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Due to a weaker, second-order nature of this resonance, the eccentricity increase is rather small. | Due to a weaker, second-order nature of this resonance, the eccentricity increase is rather small. |
Incidentally in this particular system, this is enough for the ice giants to cross orbits and scatter off of each other, but not off of one of the gasgiants?. | Incidentally in this particular system, this is enough for the ice giants to cross orbits and scatter off of each other, but not off of one of the gas. |
. It appears that somewhat larger eccentricities are needed. | It appears that somewhat larger eccentricities are needed. |
Testing each initial condition with a large number of numerical simulations, as discussed above, is rather time-consuming. | Testing each initial condition with a large number of numerical simulations, as discussed above, is rather time-consuming. |
Consequently, it is worthwhile to quantify the amplitudes of eccentricity jumps due to various resonance crossings before-hand if possible. | Consequently, it is worthwhile to quantify the amplitudes of eccentricity jumps due to various resonance crossings before-hand if possible. |
For this set of initial conditions, under the assumption of adiabatic migration, the eccentricity jumps are deterministic and can be estimated analytically (Henrard 1982). | For this set of initial conditions, under the assumption of adiabatic migration, the eccentricity jumps are deterministic and can be estimated analytically (Henrard 1982). |
Following the treatment of (Peale 1986, see also Murray Dermott we consider the planar internal first-order j:(j—1) 1999),resonant Hamiltonian where is the mean longitude, 7=c is the longitude of perihelion, A=(mmo)/(m+mo)/G(mom)a T=A(1—V1-—e?) are their respective Poincaré conjugate momenta, and the prime designates the outer planet. | Following the treatment of (Peale 1986, see also Murray Dermott 1999), we consider the planar internal first-order $j:(j-1)$ resonant Hamiltonian where $\lambda$ is the mean longitude, $\gamma=\varpi$ is the longitude of perihelion, $\Lambda = (m \ m_{\odot})/(m + m_{\odot})\sqrt{G(m_{\odot}+m)a}$ $\Gamma = \Lambda (1-\sqrt{1-e^2})$ are their respective Poincaré conjugate momenta, and the prime designates the outer planet. |
The secular changes in mean longitude and longitude of perihelion are accounted for by the last four terms, while f(a/a’) arises from the classical expansion of the planetary disturbing potential and is a function of Laplace coefficients and their derivatives. | The secular changes in mean longitude and longitude of perihelion are accounted for by the last four terms, while $f(a/a')$ arises from the classical expansion of the planetary disturbing potential and is a function of Laplace coefficients and their derivatives. |
γάμο on observations of the local universe. | based on observations of the local universe. |
Differential counts are very sensitive to the exact shape of the PALL eatures. which are crucely modelled here. | Differential counts are very sensitive to the exact shape of the PAH features, which are crudely modelled here. |
One could. also imagine that the cliscrepancy is partly due to the grain size distribution/chemical composition evolution with redshift. as the redshift) distribution of ISOCAM ealaxies has a meclian z 0.7. | One could also imagine that the discrepancy is partly due to the grain size distribution/chemical composition evolution with redshift, as the redshift distribution of ISOCAM galaxies has a median $z \approx$ 0.7. |
We plan to investigate these issues in more detail but that is bevond the scope of this paper. | We plan to investigate these issues in more detail but that is beyond the scope of this paper. |
Finally. »eceause of the way interactions are mocoelled. cach carly vpe galaxy undergoes a starburst after its host halo has just. collapsed. and it is not obvious that ISOCAM sources (sce figure 3)) in which the vast majority are LItCis (not ULIRGs). are properly described by such a violent. process. | Finally, because of the way interactions are modelled, each early type galaxy undergoes a starburst after its host halo has just collapsed, and it is not obvious that ISOCAM sources (see figure \ref{figinf}) ) in which the vast majority are LIRGs (not ULIRGs), are properly described by such a violent process. |
Dvnamical interactions (which are not mocelec in. detail rere). trigecring multiple milder starbursts. with time delays oetween them. might. provide a more realistic description of hese sources and this is another issue we plan to investigate in the future. | Dynamical interactions (which are not modeled in detail here), triggering multiple milder starbursts, with time delays between them, might provide a more realistic description of these sources and this is another issue we plan to investigate in the future. |
Another constraint on our models comes from the recdshift distributions. | Another constraint on our models comes from the redshift distributions. |
Their shapes seem to quite nicelv match the observations in the L band with a mean redshift’ of the distribution of ~ 0.6. (see top left. panel of figure 5)). | Their shapes seem to quite nicely match the observations in the I band with a mean redshift of the distribution of $\sim$ 0.6, (see top left panel of figure \ref{figred}) ). |
In the farHt (GO microns). the agreement with cata eathered in the north ecliptie pole region (NEPR) is also fairly convincing. | In the far–IR (60 microns), the agreement with data gathered in the north ecliptic pole region (NEPR) is also fairly convincing. |
As predicted in Silk Devriendt. (2000). inclusion of the cosmological constant. A. has shifted the nearLR ancl GO micron peaks towards higher recshifts aud produced a highrecshift tail in the L band. bringing the models into closer agreement with the data. | As predicted in Silk Devriendt (2000), inclusion of the cosmological constant, $\Lambda$, has shifted the near–IR and 60 micron peaks towards higher redshifts and produced a high–redshift tail in the I band, bringing the models into closer agreement with the data. |
Disks and early-types are found in comparable proportion in the I-band. with a slight domination of disks for z>0.3 and up to z—1.5. | Disks and early-types are found in comparable proportion in the I-band, with a slight domination of disks for $z>0.3$ and up to $z=1.5$. |
There are at least a couple of reasons for this behaviour. | There are at least a couple of reasons for this behaviour. |
First. in this redshift range. spheroids are already old. their star formation rates are very low and therefore their I-band luminosity comes from an old and dim stellar population. | First, in this redshift range, spheroids are already old, their star formation rates are very low and therefore their I-band luminosity comes from an old and dim stellar population. |
secondly. spheroids in a massive starburst phase at. these redshift’ experience high dust. absorption. which reduces | Secondly, spheroids in a massive starburst phase at these redshift experience high dust absorption, which reduces |
independent. but recent results. (Bullockctal.2001.. Wechlseretal. 2002)) show a correlation between these parameters. | independent, but recent results \citealt{Bul:01}, , \citealt{We:02}) ) show a correlation between these parameters. |
The NEW clensity distribution is then a one-parameter familv. namely the virial mass Ale; | The NFW density distribution is then a one-parameter family, namely the virial mass $M_{vir}$. |
From Wechlseret.al.(2002) we take the relations linking AJ,.;, to the concentration parameter e(=res£r). rs and py. at redshift z=0 and for a Universe with A=0.7 and £3,=0.3. starting with Al,;,=ο Ovhere Aj, is the virial overdensity ancl its valueana, is about 337 at 2=0. py is the critical density of the Universe and rà; is the virial radius): NEW halo has then a central density cusp. with prowox [or r+0. and a prolile/amplitude which is controlled v a free parameter AM, | From \citet{We:02} we take the relations linking $M_{vir}$ to the concentration parameter $c~(=r_{vir}/r_s)$, $r_s$ and $\rho_s$, at redshift $z=0$ and for a Universe with $\Lambda = 0.7$ and $\Omega_0 = 0.3$, starting with $M_{vir} \equiv \frac {4}{3} \pi \Delta_{vir} \rho_{c} r_{vir}^3$ (where $\Delta_{vir}$ is the virial overdensity and its value is about 337 at $z=0$, $\rho_{c}$ is the critical density of the Universe and $r_{vir}$ is the virial radius): NFW halo has then a central density cusp, with $\rho_{\rm NFW}
\propto r^{-1}$ for $r \rightarrow 0$, and a profile/amplitude which is controlled by a free parameter $M_{vir}$. |
Notice that in principle. aciabatic contraction of the »imordial dark matter halo due to barvon infall should be aken into account. but since the ellect is to render the malo evenmore concentrated. aggravating thus the known xoblems of the NEW haloes. we neglect. it. | Notice that in principle, adiabatic contraction of the primordial dark matter halo due to baryon infall should be taken into account, but since the effect is to render the halo evenmore concentrated, aggravating thus the known problems of the NFW haloes, we neglect it. |
We constrain the virial halo mass to be Mz,δ1 AL. in that. for a low luminosity spiral. it must. presumably » substantially lower than that of the Alilky Was and other very luminous galaxies. for which it is safely estimated Als&21095 M. (Chengalur.Salpeter& Terzian 1993.. Wilkinson&Evans 1999)). | We constrain the virial halo mass to be $M_{vir}<8 \times 10^{11}$ $_{\odot}$, in that, for a low luminosity spiral, it must presumably be substantially lower than that of the Milky Way and other very luminous galaxies, for which it is safely estimated: $M_{vir} \simeq 2 \times 10^{12}$ $_{\odot}$ \citealt{C:93}, \citealt{W:99}) ). |
This constraint allects only N7339 and ESO 79-C1H. due to the relatively [limited extension of their rotation curves. which prevents to rule out large AC'DAL haloes. | This constraint affects only N7339 and ESO 79-G14, due to the relatively limited extension of their rotation curves, which prevents to rule out large $\Lambda$ CDM haloes. |
ltecent numericalsimulations by Mooreetal.(1998). vielded a more concentrated. density. profile: where p, and ry are the characteristic density ancl the scale racius of the distribution. | Recent numericalsimulations by \citet{Mo:98} yielded a more concentrated density profile: where $\rho_s$ and $r_s$ are the characteristic density and the scale radius of the distribution. |
This density cistribution has an even sleeper cusp (ptsXV Laefor r: 0) than the previous one. | This density distribution has an even steeper cusp $\rho_{\rm Moore}
\propto r^{-1.5}$ for $r \rightarrow 0$ ) than the previous one. |
Similarly to the NEW halo. we consider this profile as having only one free parameter. | Similarly to the NFW halo, we consider this profile as having only one free parameter. |
Following Mooreetal. (1999)... we define exij;ose as being L8 times smaller than exp: it is then derived from Iq. 9.. | Following \citet{Mo:99}, we define $c_{\rm Moore}$ as being 1.8 times smaller than $c_{\rm NFW}$; it is then derived from Eq. \ref{cmvir}. |
For a given virial radius. the scale radius 7; of the Moore halo will then be 1.8 times larger than its corresponding quantity for the NEW halo. | For a given virial radius, the scale radius $r_s$ of the Moore halo will then be 1.8 times larger than its corresponding quantity for the NFW halo. |
p, can be derived from: Also in this case we constrain the virial halo mass to be lower than 8-104 M. | $\rho_s$ can be derived from: Also in this case we constrain the virial halo mass to be lower than $8 \times 10^{11}$ $_{\odot}$. |
Early studies of rotation curves (Bosma1981) noted the fact that the ratio between the surface density and the dark matter surface density is approximately constant in the outer parts of galaxies (but. see. Corbelli&Salucei2000)). | Early studies of rotation curves \citep{Bo:81}
noted the fact that the ratio between the surface density and the dark matter surface density is approximately constant in the outer parts of galaxies (but see \citealt{Cor:00}) ). |
This led to the hypothesis that dark matter could in some wav be associated with the disc and cistributed in the same manner: this is what is reasonable to expect in the case of models considering for instance HI» clumps as a component of dark matter (Pfenniger.Combes1994). | This led to the hypothesis that dark matter could in some way be associated with the disc and distributed in the same manner; this is what is reasonable to expect in the case of models considering for instance $_2$ clumps as a component of dark matter \citep*{Pf:94}. |
. In this case the scaling factor for the contribution to the rotation curve is a free parameter. | In this case the scaling factor for the contribution to the rotation curve is a free parameter. |
According to MOND. the law of Mocified Newtonian Dynamics (Ailerom—1983)... there exists ᾱ certain acceleration ay below which Newton's law of gravity is no longer valid ancl the expression for the gravitational acceleration reads: where A(r) account for the stellar and gaseous components and au=12;10cms 7 (Begeman.Brocils&Sanders1991 ). | According to MOND, the law of Modified Newtonian Dynamics \citep{Mi:83}, , there exists a certain acceleration $a_0$ below which Newton's law of gravity is no longer valid and the expression for the gravitational acceleration reads: where $M(r)$ account for the stellar and gaseous components and $a_0=1.2 \times 10^{-8}$ cm $^{-2}$ \citep*{Beg:91}. |
. The fits were performed by a X7-mininisation. considering both the rotational velocities and their logarithmic gradients (Vo=at), which bear a crucial information on the matter distribution in a ealaxy (see Persic&Salucci1990)). | The fits were performed by a $\chi^2$ -minimisation, considering both the rotational velocities and their logarithmic gradients $\nabla= \frac {d{\rm log} V(r)}{d{\rm log}r}$ ), which bear a crucial information on the matter distribution in a galaxy (see \citealt{PS:90}) ). |
The total 47 value to be mininiised then is xz,=No,|Ns | The total $\chi^2$ value to be minimised then is $\chi^2_{tot}=\chi^2_{vel}+\chi^2_{\nabla}$. |
It is worthwhile to point out that the X7 values should only ος considered as à way to compare the dilferent fits within he same galaxy. rather than a probability indicator. because he choice of the error bars is quite subjective and. we plot wo points per beam. so the points are not independent: the goodness of a particular mass model is also related to the raction of observational points that it hits within 1 σ as well as the ones that it baclly misses. | It is worthwhile to point out that the $\chi^2$ values should only be considered as a way to compare the different fits within the same galaxy, rather than a probability indicator, because the choice of the error bars is quite subjective and we plot two points per beam, so the points are not independent; the goodness of a particular mass model is also related to the fraction of observational points that it hits within 1 $\sigma$ as well as the ones that it badly misses. |
μα Figs. | In Figs. |
9 to 13 we show. for each galaxy. the results of he fits. the residuals (Vins— τω) of the fits and the 10 obabilitv contours in parameter space. | \ref{116} to \ref{7339} we show, for each galaxy, the results of the fits, the residuals $V_{obs}-V_{model}$ ) of the fits and the 1 $\sigma$ probability contours in parameter space. |
The case of NGC 7339 will be discussed in Appendix A. The Burkert prolile so as any cored profile has the vest fits to the rotation curves. with no systematic deviation rom the observed rotation curves seen in all galaxies. | The case of NGC 7339 will be discussed in Appendix A. The Burkert profile – so as any cored profile – has the best fits to the rotation curves, with no systematic deviation from the observed rotation curves seen in all galaxies. |
None of our LOO data points (considering the five galaxies ogether) is inconsistent with this model. having a residual weer than 3 σ (where & is the observational error). | None of our $\sim$ 100 data points (considering the five galaxies together) is inconsistent with this model, having a residual larger than 3 $\sigma$ (where $\sigma$ is the observational error). |
The stellar banc mass-to-light ratios. which lie between 0.5 and LS. are consistent with population synthesis mocoels (ce... Dell&deJong 2001)). | The stellar I-band mass-to-light ratios, which lie between 0.5 and 1.8, are consistent with population synthesis models (e.g., \citealt{BdJ:01}) ). |
The core radii are in the range (0.7 μεν and the central densities are between (0.4 3) 10?! e em. | The core radii are in the range (0.7 – 2.3) $\times~ r_{opt}$, and the central densities are between (0.4 – 3) $\times$ $^{-24}$ g $^{-3}$. |
In Fig. | In Fig. |
LE we plot the galaxies of our sample in the poMoore plane of Burkert (1995).. slightly adapted to spiral galaxies by Salucci&Burkert (2000): despite à certain scatter. they roughly follow the relation. which certainly has animplication for the nature of dark matter. | \ref{bur} we plot the galaxies of our sample in the $\rho_0~-~r_{core}$ plane of \citet{B:95}, , slightly adapted to spiral galaxies by \citet{SB:00}: : despite a certain scatter, they roughly follow the relation, which certainly has animplication for the nature of dark matter. |
The minimum X values for the NEW haloes are significantly higher than for theBurkert haloes. | The minimum $\chi^2$ values for the NFW haloes are significantly higher than for theBurkert haloes. |
The former fail to reproduce both the velocities and. the shape of the observed. rotation curves. | The former fail to reproduce both the velocities and the shape of the observed rotation curves. |
Moreover. there is à systematic | Moreover, there is a systematic |
latter correlation test GL, versus Adio). the result was the same with both Adio values for VY Aqr. | latter correlation test $L_{x}$ versus $M_{WD}$ ), the result was the same with both $M_{WD}$ values for VY Aqr. |
Excluding GW Lib decreased the significance to 91 per cent (L, versus AZo nae)and to 63 per cent CL, versus Adi). | Excluding GW Lib decreased the significance to 91 per cent $L_{x}$ versus $kT_{max}$ )and to 63 per cent $L_{x}$ versus $M_{WD}$ ). |
We have analvsed the X-ray spectra of 13 dwarl novae with accurate parallax-based distance estimates. and derived the most accurate shape for the X-ray luminosity function. of DNe in the 210 keV band to date due to accurate distance measurements and due to the fact that we did not use an X-ray selected sample. | We have analysed the X-ray spectra of 13 dwarf novae with accurate parallax-based distance estimates, and derived the most accurate shape for the X-ray luminosity function of DNe in the 2–10 keV band to date due to accurate distance measurements and due to the fact that we did not use an X-ray selected sample. |
The derived X-ray luminosities are. located between ~ LO 107 erg showing. a peak at ~ 10"'E erg +. | The derived X-ray luminosities are located between $\sim$ $^{28}$ $^{32}$ erg $^{-1}$, showing a peak at $\sim$ $^{30}$ erg $^{-1}$. |
Thus. we have obtained peal: luminosities which are lower compared to other previous studies of CV. luminosity functions. | Thus, we have obtained peak luminosities which are lower compared to other previous studies of CV luminosity functions. |
The shape of the X-ray luminosity function of the source sample suggests that the two following scenarios are possible: 1) the sample can be described by a power law with a single a slope. but the sample becomes more incomplete below ~ 3. 107"t cre + than it+ is+ above this. limit.. or. 2) the shape of the real X-ray luminosity function of dwarf novae is à broken power law with a break at around. 3 107 org | The shape of the X-ray luminosity function of the source sample suggests that the two following scenarios are possible: 1) the sample can be described by a power law with a single $\alpha$ slope, but the sample becomes more incomplete below $\sim$ 3 $\times$ $^{30}$ erg $^{-1}$ than it is above this limit, or, 2) the shape of the real X-ray luminosity function of dwarf novae is a broken power law with a break at around 3 $\times$ $^{30}$ erg $^{-1}$. |
The integrated. luminosity between 1 107 erg s and the maximum luminosity of the sample. 1.50. 1077 ere lods LAS 107 erg . | The integrated luminosity between 1 $\times$ $^{28}$ erg $^{-1}$ and the maximum luminosity of the sample, 1.50 $\times$ $^{32}$ erg $^{-1}$, is 1.48 $\times$ $^{32}$ erg $^{-1}$. |
In order to better constrain the integrated luminosity and. the slope of the N-ray luminosity function. more dwarf novae need to be included in the sample. | In order to better constrain the integrated luminosity and the slope of the X-ray luminosity function, more dwarf novae need to be included in the sample. |
Εις. we suggest more future N-rav. imaging observations of dwarf novae in the 210 keV band with accurate distance measurements. | Thus, we suggest more future X-ray imaging observations of dwarf novae in the 2–10 keV band with accurate distance measurements. |
The total X-ray emissivity of the sample within a radius of 200 pe is 1.81. 1075 erg LN (210 keV). | The total X-ray emissivity of the sample within a radius of 200 pc is 1.81 $\times$ $^{26}$ erg $^{-1}$ $^{-1}_{\odot}$ (2–10 keV). |
This accounts for ~ 16 per cent of the total X-ray. emissivity of CVs as estimated by(2006).. and. ~ 5 per cent of the Galactic Riclge X-ray enmissivitsy. | This accounts for $\sim$ 16 per cent of the total X-ray emissivity of CVs as estimated by, and $\sim$ 5 per cent of the Galactic Ridge X-ray emissivity. |
Vhe X-ray luminosities and the inclinations of our sample do not show anti-correlation which has been seen in other previous correlation studies. but a strong correlation is seen between the X-ray luminosities and the orbital periods. | The X-ray luminosities and the inclinations of our sample do not show anti-correlation which has been seen in other previous correlation studies, but a strong correlation is seen between the X-ray luminosities and the orbital periods. |
Also. evidence for a correlation between the white dwarf masses and the shock temperatures exists. | Also, evidence for a correlation between the white dwarf masses and the shock temperatures exists. |
In the future. larger cbwarl nova samples are needed in order to confirm these results. | In the future, larger dwarf nova samples are needed in order to confirm these results. |
This research has made use of cata obtained from. the satellite. a collaborative mission between the space agencies of Japan (JANA) and the USA (NASA). | This research has made use of data obtained from the satellite, a collaborative mission between the space agencies of Japan (JAXA) and the USA (NASA). |
JO acknowledges support from. STEC. | JO acknowledges support from STFC. |
. Part of this work ds based on observations obtained with XAZAZ-Neiwlon. an ESA science mission with instruments and contributions directly Funded by ESA Member States and the USA (NASA). | Part of this work is based on observations obtained with , an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA). |
We thank the reviewer M. Hoevnivtsev for his helpful comments on this paper. | We thank the reviewer M. Revnivtsev for his helpful comments on this paper. |
Iu recent vears. the WWide-Field Camera (WFC) aud theExplorer All-Skyv Monitor AASMD have identified two classes of faint low-1ass N-vav binary (LAINB) that iav be closely related: ow-huuimositv trausicuts and low-huuinosity bursters. | In recent years, the Wide-Field Camera (WFC) and the All-Sky Monitor ASM) have identified two classes of faint low-mass X-ray binary (LMXB) that may be closely related: low-luminosity transients and low-luminosity bursters. |
The low-hnuumositv transients consist of a rather Inhomogeneous eroup of about 15 LAINBs with outbursts hat last from a few davs to several mouths. aud. ect no xighter than a few times 1076st. which corresponds ο an accretion rate of 10 citep|e.g..|[zand(0.. | The low-luminosity transients consist of a rather inhomogeneous group of about 15 LMXBs with outbursts that last from a few days to several months, and get no brighter than a few times $10^{36}$, which corresponds to an accretion rate of $10^{-10}$ \\citep[e.g.,][]{zand00}. |
This distinguishes the low-huninosity ransicuts frou more casily detected transient LAINBs such as Aql X-1 that usually exhibit outbursts brighter han 1075 and from which faint outbursts are ess conumuon (e.g.Simon2002). | This distinguishes the low-luminosity transients from more easily detected transient LMXBs such as Aql X-1 that usually exhibit outbursts brighter than $10^{37}$, and from which faint outbursts are less common \citep[e.g.,][]{sim02}. |
. The faintuess of these outbursts has been attributed to average mass trausfor rates of M<10HALL citepzando0, kingdO.. | The faintness of these outbursts has been attributed to average mass transfer rates of $\dot{M} \lesssim 10^{-11}$ \\citep{zand00, king00}. |
The low-luninosity trausieuts have attracted particular attention because they include the our known accreting millisecond X-ray pulsars (Wijuaudsetal.2002:Maiurkwiurdt 2003). | The low-luminosity transients have attracted particular attention because they include the four known accreting millisecond X-ray pulsars \citep{wk98,mar02,gal02,mar03}. |
. For this reason. it as been livpothesized that the low average accretion rates allow relatively strong C105 Cass) maenetie fields to oersist on the surfaces of the ucutron stars among these LMXDSs. whereas the surface field is buried im svstemis with Neher accretion rates (Cunuuus. Zweibel Bildsten 2001). | For this reason, it has been hypothesized that the low average accretion rates allow relatively strong $> 10^{8}$ Gauss) magnetic fields to persist on the surfaces of the neutron stars among these LMXBs, whereas the surface field is buried in systems with higher accretion rates (Cumming, Zweibel, Bildsten \nocite{cum01}. |
. The low-huuinositv X-rav bursters are sources frou which bright thermomuclear X-ray bursts (sce Lewin. vau Daradijs. Taam 1993 for a lave been observed with the WWEC. and vet there was no evidence for X-ray emission roni persistent accretion at the time of the burst (Cocchietal.2001:Cornelisse 2002a.b). | The low-luminosity X-ray bursters are sources from which bright thermonuclear X-ray bursts (see Lewin, van Paradijs, Taam 1993 for a \nocite{lvt93} have been observed with the WFC, and yet there was no evidence for X-ray emission from persistent accretion at the time of the burst \citep{coc01,cor02a,cor02b}. |
. This sample of low- bursters may represent a large population of uidiscovered neutron star X-ray binaries. depending upon iow often these systems produce bursts. | This sample of low-luminosity bursters may represent a large population of undiscovered neutron star X-ray binaries, depending upon how often these systems produce bursts. |
The intervals )etwoeen X-ray bursts is a strong function of the accretion rate per unit area onto the neutron star. which determines tow quickly a sufficicut column of material is collected for icliumi burning to become unstable (e.g..Bildsten2000). | The intervals between X-ray bursts is a strong function of the accretion rate per unit area onto the neutron star, which determines how quickly a sufficient column of material is collected for helium burning to become unstable \citep[e.g.,][]{bil00}. |
. Since the nuclear energy iu accreted materialis at most of the exavitational cucrey that it emitted during accretion (Lewinetal.1993).. sufficient. nuclear energy to produce an easily-detectable 10°? cre burst is collected iu 5 hows ou the suface of a neutron star producing 10°? oof X-rays through accretion (corresponding tfo an accretion rate of 10tt Ly) while such a burst could uot occur for several vears on a star enüitting persistently at 1072 (104! 1jj | Since the nuclear energy in accreted material is at most of the gravitational energy that it emitted during accretion \citep{lvt93}, sufficient nuclear energy to produce an easily-detectable $10^{39}$ erg burst is collected in 5 hours on the surface of a neutron star producing $10^{35}$ of X-rays through accretion (corresponding to an accretion rate of $10^{-11}$ ), while such a burst could not occur for several years on a star emitting persistently at $10^{32}$ $10^{-14}$ ). |
Unfortuately. thePeppoSAX WEC could only place an upper limit of 1079 oon the huunositv of persisteutly faint bursters close to the Galactic center. so estimates of the frequency of bursts from the low-huunosity bursters are uncertain bv several orders of magnitude. | Unfortunately, the WFC could only place an upper limit of $10^{36}$ on the luminosity of persistently faint bursters close to the Galactic center, so estimates of the frequency of bursts from the low-luminosity bursters are uncertain by several orders of magnitude. |
Therefore. the nuuber of low-huninosity bursters can oulv be guessed within a factor of 1000. and it is still possible that παν of them are also faint N-rav trausicuts. | Therefore, the number of low-luminosity bursters can only be guessed within a factor of 1000, and it is still possible that many of them are also faint X-ray transients. |
ds both a faiut transient svsteni aud οΓΕ, | is both a faint transient system and a low-luminosity |
The spectrum of the secimentecd powder is in good agreement with the result οἱ the calenlation for a CDE with PE environment (see velο indicatinglhepresenceofaconsiderab(sdmnaet oar Π Πω [romsphericals | The spectrum of the sedimented powder is in good agreement with the result of the calculation for a CDE with PE environment (see \\ref{f:PE_CDE}) ), indicating the presence of a considerable fraction of grains with shapes far away from spherical symmetry. |
Conversely. we find that the mean CDE is a good representation of the shape distribution of a real calcite powder with grain sizes within the Ravleigh limit. | Conversely, we find that the mean CDE is a good representation of the shape distribution of a real calcite powder with grain sizes within the Rayleigh limit. |
This justilies the use of (his shape distribution for the comparison with observed. data αἱ large wavelengths. | This justifies the use of this shape distribution for the comparison with observed data at large wavelengths. |
In the following (see sect.44.1) we wil apply this model to simulate dust spectra at low temperatures. | In the following (see 4.1) we will apply this model to simulate dust spectra at low temperatures. |
The laboratory spectra presented so far are especially relevant [or a comparison with FUR spectra of cold dust in objects such as the planetary nebula NGC 6302. | The laboratory spectra presented so far are especially relevant for a comparison with FIR spectra of cold dust in objects such as the planetary nebula NGC 6302. |
Ilence. we shall first give a brief review οἱ some key properties of the dust spectrum of this object. | Hence, we shall first give a brief review of some key properties of the dust spectrum of this object. |
694m and another possible broad band feature in (he range of jum. While the jum feature was attributed to crystalline ice (see also Waters οἱ 11996). the uim feature was suspected to be due to crvstalline forsterite and the question for (he carrier of the jmi band remained open. | Barlow (1997) observed features 65 and $\mu$ m and another possible broad band feature in the range of $\mu$ m. While the $\mu$ m feature was attributed to crystalline ice (see also Waters et 1996), the $\mu$ m feature was suspected to be due to crystalline forsterite and the question for the carrier of the $\mu$ m band remained open. |
Molster οἱ ((2001) identified the e-65 mn band as a blend. of diopside and crystalline water ice with enstatite. | Molster et (2001) identified the $\sim$ $\mu$ m band as a blend of diopside and crystalline water ice with enstatite. |
They also confirmed the reality. of the broad feature around jn and suspected that there should be a verv cold dust component present in the nebula. | They also confirmed the reality of the broad feature around $\mu$ m and suspected that there should be a very cold dust component present in the nebula. |
Kemper et ((2002a.b) first. assigned and the broad emission band around 90;0n to cold calcite dust. | Kemper et (2002a,b) first assigned and the broad emission band around $\mu$ m to cold calcite dust. |
As already mentioned above. the temperature of the carbonate dust was asstumecl to be in the 30-GOIXIX range. | As already mentioned above, the temperature of the carbonate dust was assumed to be in the K range. |
A mass [raction ol less than (for calcite as well as for dolomite) has been derived by these autliors. | A mass fraction of less than (for calcite as well as for dolomite) has been derived by these authors. |
We retrieved an ISO-SWS and an I5O-INS spectrum of NGC 6302 from the ISO archive and reduced it by means of the OLP version 10.0. | We retrieved an ISO-SWS and an ISO-LWS spectrum of NGC 6302 from the ISO archive and reduced it by means of the OLP version 10.0. |
From (he composite spectrum (ranging from jm). we subtracted a combination of Planck functions for temperatures of 30. 55 and IXIx. The remaining residual dust emission. is shown in for the jan range: the | From the composite spectrum (ranging from $\mu$ m), we subtracted a combination of Planck functions for temperatures of 30, 55 and K. The remaining `residual dust emission' is shown in \\ref{f:resid1} for the $\mu$ m range; the |
These optical path mocdulations can be servo controlled using opto-electronie systems previously developed for such kinds of applications (Delageetal (2000): Olivieretal (2005))). | These optical path modulations can be servo controlled using opto-electronic systems previously developed for such kinds of applications \cite{D}; ; \cite{O2}) ). |
Lt allows to monitor. with a manometric accuracy. the lincarity of the optical path variation as a function of time. | It allows to monitor, with a nanometric accuracy, the linearity of the optical path variation as a function of time. |
For cach scan. the olfset (OO,fLyay is set in order to display. the signal. that would be observed at. y positionPI forB a classical spatial hypertelescope. | For each scan, the offset $\frac{(OO_i)_y \cdot y}{f}$ is set in order to display the signal that would be observed at $y$ position for a classical spatial hypertelescope. |
Subsets and Splits