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B. f. where both the singlet aud pair magnitudes have con. plotted versus ejection angle. it is clear that Hs ts not thecasc. | 4, where both the singlet and pair magnitudes have been plotted versus ejection angle, it is clear that this is not the. |
. Thus it is not possible to explain the changes in apparent magnitude using his rotation aremment. | Thus it is not possible to explain the changes in apparent magnitude using this rotation argument. |
Since the ejection angle is also related to time (and the age of each triplet). the above relations can also be explained by time variations in the redshifts aud maguitudes. | Since the ejection angle is also related to time (and the age of each triplet), the above relations can also be explained by time variations in the redshifts and magnitudes. |
Such a model was suggested by NarlikarandDas(1980) who proposed that QSOs might be born out of eveuts in ealactic nuclei as matter was born out of the Big Bang. | Such a model was suggested by \citet{nar80b} who proposed that QSOs might be born out of events in galactic nucleii as matter was born out of the Big Bang. |
Thev have variable mass. increasing from zero at birth. and a high redshift that decreases with age (seeArp1905,section9).. | They have variable mass, increasing from zero at birth, and a high redshift that decreases with age \citep[see][section 9]{arp98}. |
Fie. | Fig. |
3 has been re-plotted in Fig. | 3 has been re-plotted in Fig. |
5(a) aud (1). with age plotted on a logavitlinic scale. | 5(a) and (b), with age plotted on a logarithmic scale. |
Although the curves cau only be estimated for ages in excess of Lx10° ves. they imply that by. ~105 yrs the apparent magnitudes will have approached those expected for normal galaxies at the distance of NGC 1068. and the iutrinsic redshift component will have largely disappeared. | Although the curves can only be estimated for ages in excess of $\times10^{6}$ yrs, they imply that by $\sim10^{8}$ yrs the apparent magnitudes will have approached those expected for normal galaxies at the distance of NGC 1068, and the intrinsic redshift component will have largely disappeared. |
This indicates that the intrinsic redshift component isa trausicut one. and therefore the QSOs themselves. are short-lived (107—105 325) compared to the age of the Universe (~1.ς1010 333).On the other hand. the rate of change of the redshift would only be At~201089 xy. J| which would be impossible to detect. | This indicates that the intrinsic redshift component is a transient one, and therefore the QSOs themselves, are short-lived $10^{7} - 10^{8}$ yrs) compared to the age of the Universe $\sim1.4\times10^{10}$ yrs).On the other hand, the rate of change of the redshift would only be $\Delta z \sim2\times10^{-6}$ $^{-1}$ , which would be impossible to detect. |
Compact objects reported close to nearby active galaxies (Arp1900:Burbidge1999:Clin have measured redshifts z = 2. This | Compact objects reported close to nearby active galaxies \citep{arp99b,bur99,chu98} have measured redshifts z $\lesssim2$ This |
ds the cluster huninosity function. aud M measures the cluster richuess. | is the cluster luminosity function, and $\Lambda_{cl}$ measures the cluster richness. |
The parameters a aud r2 are the apparent magnitude corresponding to the characteristic Iumuinositv of the cluster galaxies. aud the projected value of the cluster characteristic scale leugth. | The parameters $m^*$ and $r_c$ are the apparent magnitude corresponding to the characteristic luminosity of the cluster galaxies, and the projected value of the cluster characteristic scale length. |
Frou this model one can write an approximate likelihood £ of having a cluster at a given position as The matched filter algorithin is obtained using a series of à functions to represeut the discrete distribution of galaxies in a given catalog. instead of the continuous functiou Dir.in). | From this model one can write an approximate likelihood ${\cal L}$ of having a cluster at a given position as The matched filter algorithm is obtained using a series of $\delta$ functions to represent the discrete distribution of galaxies in a given catalog, instead of the continuous function $D(r,m)$. |
The application of the filter to an input galaxy catalog is therefore accomplished by evaluating the sti where Pry) is the angular weighting fuuctiou (radial filter). and ο). is the huninosity weighting faction (flux filter). at every poiut (7./) iu the survey. aud over a range of redshifts (which corresponds to a range of à. aud ny values]. | The application of the filter to an input galaxy catalog is therefore accomplished by evaluating the sum where $P(r_k)$ is the angular weighting function (radial filter), and $L(m_k)$ is the luminosity weighting function (flux filter), at every point $(i,j)$ in the survey, and over a range of redshifts (which corresponds to a range of $r_c$ and $m^*$ values). |
Iun practice. since the optima flux filter Lip)=OCH}δα) has a divergent integral at thle faint maeuitude limit when ois a Schechter function (Schechter 1976). 1f Is necessary o niodifv this filter. | In practice, since the optimal flux filter $L(m_k) = \phi(m_k-m^*) / b(m_k)$ has a divergent integral at the faint magnitude limit when $\phi$ is a Schechter function (Schechter 1976), it is necessary to modify this filter. |
The solution propose by P96 is to introduce a oower-LI;w. cutoff of the form 1Snn that. with i=(ταν would correspond to an extra weighting by he dux ofthe galaxy. | The solution proposed by P96 is to introduce a power-law cutoff of the form $10^{-\beta(m-m^*)}$ that, with $\beta=0.4$, would correspond to an extra weighting by the flux of the galaxy. |
The optimal radial filter is given w the assuned cluster projected radial profile. | The optimal radial filter is given by the assumed cluster projected radial profile. |
Here a modified IHubble profile is used. truucated at au arbitrary radius which is large compared to the cluster core radius. | Here a modified Hubble profile is used, truncated at an arbitrary radius which is large compared to the cluster core radius. |
Therefore the fiux aud radial filter have the form aud where o(m—nes) is taken to be a Schechter function. r, is the value of the projected cluster core radius. andr... is the arbitrary cutoff radius; | Therefore the flux and radial filter have the form and where $\phi(m-m*)$ is taken to be a Schechter function, $r_c$ is the value of the projected cluster core radius, and $r_{co}$ is the arbitrary cutoff radius. |
One further correction to the algorithia is required. | One further correction to the algorithm is required. |
The normalization adopted for the fiux filter (equation 21 in P96) is iu fact only strictly correct for a pure backeround distribution. but mtroduces an error in the redshift estimate of cluster candidates when an overdensity of galaxies is preseut. | The normalization adopted for the flux filter (equation 21 in P96) is in fact only strictly correct for a pure background distribution, but introduces an error in the redshift estimate of cluster candidates when an overdensity of galaxies is present. |
To compensate for this effect. aud obtain a corrected filter 54,0.j). the "Olde procedure proposet by P96 (their equations 22 - 26) was adopted here. | To compensate for this effect, and obtain a corrected filter $S_{corr}(i,j)$, the same procedure proposed by P96 (their equations 22 - 26) was adopted here. |
The matched filter aleorithiu described above is at the core of the EIS cluster searching pipcline that was iiplemeuted to process the galaxy catalogs produced by the EIS data reduction pipeline. | The matched filter algorithm described above is at the core of the EIS cluster searching pipeline that was implemented to process the galaxy catalogs produced by the EIS data reduction pipeline. |
In this section the details about its iuplemieutation. and the methods adopted to ideutifv sjenificaut cluster candidates are described. | In this section the details about its implementation, and the methods adopted to identify significant cluster candidates are described. |
By evaluating the sin 5,0.7) for cach clement of a two-dimensional array (7./) a filtered image (hereafter he “Likelihood map”. see Fie. | By evaluating the sum $S_{corr}(i,j)$ for each element of a two-dimensional array $(i,j)$ a filtered image (hereafter the “Likelihood map”, see Fig. |
3 for an example) of the ealaxy cataog is created. | \ref{fig:likelihood} for an example) of the galaxy catalog is created. |
The elements (7.7) correspou O0 a series of equally spaced points that cover the entire νο AYOS. | The elements $(i,j)$ correspond to a series of equally spaced points that cover the entire survey area. |
At cach point C.7) he suu is evaluated a junber of iues. with the radial aud flux filters tune o differeut cluster redshift values (this will. hereafter called. the "filter. redshift”). | At each point $(i,j)$ the sum is evaluated a number of times, with the radial and flux filters tuned to different cluster redshift values (this will hereafter be called the “filter redshift”). |
The ninm adopte filter redshift is :,,;,=0.2. while the maxima redshift Lou Us determined by finding the redshift value at which he appareut characteristic magnitude mz) becomes coniparable to the lamiting magnitude of the catalog. | The minimum adopted filter redshift is $z_{min} = 0.2$, while the maximum redshift $z_{max}$ is determined by finding the redshift value at which the apparent characteristic magnitude $m^*(z)$ becomes comparable to the limiting magnitude of the catalog. |
This approach ooOgives à i442=1.3 for the typical laniting uaenitude of £=23. | This approach gives a $z_{max} =
1.3$ for the typical limiting magnitude of $I = 23$ . |
The characteristic luuünositv. AL aud the cluster core radius are assumed to remain fixec in phivsical units. aud also the luninositv function fainut-end slope. a. is fixed. | The characteristic luminosity $M^*$ and the cluster core radius are assumed to remain fixed in physical units, and also the luminosity function faint-end slope, $\alpha$, is fixed. |
The observable quantities i” arc rare assumed to vary with redshift as in an Wy=75 /λῃο Ου] standard cosnologyv. | The observable quantities $m^*$ and $r_c$ are assumed to vary with redshift as in an $_0$ =75, $\Omega_0$ =1 standard cosmology. |
The adopte cluster pavanetors. taken from P96. are &.=1hUspe. Pog=1h !Mpe. and My-22.33. | The adopted cluster parameters, taken from P96, are $r_c = 1h^{-1}$ kpc, $r_{co} =
1h^{-1}$ Mpc and $M^*_I = -22.33$. |
The valuc of AM, was corrected to the Cousins svstem adopting the transformation given in P96. | The value of $M^*_I$ was corrected to the Cousins system adopting the transformation given in P96. |
The conversion frou the characteristic Dunünositv to the observable apparent maeguitude 11) requires a asstuuption to be made on the dF-correction of the ealaxies. | The conversion from the characteristic luminosity to the observable apparent magnitude $m^*$ requires an assumption to be made on the K-correction of the galaxies. |
Both a non-evolving galaxy model. aud a noclel with passive evolution of the stellar population have been considered. | Both a non-evolving galaxy model, and a model with passive evolution of the stellar population have been considered. |
The former is based on a template spectitun of an elliptical galaxy. taken frou, Coleman (1980). while for the latter svuthetic spectra. obtained with Druzual aud Charlot stellar population svuthesis code (Bruzual Charlot 1993). for a galaxy with solar metallicitv. a star formation listory with a single instantaneous burst of star formation. and a present age of 12 Cir. were used. | The former is based on a template spectrum of an elliptical galaxy, taken from Coleman (1980), while for the latter synthetic spectra, obtained with Bruzual and Charlot stellar population synthesis code (Bruzual Charlot 1993), for a galaxy with solar metallicity, a star formation history with a single instantaneous burst of star formation, and a present age of 12 Gyr, were used. |
It is iuportaut to empliasize that the choice of a I&-correction model does not significantly inpact the cluster detections. | It is important to emphasize that the choice of a K-correction model does not significantly impact the cluster detections. |
The pixel size of the Likelibood maps (1.6. the spacing between adjacent (7./) array elements) is taken to be 26.3 aresec. corresponding to the value of the projected cluster core radius. for a cluster at a redshift of 0.6. | The pixel size of the Likelihood maps (i.e. the spacing between adjacent $(i,j)$ array elements) is taken to be 26.3 arcsec, corresponding to the value of the projected cluster core radius, for a cluster at a redshift of 0.6. |
Ideally. one would like to have a varyiug pixel size. corresponding to a fixed fraction of a cluster projected core radius at all filter redshifts. | Ideally, one would like to have a varying pixel size, corresponding to a fixed fraction of a cluster projected core radius at all filter redshifts. |
Towever this would complicate the comparison between Likelihoodmaps obtained with different filter redshift. aud since this comparison is extremely useful | However this would complicate the comparison between Likelihoodmaps obtained with different filter redshift, and since this comparison is extremely useful |
perform two analyses, one on the original particle distribution and one on a shifted half a domain grid cell in each direction, taking (byperiodicity into account) particle distribution; we otherwise keep all other parameters constant at DomGrid— 64, DomRef—5.0 and RefRef—5.0. | perform two analyses, one on the original particle distribution and one on a shifted (by half a domain grid cell in each direction, taking periodicity into account) particle distribution; we otherwise keep all other parameters constant at $\DomGrid = 64$ , $\DomRef = 5.0$ and $\RefRef = 5.0$. |
The ratio of the resulting mass-functions and spin-parameter distributions are shown in figure 3.. | The ratio of the resulting mass-functions and spin-parameter distributions are shown in figure \ref{fig:shift_global}. |
Small deviations can be seen in the mass function, which is however of the order of 1—3 per cent in a few bins and then mostly due to one or two haloes changing bin. | Small deviations can be seen in the mass function, which is however of the order of $1-3$ per cent in a few bins and then mostly due to one or two haloes changing a bin. |
'This translate into some scatter in the spin parametera distribution, however the seemingly large deviation on the low- and high-spin end are artificially enhanced due to the low number counts in the bins, we also excluded halos with less than 100 particles (cf. refsubsec:refcrit)) from this plot. | This translate into some scatter in the spin parameter distribution, however the seemingly large deviation on the low- and high-spin end are artificially enhanced due to the low number counts in the bins, we also excluded halos with less than 100 particles \\ref{subsec:refcrit}) ) from this plot. |
The scatter here is also of the order of 1—3 per cent. | The scatter here is also of the order of $1-3$ per cent. |
It is important to keep in mind that the derived properties can vary to this extent simply due to numerical effects. | It is important to keep in mind that the derived properties can vary to this extent simply due to numerical effects. |
However, the main point of this comparison is to verify that, up to numerical effects caused by the perturbed density field, we do not introduce any systematic deviation in the statistical properties. | However, the main point of this comparison is to verify that, up to numerical effects caused by the perturbed density field, we do not introduce any systematic deviation in the statistical properties. |
We now focus on the choice ofDomGrid,, i.e. the size of the domain grid. | We now focus on the choice of, i.e. the size of the domain grid. |
To this extent, we employ one process and vary the domain grid and the refinement criterion oon the domain grid. | To this extent, we employ one process and vary the domain grid and the refinement criterion on the domain grid. |
In figure 4 we show the deduced mass-function and spin parameter distribution for the four cases. | In figure \ref{fig:serial_dom} we show the deduced mass-function and spin parameter distribution for the four cases. |
As can be seen, the impact of the domain grid choice is negligible; small deviations can be seen at the high mass end of the mass function in B50, however these are caused by (of order) one halo changing bin across runs with varied parameters. | As can be seen, the impact of the domain grid choice is negligible; small deviations can be seen at the high mass end of the mass function in B50, however these are caused by (of order) one halo changing bin across runs with varied parameters. |
We will discuss the drop of N(AM) at the low M end of the mass function below. | We will discuss the drop of $N(\Delta M)$ at the low $M$ end of the mass function below. |
In view of the parallel strategy the insensitivity of the results to the choice of the domain grid is reassuring, as we can use a rather coarse domain grid and start the refinement hierarchy from there; note that the domain grid will be allocated completely on each CPU and hence choosing a fine grid would lead to a large number of empty cells in the parallel version. | In view of the parallel strategy the insensitivity of the results to the choice of the domain grid is reassuring, as we can use a rather coarse domain grid and start the refinement hierarchy from there; note that the domain grid will be allocated completely on each CPU and hence choosing a fine grid would lead to a large number of empty cells in the parallel version. |
It can also be seen that the choice of the refinement criterion on the domain grid has no impact. | It can also be seen that the choice of the refinement criterion on the domain grid has no impact. |
This is because the domain grid is coarse enough as to justify refinement everywhere anyhow. | This is because the domain grid is coarse enough as to justify refinement everywhere anyhow. |
In fact, only very pronounced underdense regions would not trigger refinements for the choices of parameters and particle resolution. | In fact, only very pronounced underdense regions would not trigger refinements for the choices of parameters and particle resolution. |
Now we investigate the effect of choosing a different refinement criterion (RefRef)) for the refined grids. | Now we investigate the effect of choosing a different refinement criterion ) for the refined grids. |
We limit ourselves to a domain grid of DomGrid—128 cells per dimension and use a refinement criterion on the domain grid of DomRef—1.0 with one process. | We limit ourselves to a domain grid of $\DomGrid = 128$ cells per dimension and use a refinement criterion on the domain grid of $\DomRef = 1.0$ with one process. |
We then vary the refinement criterion on the refined grids from RefRef—5.0 (this corresponds to analysis already used above when investigating the impact of the domain grid and its refinement criterion) to RefRef—2.0 in steps of 1.0. | We then vary the refinement criterion on the refined grids from $\RefRef =
5.0$ (this corresponds to analysis already used above when investigating the impact of the domain grid and its refinement criterion) to $\RefRef
= 2.0$ in steps of $1.0$. |
In figures5 and 6 we show the effect of varying the refinement criterion on the refined grids on the | In figures\ref{fig:serial_mass_ref} and \ref{fig:serial_lambda_ref} we show the effect of varying the refinement criterion on the refined grids on the |
feature at 1.611 pam had a cot[used appearancee. | feature at 1.611 $\mu$ m had a confused appearance. |
More recently. the analysis of the FeH bauds in M and L cdwarfs has been exteuced by Cushiugeal.(2003).. who showed that the 1.625 jun feature is the strongest amoung the [οιr features promient du t1ο H-baud 'eelon. | More recently, the analysis of the FeH bands in M and L dwarfs has been extended by \citet{cus03}, who showed that the 1.625 $\mu$ m feature is the strongest among the four features prominent in the $H$ -band region. |
Thus all the four features (1.583. 1.591. 1.611. and 1.625 jan) in our speetra of L3 - L5 dwa‘fs can be identiied with FeH. A question is ifthe same iclentification cai e applied to t1ο L6.5 aud Ls dwaTs in our sample. | Thus all the four features (1.583, 1.591, 1.611, and 1.625 $\mu$ m) in our spectra of L3 - L5 dwarfs can be identified with FeH. A question is if the same identification can be applied to the L6.5 and L8 dwarfs in our sample. |
Another possibility is that the 1.63 as well as 1.67 pun feature itin the L6.5 axd L3 dwars ds due to methane bands. | Another possibility is that the 1.63 as well as 1.67 $\mu$ m feature in the L6.5 and L8 dwarfs is due to methane bands. |
We have aready attribute hese features as dle lo 1jethanele iL the L62 dwarf 2NLASSO92043:) (Nakajima.Tsuji&Yanagisawa2001).. hiitl e possible preseuce of he FeH bands was uot shown at that time. | We have already attributed these features as due to methane in the L6.5 dwarf 2MASS0920+35 \citep{nak01}, but the possible presence of the FeH bands was not known at that time. |
The 1.63 ja 1 feature cau also be seei in all the L dwaTs whi the 1.67 jan feature is not clear in te L3 - L5 dwarls and aypears firs in L6.5. | The 1.63 $\mu$ m feature can also be seen in all the L dwarfs while the 1.67 $\mu$ m feature is not clear in the L3 - L5 dwarfs and appears first in L6.5. |
Also. 1je. 1.58: and 1.591 jun ealures are rather we:uk in the L6.5 clwarls. | Also, the 1.583 and 1.591 $\mu$ m features are rather weak in the L6.5 dwarfs. |
These observations nay supσοι th appearance of nethane at L6.5. aud it is at least possible that FeH ane CH, both contribute to the 1.63 ane 1.67 jun features in L6.5 cdwarfs. | These observations may support the appearance of methane at L6.5, and it is at least possible that FeH and $_4$ both contribute to the 1.63 and 1.67 $\mu$ m features in L6.5 dwarfs. |
To settle this problem. higher resolution wo‘ks are ueeded. | To settle this problem, higher resolution works are needed. |
As for 2MLASS:|523-4-30. abseice of the 1.61 pan feature suMODooests that the 1.63 jun ealtire may be due o CH, rather th:ui FeH. The nethane features at 1.63 and 1.67 san are clearly seet iu the T2 dwar- alid strerethenu rapidly oward the later T. dwarls. | As for 2MASS1523+30, absence of the 1.61 $\mu$ m feature suggests that the 1.63 $\mu$ m feature may be due to $_4$ rather than FeH. The methane features at 1.63 and 1.67 $\mu$ m are clearly seen in the T2 dwarf and strengthen rapidly toward the later T dwarfs. |
Where mehaue bands appear i qitical isse iu the spectral Classification of brown dwarfs. since methane is a seusitive Indica eluperatre. | Where methane bands appear is a critical issue in the spectral classification of brown dwarfs, since methane is a sensitive indicator of temperature. |
Αλλοιeh the methare bh:uxds depeud ou gravity as well as ou temperatu'e. 'eased abuicdauce o- juethaue at tje hiehergravity is compensated lor ο some extent by 'eased Ho CLA whic1 also Increases wit Leravity (Tsuji.Nakajima&Yanagisawa|2001). | Although the methane bands depend on gravity as well as on temperature, the increased abundance of methane at the highergravity is compensated for to some extent by the increased $_2$ CIA which also increases with gravity \citep{tsu03}. |
. It is to be noted that the H-hud region is by no means a good continuum window. bu is containinatec| by molecular bands of uukuowu origin in addition to FeH especially in L dwa‘Ts. | It is to be noted that the $H$ -band region is by no means a good continuum window, but is contaminated by molecular bands of unknown origin in addition to FeH especially in L dwarfs. |
Also. we notice a large clilference in tle spectra of the same spectral type. iuuelv between the L cwarfs 2M.ASS1507—16 and SD55221t)4-00: The H band spectrum of 2ALASSIS07—16 is rather lat as are those of the other L dwarls whie that of SD$522192-00 shows a promiuent peaking cente'ecd at about 1.7 jan. The reason for this difference is unknown at. present. | Also, we notice a large difference in the spectra of the same spectral type, namely between the L dwarfs $-$ 16 and SDSS2249+00: The $H$ band spectrum of $-$ 16 is rather flat as are those of the other L dwarfs while that of SDSS2249+00 shows a prominent peaking centered at about 1.7 $\mu$ m. The reason for this difference is unknown at present. |
In Fig.3. we show the region between J aud H bands. | In Fig.3, we show the region between $J$ and $H$ bands. |
H is to be noted that water "πας alt Lit gan can be well measured in the five spectra observed by CISCO despite the water "πας in the Earth's atmosphere. | It is to be noted that water bands at 1.4 $\mu$ m can be well measured in the five spectra observed by CISCO despite the water bands in the Earth's atmosphere. |
Unfortunately. only the edge of the 1.1 jam bands. arising in hotter environments of these cool chwarfs than the Earth's atmosphere. can be ueasured in the three objects observed by IRC'S. | Unfortunately, only the edge of the 1.4 $\mu$ m bands, arising in hotter environments of these cool dwarfs than the Earth's atmosphere, can be measured in the three objects observed by IRCS. |
One interesting result is that the HeO 1.1 gam baxl does uot appear stronger in the LS dwarf 2MASSI523430 than in the L6.5 dwarls 2MASSI71L+2°2 and .. | One interesting result is that the $_2$ O 1.4 $\mu$ m band does not appear stronger in the L8 dwarf 2MASS1523+30 than in the L6.5 dwarfs 2MASS1711+22 and . |
deficit is. about) half the positive superhump excess (Section 1.2). | deficit is about half the positive superhump excess (Section 1.2). |
As the 3028-s period. is about 0.9 percent longer than the 3001-s period. the corresponding ratio in V1405. Aq! (0.79) is somewhat larger than this. | As the 3028-s period is about 0.9 percent longer than the 3001-s period, the corresponding ratio in V1405 Aql (0.79) is somewhat larger than this. |
We have checked the connection between the deficit. ancl excess of negative and positive superhumps in CVs and found a new relation. | We have checked the connection between the deficit and excess of negative and positive superhumps in CVs and found a new relation. |
Table 2 presents the data on the periods in svstens showing both types of superhumps. | Table 2 presents the data on the periods in systems showing both types of superhumps. |
In Fig. | In Fig. |
5 we show this ratio for these svstenis. and we see a clear trend as a function of orbital period. | 5 we show this ratio for these systems, and we see a clear trend as a function of orbital period. |
Our result for VI405 Aq! fits this trend very well. | Our result for V1405 Aql fits this trend very well. |
Our results strengthen. the observational link between positive and negative superhumps. ancl supports the idea that the two tvpes of superhumps have a similar physical origin. namelv a precessing accretion disc. | Our results strengthen the observational link between positive and negative superhumps, and supports the idea that the two types of superhumps have a similar physical origin, namely a precessing accretion disc. |
ltetter ct al. ( | Retter et al. ( |
2002b) speculated that every permanent superhump svstenm may have both kinds of superhumps. | 2002b) speculated that every permanent superhump system may have both kinds of superhumps. |
The data on V1405 Aql support this idea. | The data on V1405 Aql support this idea. |
Wood. Montgomery Simpson (2000) showed that a tilted accretion disc can explain the presence of negative superhumps. however. it is still unclear what physical force drives the dise to precess in the nodal plane (Murray Armitage 1998). | Wood, Montgomery Simpson (2000) showed that a tilted accretion disc can explain the presence of negative superhumps, however, it is still unclear what physical force drives the disc to precess in the nodal plane (Murray Armitage 1998). |
Wynn (personal communication) suggested hat a strong magnetic field of the primary white cwarf can cause the dise to precess in the nodal plane. | Wynn (personal communication) suggested that a strong magnetic field of the primary white dwarf can cause the disc to precess in the nodal plane. |
Lf this idea is confirmed. negative superhumps are expected. in every permanent superhumyp LMXIUD. as their. primaries. have magnetic fields tvpically much stronger than in CVs. | If this idea is confirmed, negative superhumps are expected in every permanent superhump LMXRB, as their primaries have magnetic fields typically much stronger than in CVs. |
Montgomery (2001) and Montgomery et. al. ( | Montgomery (2001) and Montgomery et al. ( |
in oeparation) developed: analytic expressions for. the two vpes of superhumps. | in preparation) developed analytic expressions for the two types of superhumps. |
The theoretical. values for the ratio Ó-cfc mentioned above. were found to be consistent with he observations ofCVs. | The theoretical values for the ratio $\phi$ $\epsilon_{-}/\epsilon_{+}$ mentioned above, were found to be consistent with the observations of CVs. |
Montgomery ct al | Montgomery et al. |
s results can also explain the relation shown in Fig. | 's results can also explain the relation shown in Fig. |
5 for CVs (see their Fig. | 5 for CVs (see their Fig. |
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