UV (central wavelength 2315 $\AA$ ) and 12.96 in the Far-UV $\lambda$ 1539 $\AA$ ).
If taken at face value, these magnitudes would indicate a small fading in the near-UV, and a stronger one (about a factor of two) in the Far-UV with respect to the average IUE spectrum from Selvelli and Friedjung (2003)).
If taken at face value, these magnitudes would indicate a small fading in the near-UV, and a stronger one (about a factor of two) in the Far-UV with respect to the average IUE spectrum from Selvelli and Friedjung \cite{Selv03}) ).
As however is slightly brighter than the saturation limit of the Galex detectors, those magnitudes have to be corrected for dead-time effects.
As however is slightly brighter than the saturation limit of the Galex detectors, those magnitudes have to be corrected for dead-time effects.
When this is done according to the precepts described in Morrissey et al. (2007)),
When this is done according to the precepts described in Morrissey et al. \cite{Morr07}) ),
the corrected AB magnitudes become 12.36 and 12.38 in the NUV and FUV respectively.
the corrected AB magnitudes become 12.36 and 12.38 in the NUV and FUV respectively.
This means that the NUV magnitude is still about the same as in 1988, within the uncertainties of the correction, and that in the FUV the fading has not exceeded 30 96 since 1988, which is quite remarkable.
This means that the NUV magnitude is still about the same as in 1988, within the uncertainties of the correction, and that in the FUV the fading has not exceeded 30 $\%$ since 1988, which is quite remarkable.
This is illustrated in Fig.14,, which is Fig.
This is illustrated in \ref{UV}, which is Fig.